diff --git "a/9708.json" "b/9708.json" new file mode 100644--- /dev/null +++ "b/9708.json" @@ -0,0 +1,1152 @@ +{ + "9708/astro-ph9708101_arXiv.txt": { + "abstract": "We show that the soft X-ray spectra and light curves observed with the $ROSAT$ and $EUVE$ from the closest known millisecond pulsar J0437--4715 can be interpreted as thermal radiation from two hot polar caps whose emitting layers (atmospheres) are comprised of hydrogen. The simplest model yields a uniform temperature of $(0.8-0.9)\\times 10^6$~K within a cap radius of $0.7-0.9$~km. The spectral fits indicate that the temperature may be nonuniformly distributed along the cap surface. The distribution can be approximated by a central core heated up to $(1-2)\\times 10^6$~K within a radius of $0.2-0.4$~km, surrounded by a colder rim with temperatures $(3-5)\\times 10^5$~K extending out to $2-6$~km. The polar cap interpretation implies low column densities, $(1-3)\\times 10^{19}$~cm$^{-2}$, and a high degree of ionization, $> 20$\\%, of the interstellar hydrogen towards the pulsar. The inferred bolometric luminosity of the polar caps, $(1.0-1.6)\\times 10^{30}$~erg~s$^{-1}$, is in excellent agreement with the predictions of the slot-gap model of radio pulsars developed by Arons and his coworkers. Similar polar cap radiation should be emitted by other millisecond pulsars, although in some of them (e.~g., PSR B1821--24) the soft X-ray flux is dominated by the nonthermal radiation from pulsar magnetospheres. ", + "introduction": "The bright 5.75~ms pulsar J0437--4715 was discovered by Johnston et al.~(1993) during the Parkes survey for millisecond pulsars. The pulsar is in a 5.74~d binary orbit with a cool ($T_{\\rm color}\\simeq 4000$~K), low-mass ($\\sim 0.2 M_\\odot$) white dwarf companion and is surrounded by a bow-shock nebula (Becker et al.~1993; Bailyn 1993; Bell, Bailes \\& Bessel 1993; Danziger, Baade \\& Della Valle 1993; Bell et al.~1995). % It is an old object, with a characteristic age $\\tau= P/2\\dot{P}\\simeq 5\\times 10^9$~yr, low magnetic field $B\\sim 3\\times 10^8$~G, and rotational energy loss $\\dot{E}=4\\times 10^{33}$~erg~s$^{-1}$. The observed dispersion measure of 2.65~pc~cm$^{-3}$ implies a distance $d\\simeq 100-180$~pc, making this the closest known millisecond pulsar. Sandhu et al.~(1997) reported the distance $d=178\\pm 26$~pc, from parallax measurements. The pulsar shows significant radio emission over at least 80\\% of the pulse period, with a complicated mean pulse shape varying with radio frequency. Variation of the linear polarization position angle within the mean pulse interpreted in terms of the rotating vector model (Radhakrishnan \\& Cooke 1969) yields the angle between the observer's line of sight and the rotation axis, $\\zeta\\simeq 40^\\circ$, and the angle between the magnetic and rotation axes, $\\alpha\\simeq 35^\\circ$ (Manchester \\& Johnston 1995). $ROSAT$ observations of \\psr~ with the Position Sensitive Proportional Counter (PSPC) have revealed (Becker \\& Tr\\\"umper 1993; hereafter BT93) that this is also a bright (count rate $=0.204\\pm 0.006$~s$^{-1}$) soft X-ray pulsar with a single broad pulse and a pulsed fraction $\\fp =33\\pm 3\\%$ in the PSPC energy range $0.1-2.4$~keV. BT93 found that the pulsed fraction varies with photon energy $E$ and peaks in the range $0.6-1.1$~keV, reaching $53\\pm 6\\%$. Although the pulsar spectrum can be fitted with a single power law, % indicative of a non-thermal origin of the soft X-ray radiation, the energy dependence of the pulsed fraction in the narrow energy range makes this interpretation hardly plausible. BT93 showed also that a single blackbody model does not fit the PSPC spectrum leaving a residual hard excess above 0.4~keV. They suggested that the spectrum can consist of two components: a power law, representing magnetospheric or nebular emission, % and a blackbody component of the % temperature $T\\sim 1.7\\times 10^6$~K emitted from an area of $\\sim 0.08\\, d_{180}^2~{\\rm km}^2$ ($d_{180}=d/180~{\\rm pc}$); this thermal component was suggested by BT93 to be radiated from a hot spot on the neutron star (NS) surface. \\psr~ was also detected with the $EUVE$ Deep Survey Instrument (DSI) in the Lexan band filter ($E\\simeq 0.05-0.2$~keV) by Edelstein, Foster \\& Bowyer (1995) and Halpern, Martin \\& Marshall (1996; hereafter HMM96). According to Edelstein et al.~(1995), the DSI count rate of the source is $0.0143\\pm 0.0008$~s$^{-1}$, whereas HMM96 report the count rate of $0.00973\\pm 0.00017$~s$^{-1}$ obtained with a much longer exposure (496~ks vs.~72~ks). Edelstein et al.~(1995) ruled out the single power-law spectral model based on the improbably high hydrogen column density, $n_H=2.5\\times 10^{20}$~cm$^{-2}$, required by this model. They claimed that both the $ROSAT$ flux below 0.4~keV and the $EUVE$ flux could arise from an isothermal blackbody with a temperature $\\sim 5.7\\times 10^5$~K, an emitting area of $\\sim 3$~km$^2$, and an absorbing column of $n_H=5\\times 10^{19}$~cm$^{-2}$. On the other hand, HMM96 concluded that the combined analysis of the $ROSAT$ spectrum and $EUVE$ flux is consistent with a single power-law spectrum of photon index $\\gamma=2.2-2.5$ and intervening column density $n_H=(5-8)\\times 10^{19}$~cm$^{-2}$. Alternatively, the combined data can be interpreted as comprised of two components, e.~g., a power law and a blackbody component emitted from a hot polar cap of radius 50--600~m and temperature $(1.0-3.3)\\times 10^6$~K. HMM96 observed the pulsar in the high time resolution mode, which enabled them to obtain the light curve and to measure the pulsed fraction $\\fp = 27\\pm 5\\%$ in the $EUVE$ DSI spectral range. Thus, although very important data have been collected and analyzed, the true nature of the soft X-ray radiation from \\psr~ remains elusive. The only firmly established facts are that the radiation is pulsed (pulsed fraction apparently depends on energy), and the spectrum cannot be fitted with a single blackbody model. The main question is whether the radiation is of a nonthermal (magnetospheric? nebular?) origin or at least a fraction of it can be interpreted as thermal (or thermal-like) radiation from some heated regions (polar caps?) on the NS surface. Virtually all the (different) models of radio pulsars (e.~g., Cheng \\& Ruderman 1980; Arons 1981; Michel 1991; Beskin, Gurevich \\& Istomin 1993) predict these objects to have {\\em polar caps} (PCs) around the NS magnetic poles heated up to X-ray temperatures by relativistic particles and gamma-quanta impinging onto the pole regions from the acceleration zones. A conventional assumption about the PC radius is that it is close to the radius within which open magnetic field lines originate from the NS surface. For \\psr, it gives $\\Rpc =1.9\\, (\\rns /10\\,{\\rm km})^{3/2}\\,(P/5.75\\,{\\rm ms})^{-1/2}$ km. Expected PC temperatures, $T_{\\rm pc}\\sim 5\\times 10^5 - 5\\times 10^6$~K, and luminosities, $L_{\\rm pc} \\sim 10^{28}-10^{32}$~erg~s$^{-1}$, are much less certain, being strongly dependent on the specific pulsar model. Thus, one cannot firmly predict PC properties because of the lack of a well-established pulsar model --- rather a theoretical model should be chosen based on X-ray observations of radio pulsars. X-ray observations of other old pulsars (e.~g., PSR B1929+10 -- Yancopoulos, Hamilton \\& Helfand 1994; PSR B0950+08 -- Manning \\& Willmore 1994) do show pulsed X-ray radiation which, in principle, could be the thermal PC radiation. However, the number of photons collected has been too small to make firm conclusions, and the opposite hypothesis, that this radiation is of a magnetospheric origin (\\\"Ogelman 1995; Becker \\& Tr\\\"umper 1997 -- hereafter BT97), cannot be excluded. In principle, one could also observe the PC radiation from younger pulsars and use these data to discriminate between different PC models. Indeed, there are some indications that hard components of the soft X-ray spectra of, e.~g., PSR B0656+14 and PSR B1055--52 (Greiveldinger et al.~1996) may consist of two subcomponents, a power law and a thermal subcomponent corresponding to emission from PCs of temperatures $\\Tpc \\sim 1.5\\times 10^6$~K (for PSR B0656+14) and $\\sim 3.7\\times 10^6$~K (for PSR B1055--52). This interpretation, however, is not unique because it is difficult to separate the hard component from the soft one which is believed to originate from the (cooler) entire NS surface, and even more difficult to separate the two subcomponents of the hard spectral tail. One cannot also exclude {\\em a priori} that the nonthermal radiation from the pulsar magnetosphere or a pulsar-powered compact nebula contributes to, or even dominates, the observed soft X-ray flux of \\psr. It is natural to assume that the spectrum of this radiation can be approximated by a power law in the relatively short $EUVE$-$ROSAT$ range. One could also expect the magnetospheric (but not nebular) radiation to be pulsed with the radio pulsar period; the pulses should be, as a rule, narrower than those of the thermal PC radiation, and the shape of the light curve should not vary considerably with photon energy. Thus, if more data confirm the conclusion of BT93 that the pulsed fraction depends on $E$, the radiation of \\psr~ should either be thermal or consist of the thermal and nonthermal components. In the latter case, since the sources of the thermal and nonthermal radiation are expected to be spatially separated, it is natural to expect that pulsations of these two components should be phase-shifted due to the difference of travel times and aberration. For instance, if the nonthermal component is generated at a distance comparable to the light cylinder radius ($R_{\\rm lc}\\sim Pc/2\\pi \\sim 3\\times 10^7$ cm), the time delay leads to a phase shift of about 0.2 of the pulsar period. No energy-dependent phase shifts have been reported for this source, although it can be explained by the poor photon statistics in hard PSPC channels. As in the case of thermal PC radiation, the current theoretical models of nonthermal high-energy pulsar emission (e.~g., Sturner, Dermer \\& Michel 1995; Romani 1996) are not enough elaborated to predict the intensity, spectrum and light curve for a given pulsar. There exist empirical formulae (Seward \\& Wang 1988; \\\"Ogelman 1995; BT97) which relate the (presumably nonthermal) X-ray luminosity to the pulsar parameters, e.~g., to the period and magnetic field. For instance, \\\"Ogelman (1995) pointed out that the observed soft X-ray luminosities satisfy, for 7 pulsars, the following equation: $L_x\\simeq 6.6\\times 10^{26}\\, (B_{12}/P^2)^{2.7}$~erg~s$^{-1} \\propto \\dot{E}^{1.35}$. % BT97 fitted the luminosities of 26 pulsars observed in the $ROSAT$ range with the dependence $L_x \\simeq 0.001 \\dot{E}$. These pulsars represent a very wide range of spin-down luminosities ($10^{33}-10^{39}$~erg~s$^{-1}$), ages ($10^3-7\\times 10^9$~yr), magnetic field strength ($10^8-10^{13}$~G) and spin periods ($1.6-530$~ms). The inferred dependence indicates that for most pulsars the bulk of observed radiation is directly connected with pulsar mechanisms, i.~e., with production and acceleration of relativistic particles which carry away the NS rotational energy. The fact that X-rays from some of these pulsars (e.~g., Crab) are certainly nonthermal may allow one to assume that the radiation detected with the $ROSAT$ from most pulsars is of a nonthermal origin. On the other hand, a correlation between $L_x$ and $\\dot{E}$ should take place also for radiation emitted by the pulsar PCs because their (thermal) luminosity is also provided by relativistic particles generated in the magnetosphere and accelerated towards the NS surface, so that the PC luminosity should be a fraction of $\\dot{E}$ which goes to heating of the PCs. Moreover, both theoretical (see Section 5.1) and observational estimates of the PC luminosities may happen to be very close to the values predicted by the $L_x(\\dot{E})$ dependence obtained by BT97. For instance, the luminosity of the thermal component in the power-law plus blackbody fit for \\psr, $L_{\\rm bol}=2.2\\times 10^{30}\\, d_{180}^2$~erg~s$^{-1}$ (BT93), is close to the predicted value, $4\\times 10^{30}$~erg~s$^{-1}$. Another example is PSR B1929+10 whose luminosity $L_{\\rm bol}\\simeq 1.2\\times 10^{30}\\, d_{250}^2$~erg~s$^{-1}$ inferred from the blackbody fit of both $ROSAT$ and $ASCA$ data (Yancopoulos et al.~1994; Wang \\& Halpern 1997) perfectly matches the dependence derived by BT97. In general case, we may expect that the pulsar X-ray radiation contains both thermal (PC) and nonthermal components, both growing with $\\dot{E}$, and the relation between the thermal and nonthermal fluxes may be different for different objects, depending, in addition to $\\dot{E}$, on other intrinsic pulsar parameters (e.~g., pulsar period $P$, magnetic inclination $\\alpha$), as well as on the rotational inclination $\\zeta$ (because the radiation beam widths are different for the thermal and nonthermal components). Thus, one cannot rule out that the PC component dominates in some cases, and we explore this possibility for \\psr. An important evidence on the nature of radiation of \\psr~ could come from deep observations of its high-energy tail, at energies above $1-2$~keV. Although this object has been detected by $ASCA$ (two 20~ks observations) below 3~keV, the results are still not very conclusive because of poor choice of observing modes which did not take into account the presence of a neighboring Seyfert galaxy. Nevertheless, preliminary results show that a spectrum of $\\sim 400$ pulsar photons extracted from the CCD away from the Seyfert galaxy is softer than a power law and resembles more a thermal-like spectrum (Kawai, Tamura \\& Saito 1996). This conclusion, however, needs verification based on longer pointed observations in a mode minimizing contamination from the Seyfert. The conclusion of BT93 and HMM96 that the PSPC spectrum of \\psr~ cannot be fitted with a single thermal component is based on the assumption that its spectrum coincides with the spectrum of the blackbody radiation. However, spectra emitted by real bodies, including stars, are always different from the Planck spectrum. In particular, if the temperature of a stellar atmosphere grows inward, and the absorption coefficient decreases with frequency (e.~g., $k_\\nu \\propto \\nu^{-3}$ for the Kramers law), then the spectrum is harder than the blackbody spectrum at high frequencies because we see deeper and hotter layers. This means that if a NS is covered with a fully ionized plasma, whose opacity can be described by the Kramers law, its spectrum is substantially harder at high energies, $h\\nu \\gapr kT_{\\rm eff}$ (e.~g., Pavlov \\& Shibanov 1978). This general property has been demonstrated in models of hydrogen and helium atmospheres of cooling NSs by Romani (1987), Pavlov et al.~(1995), Rajagopal \\& Romani~(1996), Zavlin, Pavlov \\& Shibanov (1996; hereafter ZPS96). One can expect that a NS has a purely hydrogen atmosphere if it experienced accretion of the interstellar matter --- heavy elements of the accreted matter sink down rapidly due to the strong NS gravitation (Alcock \\& Illarionov 1980). Since \\psr~ is a very old object, it is quite plausible that it accreted some matter during its long life, and in this case we may expect that the excess of the observed flux at $E\\gapr 0.4-0.6$~keV is due to the fact that \\psr, including its PCs, is covered by a fully ionized hydrogen atmosphere. The PC radiation should be inevitably pulsed unless the rotation axis coincides with either the line of sight or the magnetic axis. If it were the blackbody radiation, the shape of the light curves and the pulsed fraction $\\fp$ would remain the same at all photon energies $E$. However, the atmosphere radiation has a very important feature --- it is anisotropic, with anisotropy (and consequently light curves) depending on energy (Pavlov et al.~1994; Shibanov et al.~1995; ZPS96). This dependence is different for different chemical compositions and surface temperatures. E.~g., for atmospheres with relatively high temperatures $\\sim 10^6$~K consisting of light elements (hydrogen, helium), the anisotropy increases with increasing $E$ in the soft X-ray range. This also makes the spectra dependent on the rotation phase. The inter-dependence of the spectral and angular distributions means that the proper interpretation of the observed PC radiation implies fitting of {\\em both the spectra and the light curves with the same PC model}. The above-described properties of radiation emitted by a PC covered with an atmosphere warrant a new investigation of the soft X-ray radiation observed from \\psr, based on the NS atmosphere models. To perform this investigation, we used the $ROSAT$ and $EUVE$ data analyzed previously by BT93 and HMM96 and combined with the results of new observations carried out with the $ROSAT$ High-Resolution Imager (HRI) and $ROSAT$ PSPC (Becker et al.~1997). These observations are briefly described in Section~2. We show that with allowance for the properties of NS atmosphere radiation these joint $ROSAT$ and $EUVE$ data can be interpreted as thermal radiation of two PCs, {\\em without invoking an additional nonthermal component}. In Section~3 we explore a simplest single-temperature PC model which assumes that both PCs have equal radii $\\Rpc$ and uniform temperatures $\\Tpc$, whereas the rest part of the NS surface has a much lower temperature and is invisible in X-rays. Since the magnetic field of \\psr~ is very low, it cannot affect the radiative properties of the PC atmospheres. Therefore, we use the low-field atmosphere models developed by ZPS96. To calculate the flux as measured by a distant observer, we integrate the specific flux over the visible PC surface with allowance for the gravitational redshift and bending of the photon trajectories. For the case of small PCs ($\\Rpc\\ll\\rns$) with uniform temperatures, a convenient expression for the observable flux is given by Zavlin, Pavlov \\& Shibanov (1995; see their Eq.~[A15])). We show that this simple model is generally consistent with the observational data if the PCs are covered with hydrogen or helium, whereas the iron atmosphere model does not fit the spectra (cf. Pavlov et al.~1996b). For the hydrogen and helium models, we estimate $\\Rpc$, $\\Tpc$ and the interstellar hydrogen column density, $n_H$. The fit with the simplest PC model is still not perfect. This is not surprising because it is hard to expect the real PCs to be uniformly heated. On the contrary, due to higher heat conduction of subphotospheric layers, where the energy of accreting relativistic particles is released, the heat can propagate along the surface, eventually heating surface layers out of the ``primary'' hot spot. This should result in a larger hot region with the temperature decreasing outward. In fact, this mechanism should be more efficient just for low-field pulsars because the strong magnetic fields of ordinary pulsars greatly reduce the transverse conductivity (e.~g., Hernquist 1985). To the best of our knowledge, there have been no reliable calculations of the temperature distribution around the pulsar magnetic poles. Thus, to include this effect into consideration, we explore in Section~4 a simple model: the distribution is assumed to be a two-step function (``core+rim'') with two temperatures, $\\Tc$ and $\\Tr$, and two radii, $\\Rc$ and $\\Rr$. We show that this model is fairly consistent, for hydrogen-covered PCs, with all the data available. We discuss the results and implications of our interpretation in Section~5 and draw conclusions in Section~6. ", + "conclusions": "We have shown that both the spectra and the light curves of the soft X-ray radiation of \\psr~ observed with the $ROSAT$ PSPC and HRI and $EUVE$ DSI can be interpreted as {\\em thermal radiation from two hydrogen-covered polar caps}. The simplest, single-temperature PC model allows us to estimate typical PC radius, $\\Rpc\\sim 1$~km, and temperature, $\\Tpc\\sim 1\\times 10^6$~K. The successful fitting of the data with the two-temperature model indicates that the PC temperature may be non-uniform, decreasing from $\\Tc\\sim (1.1-1.8)\\times 10^6$~K at a PC core with a typical size $\\Rc\\sim 150-400$~m down to $\\Tr\\sim (3-5)\\times 10^5$~K at a much greater radius $\\Rr\\sim 2-6$~km, so that the heated area may comprise a considerable fraction of the NS surface. The bolometric luminosity of the two PCs is $L_{\\rm bol} = (1.0-1.6)\\times 10^{30} d_{180}^2$~erg~s$^{-1}$, that is $\\sim (2-4)\\times 10^{-4}$ of the total energy loss of the pulsar is absorbed and re-emitted by the PCs. We emphasize that the proposed interpretation is based essentially upon the {\\em neutron star atmosphere models} --- the data cannot be interpreted as purely thermal PC radiation with the simplistic blackbody model. Moreover, with the aid of the NS atmosphere models we show that the emitting layers should be depleted of heavy elements, which can be naturally explained by the plausible assumption that the old NS has experienced accretion of the hydrogen-rich matter during its long life-time. The successful fits of the $EUVE$ and $ROSAT$ light curves in different energy ranges with the PC models are possible only if the energy-dependent limb-darkening of the hydrogen atmospheres is taken into account. The shape of the light curves is also very sensitive to the bending of photon trajectories in the strong gravitational field of the NS and to the orientations of the rotational and magnetic axes of the pulsar. We obtained the satisfactory fits assuming $\\mns=1.4 M_\\odot$ and $\\rns=10$~km, with the magnetic inclination angle, $\\alpha=35^\\circ$, and viewing angle, $\\zeta=40^\\circ$, inferred from radio observations. We have demonstrated that the effect of the gravitational bending on the light curves can provide important information on the NS mass-to-radius ratio. The hydrogen column density towards the \\psr, $n_H\\sim (1-3)\\times 10^{19}$~cm$^{-2}$, as well as strong ionization of hydrogen, $\\xi > 20\\%$, estimated from the PC model fits are consistent with the ISM properties obtained from observations of other objects in the vicinity of the pulsar, whereas the power-law fit yields greater $n_H$ and lower $\\xi$. This can be considered as one more argument in favor of thermal origin of the X-ray radiation of \\psr. The inferred PC temperature, radius and, especially, luminosity are in excellent agreement with the predictions of the slot-gap pulsar model (Arons 1981). The upper limit on the PC luminosity provided by the outer-gap pulsar model (Cheng et al.~1986) exceeds the observed value. This model becomes compatible with our results if a reduced efficiency of the particle accelerator, consistent with the upper limit on the gamma-ray flux from the pulsar, is assumed. Radio pulsar models predict that heated PCs is a common phenomenon inherent to all active pulsars. By virtue of this, the dependence $L_x\\simeq 0.001 \\dot{E}$ found by BT97 for 26 pulsars detected with the $ROSAT$ % determines a fraction of the rotational energy loss re-emitted in X-rays {\\em by both nonthermal and thermal processes}. Contributions of these two mechanisms to the total X-ray flux depend on various factors, such as pulsar period, radius, magnetic field, etc. In some pulsars (e.~g., PSR B1821--24) the thermal-like soft X-ray radiation from PCs may happen to be less intensive than the nonthermal radiation generated in the pulsar magnetosphere. For instance, both the slot-gap and outer-gap models predict the luminosity of PCs of PSR B1821--24 much lower than the luminosity observed with the $ROSAT$ and $ASCA$. Although all the available observations of \\psr~ are consistent with the suggested PC interpretation, and this interpretation is supported by the theoretical pulsar models and by indirect arguments, we still cannot completely exclude that the same data, collected in the relatively narrow energy range, could be interpreted with another model --- for instance, as a combination of the PC radiation and magnetospheric radiation. Hence, the proposed interpretation can be considered as complementary to the nonthermal interpretation discussed by BT97. A crucial test would be provided by observations of this object at energies $\\gapr 1-2$~keV, which would enable one to firmly discriminate between the power-law and thermal-like spectra, or to separate their contributions. Such observations could be carried out by $ASCA$ (with a sufficiently long exposure) and by the forthcoming $AXAF$, $XMM$ and $ASTRO$-$E$ missions. Important data could be also obtained from observations of the pulsar in the UV range ($1200\\lapr \\lambda\\lapr 3000$~\\AA) with the $HST$. In particular, such observations would enable one to measure the temperature of the entire NS surface (cf.~Pavlov, Stringfellow \\& C\\'ordova 1996a) and to elucidate other possible heating mechanisms competing with the PC heating. For more accurate and reliable interpretation of the future observations, more theoretical work is highly desirable. In particular, the temperature distribution over the NS surface within and around the PCs should be investigated (and further used to fit the data), and the nonthermal X-ray spectrum and beam shape should be calculated." + }, + "9708/astro-ph9708271_arXiv.txt": { + "abstract": "\\noindent The angular correlation function $\\wth$ of faint galaxies is affected both by nonlinear gravitational evolution and by magnification bias due to gravitational lensing. We compute the resulting $\\wth$ for different cosmological models and show how its shape and redshift evolution depend on $\\Omega$ and $\\Lambda$. For galaxies at redshift greater than 1 ($R$ magnitude fainter than about 24), magnification bias can significantly enhance or suppress $\\wth$, depending on the slope of the number-magnitude relation. We show for example how it changes the ratio of $\\wth$ for two galaxy samples with different number-count slopes. ", + "introduction": "The angular correlation function of galaxies has been used to characterize the large-scale distribution of galaxies for over 2 decades. If the number density on the sky at angular position $\\hat{\\psi}$ is $n(\\hat{\\psi})$, then $\\wth$ is defined as \\begin{equation} \\wth=\\frac{\\langle n(\\hat{\\psi})n(\\hat{\\phi})\\rangle}{\\bar{n}^2}-1 \\; . \\label{wthav} \\end{equation} The 3-dimensional unit vectors $\\hat{\\psi}$ and $\\hat{\\phi}$ are used to define the angular separation $\\theta$ as $ \\hat{\\psi} \\cdot \\hat{\\phi} = \\cos(\\theta)$, and $\\bar{n}$ is the mean galaxy number density on the sky. The observed galaxy distribution is well described by a power law $\\wth\\propto \\theta^{-\\gamma}$, with slope $\\gamma\\simeq 0.8$. Since $\\wth$ is a projection on the sky of the 3-dimensional auto-correlation function $\\xi(r, z)$, the above power law for $\\wth$ has been associated with a power law for $\\xi$ with slope $1.8$ which is close to the observed value in the nearby galaxy distribution. Measurements of $\\wth$ are difficult to interpret as it involves a projection of the galaxy distribution out to redshifts of order 1 (e.g. for a sample with limiting $R$ magnitude of about 24.5). Thus the effects of evolution of large scale structure due to gravitational clustering as well as galaxy evolution need to be understood to interpret $\\wth$. Recently Villumsen (1996) has considered the effect on $\\wth$ of gravitational lensing by large-scale structure along the line of sight (Gunn 1967). Lensing increases the area of a given patch on the sky, thus diluting the number density. On the other hand, galaxies too faint to be included in a sample of given limiting magnitude are brightened due to lensing and may therefore be included in the sample. The net effect, known as magnification bias, can go either way: it can lead to an enhancement or suppression of the observed number density of galaxies, depending on the slope of the number-magnitude relation. Variations in the number density which are correlated over some angular separation alter $\\wth$. Villumsen (1996) showed how the linear evolution of density fluctuations along the line of sight can be used to compute the change in $\\wth$ due to magnification bias. The deflections of neighbouring photon trajectories due to lensing by large scale structure are very small, hence the calculation can be done in the limit of weak lensing. In this paper we compute $\\wth$ for different cosmological models taking into account the effects of nonlinear gravitational evolution and gravitational lensing. Since $\\wth$ is a 2-point statistic, even in the fully nonlinear regime it is determined completely by the 3-dimensional power spectrum. We use extensions of the proposal of Hamilton et al. (1991) to include the nonlinear evolution of the power spectrum into the calculation of $\\wth$. We also include the dependence on the cosmological parameters $\\Omega_m$ and $\\Omega_{\\Lambda}$. In Section 2 the formalism for computing $\\wth$ is presented. Section 2.1 discusses the effects of the cosmological model via the growth of density perturbations and the distance-redshift relation. Results for CDM-like power spectra for different values of $\\Omega_m$ and $\\Omega_{\\Lambda}$ are presented in Section 3. We explore ways to isolate the effect of the lensing contribution to $\\wth$ in Section 4 and conclude in Section 5. ", + "conclusions": "We have quantified the effect of gravitational lensing by large-scale structure on the angular correlation function of galaxies $\\wth$, for different cosmological models, on angular scales ranging from $1^\\prime$ to $20^\\prime$. We have taken into account nonlinear gravitational clustering which affects both the intrinsic clustering and lensing contributions to $\\wth$. We find that the ratio of the angular correlation function for red and blue galaxy samples, normalized by the inverse of the relative bias of the two samples, deviates from the value of $1$ expected in the absence of lensing at sufficiently large mean redshifts of the sample. For a mean redshift of 1.3 of the sample, this ratio rises to about $1.5$ for the model with $\\om=1$ and $\\sigma_8=1$, and continues to rise, flattening off, until it reaches $1.7$ by a mean redshift of 2. At $\\langle z \\rangle=1.3$ it rises to about $1.3$ for the model with $\\om=1$ and $\\sigma_8=0.6$, and to about $1.2$ for the open model. In the $\\Lambda-$dominated model, the ratio reaches a value of about $1.4$ at $\\langle z \\rangle=1.3$ , and continues to rise to a value of about $1.6$ at a mean redshift of 2. These values are at $1^\\prime$ angular separation; for larger angles they differ somewhat, but not to a great extent as shown in Fig.~\\ref{wratio}. The largest uncertainty in the application of this result is due to the assumption of two populations with different number count slopes but the same redshift distribution with $\\langle z \\rangle \\largapprox 1$. The effect of magnification bias on the angular correlation function could be used in future surveys, such as the ESO imaging survey (Renzini et al. 1996), to detect weak lensing by large-scale structure and constrain the cosmological parameters $\\om$ and $\\ol$. The ESO imaging survey aims to detect about 200-300 galaxies with redshifts between $1$ and $2.8$ in a field of 25 arcminutes squared, and a similar number with redshifts larger than $2.8$ in a field ten times larger. One problem with trying to detect the effect of magnification bias on $\\wth$ is that unknown physical evolution of the galaxies can modify the intrinsic clustering and introduce uncertainties which may be larger than the lensing signal. The presence of bias evolution would introduce further uncertainties. In future work, we plan to study the cross-correlation of two galaxy samples with different mean redshifts and non-overlapping redshift distributions (Moessner \\& Jain 1997). In this way we hope to minimize the importance of uncertainties due to intrinsic clustering and the exact form of the redshift distribution, and propose a measurable statistic that is dominated by the lensing contribution." + }, + "9708/astro-ph9708029_arXiv.txt": { + "abstract": "Using recent observations of the kinematics of the disk stars the dynamical state of the disk of NGC\\,488 is discussed. The disk is shown to be dynamically `cool', so that NGC\\,488 cannot have experienced many -- even minor -- mergers in the past. ", + "introduction": "It has been argued (Jore, Haynes \\& Broeils 1997) that some of the massive bulges of Sa type galaxies have not been formed during the collapse of the protogalaxy or by secular evolution of the galactic disk (Pfenniger, Combes \\& Martinet 1994), but by capture of satellite galaxies. Jore, Haynes \\& Broeils (1997) present a sample of Sa galaxies with kinematically distinct components in their inner parts such as counter-rotating disks. These might be well interpreted as debris of satellite galaxies which disintegrated while they merged with their parent galaxies. On the other hand, Ostriker (1990) has pointed to the fact that the disks of Sa galaxies are dynamically cool enough to develop spiral structure. Since galactic disks are dynamically heated during the merging process, this sets severe constraints on the accretion rate of satellites (T\\'{o}th \\& Ostriker 1992). Recently Gerssen, Kuijken \\& Merrifield (1997; hereafter referred to as GKM) have observed the kinematics of the stellar disk of NGC\\,488. NGC\\,488 is a typical Sa galaxy which is actually surrounded by dwarf satellites (Zaritsky et al. 1993), so that there might have been indeed minor merger events in the past. GKM discuss implications of their observations, in particular the shape of the velocity ellipsoid, for dynamical disk heating by molecular clouds and transient spiral density waves. But their data allow also to state Ostriker's (1990) objections to the scenario of steady accretion of small satellites in a quantitative way. For this purpose I construct in sections 2 and 3 dynamical models of the bulge and disk of NGC\\,488 using the photometric and kinematical data presented by GKM. In the final section I discuss the dynamical state of the disk and implications for the merging history of NGC\\,488. \\begin{figure}[htbp] \\begin{center} \\leavevmode \\epsffile{bulge.eps} \\caption{Surface brightness profile of the bulge of NGC\\,488 according to the bulge -- disk decomposition by GKM. The left panel shows the surface brightness profile as a de Vaucouleurs profile, while the right panel shows a surface brightness profile according to equation (1) fitted to the data ($\\Box$).} \\label{bulge} \\end{center} \\end{figure} ", + "conclusions": "In order to discuss the viability of the dynamical disk models derived in the previous section I consider two diagnostics. First, I estimate the expected {\\it vertical} scale height of the disk, even though finite scale heights have been neglected in the decompositions of the rotation curve. An isothermal self-gravitating stellar sheet follows a vertical $sech^2(z/z_0)$ density profile (cf. Binney \\& Tremaine 1987), with the vertical scale height given by \\begin{equation} z_0 = \\frac{\\sigma^2_{\\rm W}}{\\pi G \\Sigma_{\\rm d}}\\, , \\end{equation} where $\\sigma_{\\rm W}$ is the dispersion of the vertical velocity components of the stars and $\\Sigma_{\\rm d}$ denotes again the surface density of the disk. Since the bulge is so massive its gravitational force field has to be taken into account. In regions $R^2>>r_{\\rm c,b}^2$ and $z^2< 10^6$. In comparison, our sun has $\\tau \\sim 10^{11}$ and a radius $\\sim 100$~times smaller than the characteristic radius of a thick disk. The thick disk's structure is very much like the envelopes of massive stars, except that self gravity plays no role; the gravity in the disks is dominated by the black hole. The disk will remain geometrically thick if energy is dissipated fast enough. In general, this energy dissipation rate corresponds to an accretion rate which is super-Eddington. The physical parameters of an tidal disruption encounter which produce super-Eddington accretion are described in Ulmer (1997; hereafter U97). Delineating the region in which thick accretion occurs is important not only because super-Eddington tidal disruption events are the brightest and therefore, the most readily observed flares, but also because one expects the accretion of the stellar debris to be fastest when the disk is very thick. In thin disk theory, the accretion time scales as $\\alpha^{-1} (z/r)^{-2}$, where $\\alpha$ is the viscosity parameter, $z$ is the height of the disk, and $r$ is the radial distance. The thicker the disk, the faster it accretes. The outline of the paper is as follows. In \\S~\\ref{secover}, we provide an overview of the novel technique for forming evolutionary sequences out of hydrostatic models of thick accretion disks. The method relies primarily on global conservation laws and the assumption that the viscosity is everywhere positive. In \\S~\\ref{hydrostatic}, we describe the hydrostatic models which compose the evolutionary sequence. The method for evolving the models in time is presented in \\S~\\ref{evolsec}. Initial conditions are given in \\S~\\ref{initcond} for the evolutionary sequences. In \\S~\\ref{results}, we present results, including discussions of the energy spectrum, the residual mass, and the importance of the inner nozzle. Long-term emission properties of thin disks which form after the thick disk phase are described in \\S~\\ref{longterm}. Conclusions are given in \\S~\\ref{consec3}. ", + "conclusions": "} Only a fraction of the disk mass can be consumed in the initial thick disk phase, which is likely to occur for tidal disruptions around $10^6-10^7~M_\\odot$ black holes. This fraction ranges from $\\sim 0-0.99$ for disruptions around $10^6$ to $10^7 M_\\odot$ black holes, and has a typical value of $0.5-0.9$. The amount of the residual mass decreases for larger mass black holes and increases with the pericenter of the disrupted star. This residual material will accrete in a thin disk over a longer period of time. The existence of the initial thick disk phase may reduce the dimming timescale of the disk by a factor of $\\sim 2$ from estimates based on thin disks alone. Assuming that an $0.5 M_\\odot$ initial thick disk, even if the thin disks become advection dominated, the black hole mass to light ratio can rise above $M_\\odot/L_\\odot = 1 $ in no less than 20 (0.1/$\\alpha$) to 2000 (0.1/$\\alpha$) years following a tidal disruption event, depending on the black hole mass and initial orbital of the disrupted star. The long-term emission will be most prevalent around $10^6 M_\\odot$ black holes. If the tidal disruption rates in these galactic nuclei are $\\sim 10^{-4} {\\rm~yr}^{-1}$, then at least $\\sim 10\\%$ of the nuclei should exhibit the long-term UV/optical emission at a level of $\\sim 10^{\\rm 38} {\\rm~ergs~s^{-1}}$." + }, + "9708/astro-ph9708115_arXiv.txt": { + "abstract": "We report the initial results from a program to study the morphology, physical state, and kinematics of the `Diffuse Ionized Medium' (`DIM') in a sample of the nearest and brightest late-type galaxies. For each of five galaxies (NGC 2403, M 81, NGC 4395, M 51, and M 101) we have analyzed deep narrow-band H$\\alpha$ images of the entire star-forming disk, and long-slit spectra of the inner ($\\sim$ 10 kpc) disk with a resolution of 40 to 75 km s$^{-1}$. We find that the DIM covers most of the star-forming disk, and is morphologically related to the presence of high-surface-brightness gas (the giant HII regions). The DIM and the giant HII regions differ systematically in their physical and dynamical state. The DIM is characterized by enhanced emission in the low-ionization forbidden lines ([OI], [NII], and [SII]), and even the high-ionization [OIII]$\\lambda$5007 line is moderately strong in the DIM in at least three cases. This last result contrasts with upper limits on the [OIII] surface brightness in the local DIM of our own Galaxy (the `Reynolds Layer'). We directly verify the inference made by Lehnert \\& Heckman that the DIM contributes significantly to the spatially-integrated (global) emission-line ratios measured in late-type galaxies. We also find that the DIM is more disturbed kinematically than the gas in the giant HII regions. The deconvolved (intrinsic) widths of the H$\\alpha$ and [NII]$\\lambda$6584 lines range from 30 to 100 km s$^{- 1}$ (FWHM) in the DIM compared to 20 to 50 km s$^{-1}$ in the giant HII regions. The high-ionization gas in the DIM is more kinematically disturbed than the low-ionization gas: the [OIII]$\\lambda$5007 lines have intrinsic widths of 70 to 150 km s$^{-1}$. The differing kinematics implies that `the DIM' is not a single monolithic phase of the ISM. Instead, it may consist of a `quiescent DIM' with a low ionization-state and small scale-height (few hundred pc) and a `disturbed DIM' with a high ionization state and moderate scale-height (0.5 to 1 kpc). We argue that the quiescent DIM is most likely photoionized by radiation from O stars leaking out of giant HII regions (although this requires fine-tuning the opacity of galactic disks to ionizing radiation). The disturbed DIM is most likely heated by the mechanical energy supplied by supernovae and stellar winds. Since the disturbed DIM accounts for only a minority ($<$ 20\\%) of the H$\\alpha$ emission in the regions we have studied, there is no fundamental energetics problem with this model, but it does requires mechanically-heated gas to have a high areal covering factor in the inner disk (which needs to be confirmed observationally). We find no clear discontinuity between the physical and dynamical properties of the giant HII regions and the quiescent DIM. The quiescent DIM is morphologically related to the giant HII regions and there is a smooth dependence of the emission-line ratios and emission-line widths on the surface brightness of the emission. Thus, we suggest that a unified approach to the study of the DIM and giant HII regions in star-forming galaxies will prove fruitful. ", + "introduction": "The interstellar medium (ISM) in our Galaxy is known to be a multi-phase complex system. The most recently-discovered major component of the ISM is the wide-spread diffuse ionized medium (\\cite{rey90,kul88}). This gas, usually called the `Reynolds Layer', is characterized by a relatively low ionization state compared to normal HII regions, a low surface brightness, relatively large scale height, and substantial energetic requirements. These latter are so severe that the Reynolds Layer must either soak up nearly 100\\% of the mechanical energy supplied by supernovae and stellar winds or the topology of the interstellar medium must allow a substantial fraction of the ionizing radiation produced by massive stars to escape the Galactic disk and propagate to moderate scale heights. In either case, the implications for our understanding of the structure and energetics of the interstellar medium are considerable. In several nearby normal spiral galaxies, gas with apparently similar properties to the Reynolds Layer has been found as well (see recent reviews by \\cite{det92,ran96conf,walbr96}). This gas has been variously referred to as the `Diffuse Ionized Medium' (DIM), `Diffuse Ionized Gas' (DIG), and `Warm Ionized Medium' (WIM). In this paper we will adopt the acronym `DIM'. In the best studied case of the edge-on galaxy NGC 891 (\\cite{det90,rankulhes90}), a DIM in form of filaments and bubbles can be detected out to more than a few kiloparsec from the disk mid-plane (see also Rand 1997). A DIM with similar chaotic structure can be seen in some other face-on late-type galaxies like M 33 and M 31 (e.g. Courtes et al. 1987; Walterbos \\& Braun 1994). This diffuse gas is (like the Reynolds Layer) characterized by the relative strength of low-ionization emission-lines like [OII]$\\lambda$3727, [SII]$\\lambda\\lambda$6717,6731 and [NII]$\\lambda$6584 (e.g. Dettmar \\& Schulz 1992; Hunter \\& Gallagher 1997; Ferguson et al. 1996b). More generally, Lehnert \\& Heckman (1994) have shown that the integrated optical spectra of a large sample of normal late type galaxies published by Kennicutt (1992a,b) provide very suggestive evidence that the DIM is generic to such galaxies: the global value of the [SII]/H$\\alpha$ emission-line ratio is enhanced by an average factor of about 1.5 to 2.0 relative to the value characteristic of individual extragalactic HII regions. Based on the emission-line ratios in the DIM in the Milky Way and NGC 891, they estimated that the DIM would generically contribute at least 25\\% of the total H$\\alpha$ emission in such galaxies. Indeed, direct H$\\alpha$ imaging of several nearby star-forming galaxies has confirmed this estimate. The fraction of the DIM contribution is in the range of $\\sim$20\\%--40\\%, fairly independent of Hubble type (ranging from early-type spiral(Sb) to irregular) and galaxy inclination (Kennicutt et al. 1995; Ferguson et al. 1996a,b; Hoopes et al. 1996; Hunter \\& Gallagher 1997). While the DIM is apparently ubiquitous, its dynamical state and its ionization and heating mechanism(s) are still poorly understood. To improve our understanding of the structure and evolution of the interstellar medium in our own and other galaxies, we need answers to questions such as how the DIM in external galaxies is energized and how much the DIM in the Milky Way shares in common with the extragalactic DIMs. Thus, we need to verify that DIMs are indeed generic in normal star-forming galaxies, and to elucidate the properties of the DIM in such galaxies. We have therefore conducted an investigation of a carefully selected sample of the nearest, largest, and brightest normal late type galaxies. We have selected our sample of galaxies according to the following criteria: 1) Late Hubble type (Sb or later). Such galaxies have relatively strong and easily studied optical line emission (typical global H$\\alpha$ equivalent widths of 30 to 100$\\AA$---cf. Kennicutt 1992a,b). The integrated emission-line spectra of such galaxies contain a substantial contribution from low-ionization gas (Lehnert \\& Heckman 1994). 2) Nearby (distance $<$ 10 Mpc). This allows us to obtain images with the highest possible spatial resolution (1\\arcsec\\ is typically 20-50 pc), which is crucial for detecting faint filamentary emission and determining its morphology and structure. 3) Large angular size (D$_{25} \\geq$10\\arcmin), so that we have the maximum number of resolution elements across the galaxy. 4) Inclination $\\leq$ 65$\\arcdeg$. We exclude galaxies seen nearly edge-on where line-of-sight projection effects become severe. In this paper, we report our first results for five galaxies in the sample: M 51, M 81, M 101, NGC 2403 and NGC 4395. Some properties of these galaxies are summarized in Table 1. Our data consist of both narrow-band H$\\alpha$ images and long-slit spectroscopic observations. The spectra provide information on emission lines from H$\\alpha$, [SII]$\\lambda\\lambda$6717,6731, and [NII]$\\lambda$6584 in all five galaxies and on the H$\\beta$, [OIII]$\\lambda$5007, and [OI]$\\lambda$6300 lines in subsets of the sample. These data allow us to address the morphology and global energetics of the DIM, as well as its physical and dynamical state. Specifically, we focus on the measurements of the DIM contribution to the global H$\\alpha$ luminosity, the physical state of the DIM (as probed via the relative strengths of the above emission-lines), the dynamical state of the DIM, and the inter-relationship between these physical and dynamical conditions. ", + "conclusions": "\\subsection{The Energy Source for the DIM} \\subsubsection{Photoionization} The results in \\S 3.2.3 suggest that there may be more than one energy source for the DIM. For the quiescent DIM, the `leaky HII region' model is appealing on a number of grounds. First, radiation from massive stars is the most abundant ionization source in these star-forming galaxies (by about an order-of-magnitude). Second, there is a morphological correspondence between the DIM and giant HII regions (Hoopes et al. 1996; Ferguson et al. 1996a and \\S 3.1 above). Third, we have emphasized the {\\it continuity} in the physical and dynamical properties of the gas in the quiescent DIM and the HII regions (rather than the existence of any well-defined dichotomy). Fourth, the quiescent state of this component of the DIM argues against mechanical heating as the dominant ionization source. Finally, even simple models of photoionization by dilute stellar radiation can readily match the relative intensities of the low-ionization species in the quiescent DIM (as we now show). The ionization state of photoionized gas is primarily determined by the ionization parameter (U), defined as the ratio of the density of ionizing photons and electrons in the gas. The low ionization state of the quiescent DIM (and the implied low value for U) is a natural consequence of the `leaky HII region model'. Averaged over the region interior to R$_{25}$, the mean H$\\alpha$ surface-brightness of the disks of our galaxies (including both the DIM and the HII regions) is about $3 \\times 10^{-17}$ erg cm$^{-2}$ s$^{-1}$ arcsec$^{-2}$. This corresponds to a disk-averaged flux of ionizing photons of about 10$^{6}$ s$^{-1}$ cm$^{-2}$ ster$^{-1}$. If we assume that a fraction $f_{DIM}$ of these reach the DIM, then we find U $= 2\\pi \\times 10^{6} f_{DIM}/n_{e}c$ as the characteristic value of U. We can evaluate $n_{e}$ by taking a temperature appropriate for photoionized gas (T $\\sim$ 10$^{4}$ K), and assuming the DIM is in pressure-balance with the rest of the ISM. In the solar neighborhood, P/k $\\sim 10^{4}$ cm$^{-3}$ K$^{-1}$. The surface-mass-density of stars in the inner disks of our galaxies is about an order-of-magnitude higher than the value in the solar neighborhood, while the surface-mass-density of the interstellar gas is roughly similar to the local value. Provided that the scale-height of this gas is not much less than the scale-height of the stars, simple considerations of hydrostatic equilibrium then imply that the mid-plane pressures in the ISM in the inner disks of our galaxies will be of-order P/k $\\sim 10^{5}$ cm$^{-3}$ K$^{-1}$. Since these higher pressures are not directly demonstrated observationally, we will use P/k $\\sim 10^{4}$ to $10^{5}$ cm$^{-3}$ K$^{-1}$. We then obtain U $= 4 \\times 10^{-4} f_{DIM}$ to $4 \\times 10^{-5} f_{DIM}$ as the characteristic value in the mid-plane. These mid-plane values are appropriate for the quiescent DIM since (as we will show in \\S 4.2 below) it must have a relatively small scale-height. Note that the actual value of U in the DIM will be larger than the above, since the DIM is spatially-correlated with HII regions and therefore probably sees a photoionizing flux that is higher than the disk-average. On the other hand, $f_{DIM}$ is of-order 10$^{- 1}$. The implied range in U is then consistent with photoionization models that show that the key diagnostic line ratios [NII]/H$\\alpha$ and [SII]/H$\\alpha$ peak for values of U in the range $3 \\times 10^{-5}$ to $3 \\times 10^{-4}$ (e.g. Sokolowski 1993). This analysis also reveals why the leaky HII region model would have difficulties explaining the {\\it disturbed} DIM. The lack of obvious broad wings on the H$\\alpha$ or [NII] emission-lines (\\S 3.2) implies that the disturbed DIM must have a relatively high ionization state. Photoionization models show that very strong [OIII]$\\lambda$5007 emission (e.g. [OIII]/H$\\beta$ $\\gg$ 1) requires logU $> -2$. Using the same arguments as in the above paragraph, values this large for U would require $n_{e} < 0.02 f_{DIM}$, or P/k $< 4 \\times 10^{2} f_{DIM}\\ $ cm$^{-3}$ K$^{-1}$. This pressure is far below plausible ISM pressures. While the hydrostatic equilibrium condition implies that the gas pressure will drop with distance out of the disk plane, the required pressures are so low that the disturbed DIM would have to be located more than 6 pressure-scale-heights above the inner disk (inconsistent with our estimates in \\S 4.2 below). Perhaps the biggest challenge for photoionization models is explaining the pervasiveness of the DIM. Our new long-slit data showing H$\\alpha$, [NII], and [SII] emission throughout the inner disk only exacerbates this problem. Here we would like to emphasize the curious fact that the opacity of the ISM to ionizing radiation in these galaxies seems to be `fine-tuned'. That is, we can imagine two extremes for the fate of ionizing radiation produced by massive stars. First, it may all be absorbed locally around the complexes of O stars so that a deep H$\\alpha$ image would reveal small intense `islands' of high-surface-brightness surrounded by very dark `oceans' that cover most of the disk. This is manifestly not the case (cf. Figures 1 and 2). The other extreme case is one in which the ISM is optically-thin to ionizing radiation, allowing it to freely escape into the galactic halo or intergalactic space. This also seems far from reality (\\cite{lei95,gia97}). Instead, the ISM seems to be arranged such that is optically thin enough to allow the gas throughout a significant fraction of the ISM to be bathed in the diffuse radiation of the hot stars, yet optically-thick enough to capture the great majority of the ionizing radiation. Why is this so? Is this the result of some subtle feedback loop between star-formation and the structure and physical state of the ISM? \\subsubsection{Mechanical Heating} While we can not totally rule out photoionization for the disturbed DIM, we are lead to consider a totally different ionization source. The disturbed kinematics of this gas could implicate mechanical heating of the ISM by supernovae and stellar winds. We therefore consider two alternative forms for such mechanical heating: radiative shocks and turbulent mixing layers. Models of low-density radiative shocks (Shull \\& McKee 1979; Raymond 1997) show that suitably strong [OIII]$\\lambda$5007 emission (e.g. [OIII]/H$\\alpha$ $=$ 1 to 3) is attained for shock speeds $>$ 100 km s$^{-1}$. A strong shock will accelerate the shocked material to 3/4 of the shock velocity, so we would expect the velocity dispersion in the DIM to be $>$ 75 km s$^{-1}$. The typical measured line widths for the [OIII]$\\lambda$5007 line imply a 3-dimensional velocity dispersion for the disturbed DIM of 60 to 100 km s$^{-1}$, so the required shock speeds are plausible. Turbulent mixing layers (Slavin, Shull, \\& Begelman 1993) are the interface between hot gas flowing past cold clouds, and hence have a temperature that is intermediate between the cloud and hot phases. The models yield values of [OIII]$\\lambda$5007/H$\\alpha$ of 1.5 to 3 for much of the parameter space these authors explored. This specific range encompasses mixing layers with logT $=$ 5.3 and 5.5 when the `cold' clouds have logT $=$ 4, and mixing layers with logT $=$ 5.5 when the cold clouds have logT $=$ 2. Cooler mixing layers do not produce significant [OIII] emission. The turbulent mixing layer models have one advantage over the shock models in that they do not demand that the mixing layer necessarily have a large velocity dispersion. Satisfactorily strong [OIII] emission relative to H$\\alpha$ can be produced when the hot gas moves past the cold cloud at either 25 km s$^{-1}$ or 100 km s$^{-1}$ (consistent with the observed [OIII] line widths in the DIM). High relative velocities do lead to enhanced emissivity at a given gas pressure however (see below). A potential problem with both the shock and turbulent mixing layers models (especially the latter) is the low predicted [OIII]$\\lambda$5007 surface brightness for gas having the modest pressures appropriate to a normal disk galaxy ISM (P/k $\\sim$ 10$^{4}$ K cm$^{-3}$ to $\\sim$ 10$^{5}$ K cm$^{-3}$ in the inner disk regions probed by our spectra). For P/k $=$ 10$^{5}$ K cm$^{-3}$, the predicted [OIII]$\\lambda$5007 surface brightnesses from shocks with v = 100 to 150 km s$^{-1}$ are in the range $3 \\times 10^{-18}$ to $7 \\times 10^{-18}$ erg cm$^{-2}$ s$^{-1}$ arcsec$^{-2}$. This is similar to typical [OIII]$\\lambda$5007 surface-brightnesses in the DIM of about $5 \\times 10^{-18}$ to $1 \\times 10^{-17}$ erg cm$^{-2}$ s$^{-1}$ arcsec$^{-2}$. This agreement would then require the areal covering-factor of shocked gas within our slit to be unity. However, the {\\it volume filling-factor} of shocked gas in the galaxy ISM would be much smaller. The shock models of Shull \\& McKee (1979) scaled to P/k $=$ 10$^{5}$ K cm$^{-3}$ imply that the total thickness of the layer of [OIII]-emitting gas would be of-order 0.1 pc per line-of-sight through the DIM (much less than the thickness of the disk ISM). The surface-brightness agreement is not so good with the models for turbulent mixing layers. These predict [OIII]$\\lambda$5007 surface brightnesses of 10$^{-18}$ erg cm$^{-2}$ s$^{-1}$ arcsec$^{-2}$ at best, even for P/k $=$ 10$^{5}$ K cm$^{-3}$. This is an order-of-magnitude smaller than the typical observed values in the DIM, and would require the presence of many ($>$ 10) turbulent mixing layers along a typical line-of-sight through the galaxy disk. As noted in \\S 1, energizing the entire DIM seems problematic using only the mechanical energy from supernovae and stellar winds, since it requires tapping essentially {\\it all} of this mechanical energy. However,the energetic requirements for ionizing the disturbed DIM alone are substantially less. That is, the disturbed DIM contributes less than 20\\% to the total H$\\alpha$ luminosity of the DIM (\\S 3.2.2), at least in the regions probed by our spectra. Thus, mechanical heating can not be ruled out on simple energetic grounds. We have seen that the DIM is very pervasive in the inner disks of our galaxies, so mechanical heating of the DIM would require that nearly every line-of-sight through the disk intersect regions of mechanically-heated gas. It is not clear whether the topology or `porosity' (McKee \\& Ostriker 1977) of the ISM in typical disk galaxies is consistent with this requirement. Observationally, there is no direct evidence for a high covering fraction of diffuse gas mechanically heated to a temperature of $10^5$~K or more in typical galactic disks. For example, Heiles (1990) and Oey \\& Clarke (1997) have concluded that only a small fraction of the surface area of the disks of the Milky Way, M 31, and M 33 is covered by the hot superbubbles created by supernovae and stellar winds. On the other hand, the available information about soft X-ray emission from the nearby spiral galaxies suggests the existence of diffuse, relatively hot gas (T $\\ge10^6$~K), but little is known about the corresponding areal covering factor. Perhaps the best studied case of diffuse, soft X-ray emitting gas in our sample galaxies is M 101. Snowden \\& Pietsch (1995) have shown that in the inner disk of M 101, the gas has T$\\simeq 10^{5.8}$~K and a covering fraction of-order unity. To heat up the gas to this temperature requires shock speeds ($\\sim 200$ km s$^{-1}$) that are much larger than the typical line widths that we have observed in the DIM. Certainly, the hot gas seen in the X-ray data may cool and fall back into the disk thereby contributing to the DIM. \\subsection{The Dynamics and Inferred Vertical Extent of the DIM} The observed H$\\alpha$ and [NII]$\\lambda$6584 line widths in the DIM range from about 30 to 100 km s$^{-1}$ (FWHM), rather similar to the line widths in the Reynolds Layer in our own Galactic disk of $\\sim$ 30 to 60 km s$^{-1}$ (Reynolds 1985). Since we have verified that there are no significant velocity shears along the slit in the regions where we have made these line-widths measurements, the broadening must be due to small-scale `turbulent' motions of the gas. The lines are certainly broader than pure thermal broadening (which is only 22 km s$^{-1}$ FWHM for HI at T $= 10^{4}$ K). Of course, the [OIII]$\\lambda$5007 line-widths are larger still: typically 70 to 150 km s$^{-1}$ FWHM. We can use these line-widths to estimate the corresponding scale-height of the emitting gas, assuming that the gas clouds have an isotropic velocity dispersion and move in a gravitational potential in which they act as test particles. Using optical surface photometry of the disks of our sample galaxies, and taking mass-to-light ratio M$_{tot}$/L = 5 solar units for both V and B bands (e.g. Binney \\& Tremaine 1987) we estimate that the typical surface mass-density in the inner disks ranges from 100 to 400 M$_{\\odot}$ pc$^{-2}$. We further assume that the mass that sets the potential has the form $\\rho$(z) $\\alpha$ exp[-z/H] where z is the distance out of the mid-plane and H is the scale-height for the mass $\\rho$. In the disk of the Milky Way, H $\\sim$ 350 pc (Freeman 1987). Adopting this value, we have then solved for the scale height h$_{gas}$ for an ensemble of gas clouds moving in this potential with a velocity dispersion in the z-direction $\\sigma_{z} = 0.43 \\times$ (FWHM). The results are listed in Table 7. From this we see that typical scale-heights for the quiescent DIM (based on H$\\alpha$ and [NII]$\\lambda$6584 line widths) are 300 to 500 pc, which is significantly less than the values of 600 to 900 pc for the Reynolds Layer (Reynolds 1993). This is due in part to the much larger surface mass-densities in the inner regions of our galaxies and the associated deeper potential well. The estimated scale-heights for the dominant (quiescent) component of the DIM are in fact typically comparable to our assumed scale-height for the stellar disk (350 pc). Thus, we conclude that {\\it the DIM in the inner regions of our galaxies is not in any sense an `extraplanar' phenomenon}. This is interesting, since the terms `DIM' (or `DIG' or `WIM'!) and `extraplanar gas' were often used synonymously in early investigations. In fact, the quiescent state we find for the dominant component of the DIM may naturally explain why the DIM seems to be ubiquitous in late-type galaxies, but striking examples of extraplanar gas are proving to be rare (e.g. \\cite{ran96conf,ran96}). In Table 7 we also list the derived scale-heights for the disturbed DIM using our measured [OIII]$\\lambda$5007 line widths. These scale-heights are typically 400 to 800 pc, or a factor of about 1.5 to 2 greater than for the quiescent DIM. Thus the disturbed DIM might more properly be considered `extraplanar'. In this regard it is intriguing that Rand (1997) found morphological evidence for two components in the DIM in NGC 891. Since this galaxy is viewed almost exactly edge-on, Rand fit the vertical dependence of the mean H$\\alpha$ surface-brightness in terms of the sum of two exponentials. There is a bright component that provides about 84\\% of the total emission and has a scale-height of 500 pc, and a faint component that provides the other 16\\% of the emission and has a scale height 5 to 6 times larger. The bright and faint components must have different kinematics to have such different vertical distributions, and we suggest that they may correspond respectively to the quiescent and disturbed components of the DIM that we have identified kinematically. It would be important in this regard to measure the dependence of the [OIII]$\\lambda$5007/H$\\beta$ ratio on distance out of the mid-plane in NGC 891 to see in the more extended DIM component has a relatively high ionization state (as we find for the disturbed DIM). The only published [OIII] measurement (Dettmar 1992) is an upper limit at a distance of only 0.5 kpc out of the disk (in the region dominated by Rand's inner DIM component, which we would predict to be of low ionization)." + }, + "9708/hep-ph9708420_arXiv.txt": { + "abstract": "A cubic kilometer scale experiment has been proposed to detect cosmic neutrinos of energy from tens of GeV up to the highest energies observed for cosmic rays, $\\sim10^{20}\\,$eV, or possibly even beyond. Detection efficiencies depend crucially on the neutrino-nucleon cross section at these energies at which radiative corrections beyond the lowest order approximation could become non-negligible. The differential cross sections can be modified by more than $50\\%$ in some regions of phase space. Estimates of corrections to the quantities most relevant for neutrino detection at these energies give, however, less dramatic effects: The average inelasticity in the outgoing lepton is increased from $\\simeq0.19$ to $\\simeq0.24$. The inclusive cross section is reduced by roughly half a percent. The dominant uncertainty of the standard model ultra-high energy neutrino-nucleon cross section therefore still comes from uncertainties of the parton distributions in the nucleon at very low momentum fractions. ", + "introduction": "Several proposals have recently been put forward for the search for cosmic neutrinos above tens of GeV up to the high energy end of the cosmic ray spectrum, and possibly beyond. The most well developed technique is to detect Cherenkov light from the muon produced in a charged-current reaction of an ultra-high energy (UHE) muon neutrino with a nucleon in either ice or water. Several prototype detectors based on this technique have been constructed, namely DUMAND at a depth of nearly $5\\,$km in the ocean near Hawaii (now out of commission), the neutrino telescope at Lake Baikal about $1\\,$km deep, NESTOR at about $3.5\\,$km depth in the Mediterranean near Greece, and AMANDA up to $2\\,$km deep in South Pole ice~\\cite{ghs}. Other techniques have been proposed such as detection of horizontal air showers~\\cite{ydjs,bpvz}, or detection of acoustic~\\cite{acoustic} or radio waves~\\cite{radio} associated with the neutrino induced cascades. For a given neutrino flux, the efficiency of all these methods depends predominantly on the neutrino-nucleon interaction cross section. To lowest order in the electroweak (EW) coupling, this cross section has been discussed in detail in the literature, see, e.g., Refs.~\\cite{fkr,gqrs} for the most recent work. Due to uncertainties in the extrapolation of quark distribution functions in the nucleon to very small fractional momentum transfers, $x\\lesssim10^{-4}$, and large (negative) squared four-momentum transfers, $Q^2\\gtrsim10^5\\,{\\rm GeV}^2$, best estimates for energies around $10^{20}\\,$eV (in the laboratory frame) vary by factors of a few. On the other hand it is well known that higher order processes can become important or even dominant for electromagnetic (EM) interactions at UHE. For example, energy exchange between two leptons $l_1$ and $l_2$ is dominated by EM bremsstrahlung, $l_1+l_2\\rightarrow l_1+l_2+\\gamma$, for center-of-mass energies $s$ exceeding the square of the electron mass, rather than by ordinary Mott scattering, $l_1+l_2\\rightarrow l_1+l_2$. The relevant cross section rises with the logarithm of $s$. EW radiative corrections to deep inelastic neutrino-nucleon scattering have been calculated before in the literature: Ref.~\\cite{rpsn} contains a discussion of the leading log approximation for which only corrections involving photons are relevant. These corrections contain a factor $\\ln(s/m_l^2)$ where $m_l$ is the charged lepton mass and also depend on the behavior of the parton distribution functions. Single and double differential cross sections have been evaluated in Ref.~\\cite{rpsn} for neutrino energies $E_\\nu$ up to a few hundred GeV, using rough estimates of the parton distributions available back then. To compare with HERA measurements~\\cite{hera} at $s\\simeq10^5\\,{\\rm GeV}^2$, corresponding to $E_\\nu\\simeq50\\,$TeV, more recently complete analytical expressions have been presented in Refs.~\\cite{bbcr,spiesberger}. At these energies, corrections have been shown to be up to $50\\%$ in certain areas of phase space. Consequently, at energies approaching $10^{20}\\,$eV, radiative corrections could be larger still and may play an important role. We therefore found it worthwhile to extend estimates of radiative corrections to UHE, using updated parton distribution functions and with a special emphasize on the quantities relevant for UHE neutrino detection. The rest of this paper is organized as follows: In Sec.~II we introduce our estimates of radiative corrections at UHE. In Sec.~III we present numerical results. We finally summarize our findings and resulting consequences in Sec.~IV. ", + "conclusions": "As can be seen from Fig.~\\ref{F2}, at energies around $10^{20}\\,$eV the radiative corrections to the single differential charged-current neutrino-nucleon cross section $(d\\sigma/dy)(E_\\nu,y)$ are negative for $y\\lesssim0.02$ and positive otherwise. For $y\\lesssim10^{-3}$ they grow larger than 50\\%, whereas for $y\\gtrsim0.2$ they are of the order of 30\\%. The average inelasticity $\\left\\langle y\\right\\rangle$ is increased from $\\simeq0.19$ to $\\simeq0.24$ whose potential influence on the development of the neutrino induced cascade is probably the strongest effect of radiative corrections. The corrections to the total cross section are negative and roughly constant at about half a percent (see Fig.~\\ref{F3}). Apart from physics beyond the standard model (see, e.g., Refs~\\cite{bkklz,bhfpt}, uncertainties of the UHE neutrino-nucleon cross section to date therefore are by an ample margin dominated by the uncertainties in the parton distributions in the nucleon at very low momentum fractions." + }, + "9708/astro-ph9708244_arXiv.txt": { + "abstract": "We present spectra of three M subdwarfs which are common proper motion companions to F or G subdwarfs of known metallicity. The assumption that the companions have the same composition allows us to test the Gizis (1997, \\aj, 113, 806) M subdwarf classification system and its correspondence to metallicity. The results are in excellent agreement with the Gizis (1997) scale, thereby showing that the Allard \\& Hauschildt (1995, \\apj, 445, 433) Extended model atmospheres agree well in the 6200 -- 7400$\\AA$ region for cool metal-poor stars. We also show that the results are consistent with the main sequences of globular clusters using the Reid (1997, \\aj, 114, 161) distance scale. ", + "introduction": "} The metal-poor stars of the thick disk and halo provide an invaluable record of Galactic history. The main-sequence FGK subdwarfs have proven to be an important source of information on these populations (e.g., \\cite{carney}). The much cooler and fainter M subdwarfs offer an important alternative tracer group. Indeed, in addition to the possibility of observing nearby proper motion M subdwarfs, it is now feasible to obtain both photometry for the M subdwarfs in globular clusters with the Hubble Space Telescope (e.g., \\cite{seg96}) and spectra for M dwarfs and M subdwarfs at distances of a few kiloparsecs above the galactic plane with 10-meter class telescopes (\\cite{keck}). Using these objects as probes of Galactic structure, however, requires a good understanding of their properties in order to derive metallicities and luminosities. Gizis (1997, hereafter G97) has presented a spectroscopic classification scheme which is based on moderate resolution ($\\sim 3 \\AA$) spectra covering the wavelength range 6200 -- 7400 $\\AA$. Quantitative bandstrength indices measuring TiO and CaH features are used to classify stars as M V (ordinary disk stars), sdM (M subdwarfs), and esdM (extreme M subdwarfs). Comparison to the Allard \\& Hauschildt (1995) synthetic spectra allowed G97 to show that these classes correspond to $[m/H] \\sim 0.0$, $[m/H] \\sim -1.2 \\pm 0.3$, and $[m/H] \\sim -2.0 \\pm 0.5$ respectively. Comparison of the ($M_V,V-I$) HR diagram shows that HST globular cluster sequences (\\cite{seg96}) and stellar interior calculations with Allard \\& Hauschildt (1995) model atmospheres (\\cite{bcah95}) are in agreement with this scale; however, serious systematic errors could in principle affect all of these methods of estimating metallicity. We present spectra of three M subdwarfs which are companions to hotter subdwarfs of known metallicity. Our aim is to test the metallicities derived from the M subdwarf spectra by comparison with those measured for their better understood primaries. The data are presented in Section~\\ref{data}, the implications for the metallicity scale are discussed in Section~\\ref{discuss}, and the results are summarized in Section~\\ref{conclusions}. ", + "conclusions": "} We have compared the metallicities estimated directly from spectra of three M subdwarfs to the metallicities derived for their FGK subdwarf companions. We find that the metallicities based on the Gizis (1997) spectroscopic classification system are consistent with the metallicities derived by Carney et al. (1994) from high-resolution spectra. We argue that $[m/H]$ on the G97 scale corresponds to $[m/H] \\approx [Fe/H]$." + }, + "9708/gr-qc9708035_arXiv.txt": { + "abstract": "The standard approach to the numerical evolution of black hole data using the ADM formulation with maximal slicing and vanishing shift is extended to non-symmetric black hole data containing black holes with linear momentum and spin by using a time-independent conformal rescaling based on the puncture representation of the black holes. We give an example for a concrete three dimensional numerical implementation. The main result of the simulations is that this approach allows for the first time to evolve through a brief period of the merger phase of the black hole inspiral. ", + "introduction": "While the binary black hole problem is essentially unsolved, there have been many advances in our understanding of the general relativistic evolution problem. The situation of interest is how two black holes orbit each other, spiral inwards, eventually merge, and emit gravitational waves in the process. Gravitational wave detectors are being built that eventually will detect waves from such sources \\cite{detect}. Although at least as interesting from an astrophysical as opposed to purely general relativistic view point, here we will not consider neutron stars or other matter sources (e.g.\\ \\cite{matter}). To name just three areas of research related to the inspiral of two black holes: post-Newtonian methods are applicable if the system is not too strongly general relativistic \\cite{post-newt}, that is in the early phase of the inspiral; perturbation methods can model the final stages of the merger \\cite{perturb}; and full numerical relativity can be applied to the late strong-interaction phase if one imposes axisymmetry \\cite{Sm79,axisym}. What may appear surprising is that the general, three-dimensional two black hole problem has not been solved {\\em numerically} to some degree. After all, there has been a lot of theoretical work on the Einstein equations, and even though they are complicated, one might expect that eventually computers become powerful enough to at least allow some rough investigations into the general problem. However, general relativity has at least two characteristic features, gauge freedom and singularity formation, which would prevent us from solving the problem completely with current computer codes even if an infinitely fast computer became available. Here gauge freedom refers to the freedom to choose coordinates, which in numerical relativity is a hard problem since in general one cannot specify coordinates for the entire spacetime in advance. For an adequate representation of the spacetime on a numerical grid one usually has to construct the coordinates numerically as one of the steps in the numerical evolution scheme. In this dynamic construction of coordinates one has to avoid the formation of coordinate singularities, but even if these are absent, it is still possible that completely regular initial data develops physical singularities, as is the case for typical black hole spacetimes. Both problems, the choice of coordinates and the representation of black hole spacetimes on a numerical grid, are still awaiting a completely satisfactory solution. What has been achieved in 3d numerical relativity is the evolution of weak gravitational waves \\cite{wa3d1,wa3d2} (non-linear waves that are weak enough to not collapse to black holes), and the evolution of a single spherically symmetric black hole \\cite{ss3d,BrAM,daues} in Cartesian coordinates and for non-trivial slicings. Recently, long term stable evolutions have been achieved for single black holes \\cite{stable}. A single moving black hole has been simulated in \\cite{move}, which should become important for black hole collisions. Axisymmetric collisions of two black holes with a 3d code have been reported in \\cite{mi3d}. Non-axisymmetric spacetimes containing a single distorted black hole are studied in \\cite{ss3ddist}. Various other 3d projects are actively pursued, most notably by the US binary black hole alliance \\cite{grand}. In this paper we introduce a new approach that makes the study of non-axisymmetric binary black hole spacetimes possible, at least for short time intervals. The crucial new insight is that the general binary black hole initial data of \\cite{BrBr} can be evolved on $R^3$ without special inner boundary using a conformal rescaling that is constant in time. The main result is the evolution of the apparent horizon, which for appropriate black hole initial data has two components initially that are replaced by a single component during evolution. The main limitation of the current implementation is that only a brief interval of the merger phase of the inspiral of two black holes can be studied. This is due to the well-known problem of grid-stretching associated with our particular gauge choice, maximal slicing and vanishing shift. We outline below how the gauge conditions can be generalized. Here we want to demonstrate our approach with one minimal example for general black hole data, a more detailed study is left to a future publication. To summarize, initial data for two black holes of unequal mass that are moving and spinning are constructed following \\cite{BrBr}. The evolution is carried out in the 3+1 ADM formulation with maximal slicing and vanishing shift using BAM, a bifunctional adaptive mesh code for coupled elliptic and hyperbolic systems (only a fixed mesh refinement of nested grids is used here; see \\cite{BrAM,BrBr,BrBAM} about earlier versions of this code). An evolution time of about 7--10$M$ (in units of the ADM mass) is achieved. We can study the geometric information in the collapse of the lapse function, and we find the apparent horizon and show data for the transition from two outermost marginally trapped closed surfaces to a single one. In the following, we describe the method and its implementation, present the physical data, and conclude with a discussion. ", + "conclusions": "In conclusion, we have shown how the standard approach of ADM evolution with maximal slicing and vanishing shift can be applied to non-symmetric black hole data containing black holes with linear momentum and spin by using a time-independent conformal rescaling based on the puncture representation of the black holes. We discuss an example based on a concrete three dimensional numerical implementation. The main result of the simulations is that this approach allows for the first time to evolve through a brief period of the merger phase of the black hole inspiral. Looking back to the early numerical work of Smarr and Eppley on axisymmetric black hole collisions in the seventies \\cite{Sm79}, we feel that it is useful to point out what concretely can be done about the two black hole problem today, even if it is just a first step. Important issues for further investigations are extending the run time so that wave extraction becomes possible, non-vanishing shift conditions, and the implementation of an apparent horizon boundary condition. These issues will be addressed in a new collaborative computational framework, the Cactus code \\cite{cactus}, which provides the infrastructure for input-output and MPI (The Message Passing Interface) based parallelism, and which allows easy code-sharing among several authors. For example, Cactus includes various evolution modules, initial data set constructions, and analysis modules for wave extraction and apparent horizon finding \\cite{AHF:cactus,cactus}. In particular, the ADM evolution routines and the multigrid elliptic solver of BAM have already been ported to Cactus, adaptive mesh refinement is still under development. An important and immediate application of this new, more powerful framework will be the comparison of the evolution of axisymmetric black hole data in 3d with results obtained with 2d codes \\cite{axisym}. It is a pleasure to thank J. Ehlers, C. Gundlach, P. H\\\"ubner, B. Schutz, E. Seidel, P. Walker, and especially S. Brandt and B. Schmidt, for their support and many stimulating discussions. The computations were performed at the AEI in Potsdam. \\newcommand{\\bb}{B. Br\\\"ugmann} \\newcommand{\\bib}[1]{\\bibitem{#1}} \\newcommand{\\EM}{} \\newcommand{\\apny}[1]{{\\EM Ann.\\ Phys.\\ (N.Y.) }{\\bf #1}} \\newcommand{\\cjm}[1]{{\\EM Canadian\\ J.\\ Math.\\ }{\\bf #1}} \\newcommand{\\cmp}[1]{{\\EM Commun.\\ Math.\\ Phys.\\ }{\\bf #1}} \\newcommand{\\cqg}[1]{{\\EM Class.\\ Quan.\\ Grav.\\ }{\\bf #1}} \\newcommand{\\grg}[1]{{\\EM Gen.\\ Rel.\\ Grav.\\ }{\\bf #1}} \\newcommand{\\jgp}[1]{{\\EM J. Geom.\\ Phys.\\ }{\\bf #1}} \\newcommand{\\ijmp}[1]{{\\EM Int.\\ J. Mod.\\ Phys.\\ }{\\bf #1}} \\newcommand{\\JCP}[1]{{\\EM J. Comp.\\ Phys.\\ }{\\bf #1}} \\newcommand{\\jmp}[1]{{\\EM J. Math.\\ Phys.\\ }{\\bf #1}} \\newcommand{\\mpl}[1]{{\\EM Mod.\\ Phys.\\ Lett.\\ }{\\bf #1}} \\newcommand{\\np}[1]{{\\EM Nucl.\\ Phys.\\ }{\\bf #1}} \\newcommand{\\PL}[1]{{\\EM Phys.\\ Lett.\\ }{\\bf #1}} \\newcommand{\\pr}[1]{{\\EM Phys.\\ Rev.\\ }{\\bf #1}} \\newcommand{\\PRL}[1]{{\\EM Phys.\\ Rev.\\ Lett.\\ }{\\bf #1}}" + }, + "9708/astro-ph9708134_arXiv.txt": { + "abstract": "Creation of strange quark stars through strong interaction deconfinement is studied based on modern estimates of hyperon formation in neutron stars. The hyperon abundance is shown to be large enough so that if strange quark matter (SQM) is the true ground state of matter, the deconfinement density should be at most $2.5\\!-\\!3$ times the nuclear saturation density. If so, deconfinement occurs in neutron stars at birth, and {\\it all} neutron stars must be strange quark stars. Alternatively, sould observation indicate that some neutron stars have a baryonic interior, SQM is unlikely to be absolutely stable. ", + "introduction": " ", + "conclusions": "" + }, + "9708/astro-ph9708120_arXiv.txt": { + "abstract": "We propose a method to remove the mass sheet degeneracy that arises when the mass of galaxy clusters is inferred from gravitational shear. The method utilizes high-redshift standard candles that undergo weak lensing. Natural candidates for such standard candles are type Ia supernovae (SN Ia). When corrected with the light-curve shape (LCS), the peak magnitude of SN Ia provides a standard candle with an uncertainty in apparent magnitude of $\\Delta m\\simeq 0.1-0.2$. Gravitational magnification of a background SN Ia by an intervening cluster would cause a mismatch between the observed SN Ia peak magnitude compared to that expected from its LCS and redshift. The average detection rate for SN Ia with a significant mismatch of $\\ge2\\Delta m$ behind a cluster at $z\\simeq0.05-0.15$ is about $1-2$ supernovae per cluster per year at $J,I,R\\lesssim25-26$. Since SNe are point-like sources for a limited period, they can experience significant microlensing by MACHOs in the intracluster medium. Microlensing events caused by MACHOs of $\\sim10^{-4}\\,{\\rm M}_\\odot$ are expected to have time scales similar to that of the SN light curve. Both the magnification curve by a MACHO and the light curve of a SN Ia have characteristic shapes that allow to separate them. Microlensing events due to MACHOs of smaller mass can unambiguously be identified in the SN light curve if the latter is continuously monitored. The average number of identifiable microlensing events per nearby cluster ($z\\lesssim0.05$) per year is $\\sim 0.02 \\,(f/0.01)$, where $f$ is the fraction of the cluster mass in MACHOs of masses $10^{-7} < {\\rm M_{\\rm macho}}/{\\rm M}_\\odot < 10^{-4}$. ", + "introduction": "\\label{sec:intro} The mass sheet degeneracy (Falco, Gorenstein, \\& Shapiro 1985; Schneider \\& Seitz 1995) constitutes one of the fundamental uncertainties in attempts at reconstructing galaxy cluster masses from the gravitational distortion of resolved background sources. It arises because the observed gravitational shear field is insensitive to magnification. A few methods were proposed to break the degeneracy by directly measuring either the magnification or the (scaled) surface mass density. Broadhurst, Taylor, \\& Peacock (1995) suggested to compare the redshift and/or the magnitude distribution of the background sources in cluster fields with those measured in empty fields. Kaiser (1995) and Kneib \\etal (1995) showed that inferring the mean source redshift at a given apparent magnitude can help to break the degeneracy. Bartelmann \\& Narayan (1995) proposed and tested an algorithm that exploits the size distribution of the background sources and the conservation of surface brightness under gravitational lensing. An alternative cluster reconstruction technique suggested by Bartelmann \\etal (1996) breaks the mass sheet degeneracy by simultaneously taking sizes {\\em and\\/} shapes of background galaxies into account. All these methods rest upon statistical information on intrinsic background galaxy sizes. They require accurate measurements of sizes and magnitudes or surface brightnesses of lensed galaxies. In principle, a single sufficiently precise measurement of the magnification of a standard candle by a cluster can lift the mass sheet degeneracy as well. Suppose a standard candle is observed behind a lens, and its redshift is measured. Assuming a cosmological model, the luminosity distance is known. The apparent magnitude expected from its luminosity can then be compared to the observed one, and any discrepancy between the two can be attributed to the magnification by the foreground lens. Ideal standard candles are rare in cosmology. Nevertheless, there is accumulating strong evidence that SN Ia may well serve as approximate standard candles, especially if one takes the dependence of their peak magnitude on the light curve shape into account (Phillips 1993; Riess, Press, \\& Kirshner 1995 (RPK1); Riess, Press, \\& Kirshner 1996 (RPK2); Hamuy \\etal 1996a). Moreover, SN Ia are ubiquitous, point-like sources, their distinct spectral lines allow for an accurate redshift determination, and as of yet they show no cosmological evolution. Though rare in frequency, each of their aforementioned characteristics makes SN Ia excellent probes for the integrated matter density along their lines-of-sight, and the composition of the matter (\\ie its graininess, occurrence of MACHOs etc.). Weak lensing by large-scale structure was shown to not significantly affect the SN magnitude (Frieman 1996; Wambsganss \\etal 1996), but this does not apply to galaxy clusters. Similar arguments led Kovner \\& Paczy\\'nski (1988) to propose a possible identification of SNe in giant arcs as a tool to establish the time delay between different images. Giant arcs are less appealing than weak lensing by clusters because they are much more rare, and sensitive to the details of the density profile close to the cluster core. Soucail \\& Fort (1991) suggested the combination of giant arcs and a much less accurate ``standard candle'', viz.\\ the Tully-Fisher relation, to estimate the Hubble constant. This method suffers from several difficulties, \\eg a possible cosmological evolution of the Tully-Fisher relation, and the problematic determination of the inclination angle of the lensed galaxy. In the rest of this paper, we present the basic considerations one has to take into account in order to employ SN Ia for these purposes. In \\S\\ref{sec:sn_cl} we review the necessary facts about the SN population and the cluster model we use. In \\S\\ref{sec:esti} we estimate the number of expected SN in the background of a cluster, whose luminosity--distance mismatch is large enough to allow for the removal of the mass sheet degeneracy. In \\S\\ref{sec:machos} we discuss the use of a lensed SN Ia as a probe for MACHOs in the intracluster medium. We conclude with our results in \\S\\ref{sec:conc}. ", + "conclusions": "\\label{sec:conc} We have proposed to look for SN type Ia behind nearby rich clusters. The expected rate of detectable SN behind such clusters is of the order of $~10$ SN yr$^{-1}$ per cluster, about 20\\% of which can be used alone as tools to remove the mass sheet degeneracy. The rest of the detected SN can be combined together for the same purpose. Since the magnification scales like $\\theta^{-1}$ and the number of sources scales like $\\sim \\theta^2$ (neglecting magnification bias), we expect the signal-to-noise ratio to be constant. Statistical use of the SNe on large radii yields a similar signal-to-noise ratio as for the SNe that exceeded the mismatch threshold. Any of the detected SN can serve as a point source with a typical light curve shape. Any modulation of this light curve shape, or lack thereof, can put limits on the mass fraction of the cluster in the form of $10^{-7}<{\\rm M}/{\\rm M}_\\odot<10^{-4}$ MACHOs. This proposed survey shifts the focus from high redshift clusters as targets for high redshift SNe to low redshift clusters as foreground lenses for the same, and farther SNe. At the same time, the nearby observed cluster enhances the probability of finding serendipitous SNe in it. The use of the cluster SNe as point sources for MACHO search is somewhat less efficient because of the relatively small effective lensing distance $D_{\\rm eff}$, and the unknown position of the source within the cluster. The monitoring procedure of current observations will have to be modified, namely to allow for longer gaps between successive observations of the same cluster, but more observations in total during the year. The proposal demands a higher magnitude limit than the values that are currently used, both for the identification and for the follow-up. These limits, however, are not out of reach. Let us list again all the choices we have made to result in an underestimate of the mismatch SNe rate. \\begin{enumerate} \\item We do not take into account the fainter image when multiple images appear on the image plane. \\item We took the specific SN Ia rate at $z=0.4$ and used it as the representative value for higher redshifts despite the prediction of increasing rate with redshift. \\item We choose a $\\delta$-function to represent the probability distribution of the SN absolute blue magnitude. The $\\delta$-function is centered on the mean, even though the median of the distribution seems to be on the bright side. On the other hand, note we have ignored a possible Malmquist bias in the Hamuy \\etal (1996a) sample, a bias that tends to slightly shift the average towards brighter magnitudes. \\item We considered only SNe that exceed the detection (magnitude limit) threshold $5$ days (rest-frame) prior to peak luminosity. A weaker requirement (\\ie only $0.2$ magnitudes below peak) may suffice if a denser time sampling is devised, or if the SN is intrinsically more luminous than average (and therefore rises and declines more slowly). \\item We require a mismatch of twice the natural scatter in the LCS-luminosity relationship. A more modest requirement may be sufficient to establish limits on the surface mass density along the line-of-sight. \\item We only integrate in redshift up to $z=2$ because of lack of knowledge of the number density of galaxies beyond this range. Note that if we are not willing to extrapolate the galaxy number density beyond $z\\simeq1.5$, but we still like to satisfy the observed surface number density of galaxies, the number density should rise even more steeply than assumed beyond $z=1$. If this is the case, more SN Ia fall under the magnitude limit of the survey, and the rate increases correspondingly. \\end{enumerate} It seems that the main challenge for this proposed method is the high accuracy photometry needed. Two groups are currently conducting a hunt for high redshift SNe (the SN cosmology project based at Lawrence Berkeley National Laboratory (Perlmutter \\etal 1997); the high-z SN search team based in Mt. Stromolo, Australia (Schmidt \\etal 1996)). One survey fits this paper proposal if higher magnitude limits are applied (Abell cluster SN search, University of Washington). These surveys' objective is generally quite different from the investigation of galaxy clusters (nevertheless, we have just heard that the SN cosmology project intend to propose a survey in the spirit of this paper). The expertise they have acquired for determining the value of the cosmological background parameters should serve them well in pursuing the nature and distribution of matter within this background." + }, + "9708/astro-ph9708250_arXiv.txt": { + "abstract": "Comparing the frequency of typical events with that of unusual events allows one to test whether the cosmological density distribution function is consistent with the normally made assumption of gaussianity. To this end, we compare the consistency of the tail-inferred (from clusters) and measured values (from large-scale flows) of the rms level of mass fluctuations for two distribution functions: a Gaussian, and a texture (positively-skewed) PDF. We find that if we average the recent large-scale flow measurements, observations of the rms and the tail at the $10\\hmpc$\\ scale disfavor a texture PDF at $\\sim 1.5\\sigma$ in all cases. If we take the most recent measurement of the rms, that from Willick \\etal\\ (1997b), the comparison disfavors textures for low $\\Omega_0= 0.3$, and disfavors Gaussian models if $\\Omega_0= 1$ (again at $\\sim 1.5\\sigma$). Predictions for evolution of high temperature clusters can also be made for the models considered, and, as is known (e.g., Henry 1997), strongly disfavor $\\Omega_0 = 1$ Gaussian models, while we find $\\Omega_0 = 1$ marginally disfavored in texture models. Taking the suite of tests as a whole, and using all of the quoted data, it appears that textures are strongly disfavored and only the low $\\Omega_0$ Gaussian models are consistent with all the data considered. But given evidence for the internal inconsistency of the observational data, had we only used the recent Willick \\etal\\ results, the strength of our conclusion would have been reduced to the $\\sim 1\\sigma$ level. ", + "introduction": "All current scenarios for the formation of structure in the Universe assume that cosmic structures grow from initially small seed perturbations. A fundamental assumption underlying a large class of these theories is that the initial Fourier modes have random phases. This implies that the statistics of the initial density field $\\rho_i({\\mathbf{x}})$ are fully specified by the correlation function $\\xi_i(r)=\\langle\\rho_i({\\mathbf{x}})\\rho_i({\\mathbf{x}}+ {\\mathbf{r}})\\rangle_{\\mathbf{x}}$, or equivalently its Fourier transform, the power spectrum $P_i(k)$. For any given smoothing length, then, the one-point probability distribution function (PDF) of the initial density field $\\rho_i$ is Gaussian. A measurement of the frequency of rare events (clusters of galaxies) coupled with a measurement of the level of rms fluctuations can tell us directly whether or not the distribution of density fluctuations satisfy the Gaussian hypothesis. While theories based on the inflationary paradigm generally predict Gaussian initial conditions, there are many alternatives which do not. Global textures and other topological defects, for example, predict a strongly non-Gaussian distribution for the initial perturbations (e.g., \\cite{gst91}; \\cite{pst91}; \\cite{gooding92}; \\cite{pen94}; \\cite{penetal97}). Global textures arise when a global non-Abelian symmetry, such as SU(2), is spontaneously and completely broken. When these defects collapse, energy gradients in the Higgs field perturb the metric and induce gravitational fluctuations in the matter fields. These texture knots, then, act as ``causal seeds'' of structure formation. We take the texture model as our standard non-Gaussian model, since it is well studied and still viable as a theory of structure formation. In general, texture and Gaussian models give similar predictions for the \\emph{shape} of the power spectrum and the level of CMB fluctuations (\\cite{pen94}; \\cite{pen95}). In fact, it is often noted (e.g., \\cite{pst91}) that texture and Gaussian CDM models are \\emph{most} different in their PDFs. If it could be demonstrated that the initial density field were Gaussian on some suitable scale, the texture model would be directly falsified. The standard method for testing Gaussianity relies on measuring the $N$-th order reduced moments, or cumulants $\\lambda_N$ of the present-day PDF, such as skewness or kurtosis, from redshift or peculiar velocity surveys (cf., Strauss \\& Willick 1995 for a review). The variance is $\\sigma^2 \\equiv \\lambda_2\\equiv \\langle\\delta^2\\rangle$, the skewness is $\\lambda_3\\equiv\\langle\\delta^3\\rangle$, and the kurtosis is $\\lambda_4\\equiv\\langle\\delta^4\\rangle - 3\\sigma^4$, where $\\delta$ is the density contrast $\\rho/\\langle\\rho\\rangle - 1$. For a Gaussian random field, all the cumulants of order $N>2$ are equal to zero. The growth of Gaussian initial perturbations under gravitational instability naturally produces non-Gaussianity in the distribution function of the density field. In particular, the cumulants follow a scaling relation $\\lambda_N = S_N \\lambda_2^{N-1}$, where the $S_N$ are independent of scale in the mildly non-linear regime for power-law power spectra (Fry 1984ab, Bernardeau 1992), and can be calculated from $(N-1)$-th order perturbation theory (Bernardeau 1994). Significant departures from scale-invariance are taken as signs of non-Gaussianity. Unfortunately, the higher-order moments are very sensitive to finite-volume effects and non-Gaussian noise in the data (c.f. the discussion in \\cite{ks97} and references therein). Since many non-Gaussian distribution functions have long tails, it is better to \\emph{directly} compare the rms and the tail of the PDF, rather than relying on the moments, which are integrals over the PDF. Given a distribution function, the level of tail fluctuations maps directly to a level of rms fluctuations; the tail-inferred rms, then, must be consistent with the directly measured rms. Furthermore, one should look for measures which are, as closely as possible, indicators of the initial, linear density field. In this paper, we use the observed temperature distribution of clusters of galaxies as a measure of the tail of the distribution of initial density fluctuations, and galaxy peculiar velocities and redshift-space distortions as measures of the rms. Clusters of galaxies are the most massive virialized objects in the universe, with densities a hundred or more times that of the mean density of the universe, containing only a few percent of all galaxies. They are thus extreme excursions from the mean, and hence are measures of the high potential tail of the distribution of fluctuations. We compare this with the rms level of peculiar velocity fluctuations. The rms is not as sensitive to non-linear growth as are higher order statistics (Dekel, private communication). Both cluster and peculiar velocity observations are at roughly the same spatial scale, $\\sim 10\\hmpc$. Furthermore, they both measure gravitational potential fluctuations, and therefore have a similar dependence on $\\Omega_0$, which can be approximated by defining \\begin{equation} \\e{R}\\equiv\\sigma_R\\Osix, \\end{equation} where $\\sigma_R$ is the rms mass fluctuation on \\emph{tophat} filter scale $R.$ The remaining $\\Omega_0$ dependence is $\\leq 5$\\% for the peculiar velocity and redshift distortion measurements, and $\\leq 25$\\% in the range $0.3 \\leq\\Omega_0\\leq 1.0$ for the cluster measurements. This is the quantity we will attempt to determine throughout our analysis, comparing estimates of it found by various methods. In \\S 2, we examine various recent derivations of $\\e{10}$ from the Mark III peculiar velocity catalog (Willick \\etal\\ 1995, 1996, 1997a), and from \\iras\\ galaxy redshift space distortions. In \\S 3, we infer $\\e{10}$ from the redshift-zero X-ray cluster abundances, assuming a Gaussian and texture PDF, using a modified Press-Schechter approach. We calculate the $\\e{10}$ separately for $\\Omega_0=1.0$, cosmologically flat ($\\Lambda$-dominated) $\\Omega_0=0.3$, and open $\\Omega_0=0.3$. In \\S 4, we discuss the consistency of the rms and the tail-inferred measures of $\\eta$ by performing a likelihood analysis of the Gaussian and texture hypotheses. We also discuss further constraints from observations of clusters at moderate redshift. We summarize in \\S 5. ", + "conclusions": "We show that comparing the amplitude of tail fluctuations indicated by the abundances of clusters of galaxies, to the amplitude of the rms mass fluctuation as indicated by galaxy peculiar velocities and redshift space distortions, can be a strong test of the Gaussianity of the PDF of the initial density field, with a weak dependence on the slope of the density fluctuation spectrum. Observations of peculiar velocities and redshift distortions in the linear regime can measure the rms level of gravitational potential fluctuations present in the universe in the combination $\\e{R}\\equiv\\sigma_R\\Omega_0^{0.6}$, independent of galaxy bias. Various existing measurements of this quantity are not consistent with one another within their stated errors. We assume the existence of a systematic error in each measurement, and use a maximum likelihood technique to combine these recent measurements of $\\e{R}$. The systematic error leads to a $\\sim 20\\%$ uncertainty in the value of $\\e{R}$. This procedure gives $\\e{10,{\\rm rms}}=0.39\\pm 0.07 - 0.51 \\pm 0.10$, depending on the assumed slope of the mass fluctuation spectrum. We use this value, and one recent measurement by Willick \\etal\\ (1997b), which indicates $\\e{10}= 0.27-0.30\\pm 0.04$, to compare with the cluster-inferred measures of $\\e{10}$. We chose the Willick \\etal\\ (1997b) measurement because it is the most recent analysis which explicitly filters close to the $10\\hmpc$ scale on which clusters form. The value of $\\eta_{10}$ inferred from cluster abundances is calculated using the Press-Schechter approach for both a Gaussian and a texture PDF. We find a simple relation for the typical redshift of formation of clusters in both models. This is particularly important in the texture scenario, where the shallower tail of the PDF leads to significant early formation of clusters. This typical redshift is important in the relation between the virial temperature of the clusters and the comoving radius of the initial perturbation. For a texture PDF, the observed abundances of X-ray clusters implies $\\e{10}=0.20\\pm 0.03 - 0.25\\pm 0.04$ for $\\Omega_0=0.3$, and $\\e{10}=0.22-0.29 \\pm 0.03$ for $\\Omega_0=1.0$. Given a Gaussian PDF, on the other hand, the cluster abundances indicate $\\e{10}= 0.29\\pm 0.04 - 0.33 \\pm 0.05$ for $\\Omega_0=0.3$ and $\\e{10}=0.34\\pm 0.05 - 0.42\\pm 0.04$ for $\\Omega_0=1$. We calculate the relative likelihoods of each PDF for a range of cosmological parameters. Using the maximum likelihood average of the rms measurements with the cluster abundances implies the Gaussian model is favored for any $\\Omega_0$. However, using the Willick \\etal\\ (1997b) rms measurement with the clusters implies that a Gaussian is favored for open $\\Omega_0=0.3$, a Gaussian and a texture PDF of roughly equal likelihood for flat $\\Omega_0=0.3$, and a texture PDF is favored for $\\Omega_0=1$. In each case, the accepted hypothesis is accepted at a $\\sim 1.5\\sigma$ level, with likelihood ratio $\\gsim 3$. Preliminary constraints from cluster evolution indicate that low $\\Omega_0$ models are probably more viable than $\\Omega_0$=1 models. Our conclusions can are summarized in Table~\\ref{tab:summary_table}, where a \\chk~(X) indicates the PDF is favored (disfavored) over the alternative by $\\gsim 3\\times$. A $?$ next to the \\chk or X indicates the factor is $\\lsim 3$. Our results imply that the textures+CDM scenario faces some challenges from observations on $5-25h^{-1}$Mpc scales, \\emph{completely independent} of CMB measurements. Although the current state of the measurements of the rms prevents us from firmly distinguishing the Gaussian and texture models, the principle behind this method is clear. The observed cluster abundance at $z = 0$ predicts significantly different values for $\\e{10}$ for the two models. More accurate direct determinations of the measured rms from future redshift surveys and peculiar velocity measurements, or from an improved understanding of the systematic effects which are present in the current analyses, will surely enable us to \\emph{definitively} test \\emph{any} distribution function of initial density perturbations. Furthermore, upcoming high redshift X-ray cluster data will likely break any remaining degeneracy due to the assumed background cosmology. Soon, perhaps, we will finally be able to distinguish different models for the initial PDF of the universe at 10$\\hmpc$." + }, + "9708/astro-ph9708193_arXiv.txt": { + "abstract": "Dramatic torque reversals between spin up and spin down have been observed in half of the persistent X-ray pulsars monitored by the BATSE all-sky monitor on CGRO. Theoretical models developed to explain early pulsar timing data can explain spin down torques via a disk-magnetosphere interaction if the star nearly corotates with the inner accretion disk. To produce the observed BATSE torque reversals, however, these equilibrium models require the disk to alternate between two mass accretion rates, with $\\dot M_{\\pm}$ producing accretion torques of similar magnitude, but always of opposite sign. Moreover, in at least one pulsar (GX 1+4) undergoing secular spin down the neutron star spins down faster during brief ($\\sim 20$ day) hard X-ray flares -- this is opposite the correlation expected from standard theory, assuming BATSE pulsed flux increases with mass accretion rate. The $10$ day to 10 yr intervals between torque reversals in these systems are much longer than any characteristic magnetic or viscous time scale near the inner disk boundary and are more suggestive of a global disk phenomenon. We discuss possible explanations of the observed torque behavior. Despite the preferred sense of rotation defined by the binary orbit, the BATSE observations are surprisingly consistent with an earlier suggestion by Makishima \\etal (1988) for GX~1+4: the disks in these systems somehow alternate between episodes of prograde and retrograde rotation. We are unaware of any mechanism that could produce a stable retrograde disk in a binary undergoing Roche-lobe overflow, but such flip-flop behavior does occur in numerical simulations of wind-fed systems. One possibility is that the disks in some of these binaries are fed by an X-ray excited wind. ", + "introduction": "The spin evolution of an accreting magnetic star, an X-ray pulsar, magnetic CV, or T Tauri star, is thought to be regulated by torques acting between the accretion disk and the stellar magnetosphere (Rappaport \\& Joss 1977; Warner 1990; Konigl 1991). Because of the small neutron star moment of inertia, however, only the X-ray pulsars undergo accretion-induced changes in rotation frequency large enough to measured on short time scales ($\\sim$ days). They are thus ideal laboratories for studying the dynamical interaction between a magnetic star and its accretion disk. Accreting X-ray pulsars are rotating, highly magnetized ($B \\sim 10^{12}$\\,G) neutron stars that accrete material from a stellar companion, either from a stellar wind, or by Roche-lobe overflow mediated by an accretion disk. Disks may also form in wind-fed systems if the captured material has sufficient angular momentum to circularize before reaching the neutron star magnetosphere (see King 1995). The strong magnetic field disrupts the disk and forces the accreting plasma to corotate with the star at a radius where magnetic and fluid stresses roughly balance, $r_m \\sim \\mu^{4/7} \\dot M^{-2/7} (GM_x)^{-1/7} \\sim 10^{8-9}$ cm, where $\\dot M$ is the mass accretion rate, $\\mu$ is the magnetic dipole moment, and $M_x$ is the mass of the neutron star. Although the coupling between the disk and magnetosphere is complicated, and may depend on the geometry and relative orientation of the magnetic field (Wang 1997), in the simplest picture of accretion torque (Pringle \\& Rees 1972; Rappaport \\& Joss 1977) one assumes that the specific angular momentum of material captured from the inner accretion disk is somehow transported onto the star with the accreting matter. For a Keplerian disk, the pulsar will experience a spin-up torque \\begin{equation} N = \\dot M\\sqrt{GM r_m}=2\\pi I \\dot \\nu, \\label{eq:torque} \\end{equation} where $\\dot \\nu$ is the pulsar spin frequency, and $I \\sim 10^{45}{\\,\\rm gm\\,cm^2}$ is the neutron-star moment of inertia. Early observations indicated that, on average, most accreting pulsars were spinning up on a time scale $t_{su}= {\\nu/\\dot \\nu} \\sim 10^{4}$ yrs consistent with equation (\\ref{eq:torque}). This was strong evidence that X-ray pulsars must be compact stars with large magnetospheric radii (Rappaport \\& Joss 1977). This simple spin-up picture had to be modified when two well-studied pulsars, Her X-1 and Cen X-3, were found to be spinning up much more slowly than predicted by equation (\\ref{eq:torque}). Furthermore, these pulsars sometimes underwent short episodes of spin-down (Elsner \\& Lamb 1977; Ghosh \\& Lamb 1979). How could a star capturing material from a disk with the same sense of rotation actually lose angular momentum while continuing to accrete? To explain this behavior Ghosh and Lamb (1979; hereafter GL) argued that the spin-up accretion torque must decrease -- and eventually become negative -- when the stellar rotation frequency approaches the Keplerian orbital frequency of the inner accretion disk, $\\Omega_* \\simeq \\Omega_K(r_m) =(GM/r_m^3)^{1/2}$. Since most X-ray pulsars are in binaries much older than the pulsar spin-up time scale, GL argued they should have reached this near-equilibrium state. In this situation, magnetic field lines which thread the disk beyond the corotation radius (where the disk rotates more slowly than the star) are swept back in a trailing spiral and transport angular momentum outward. Stars close to equilibrium will spin up much more slowly than predicted by equation (\\ref{eq:torque}) and can even spin down while continuing to accrete. \\placefigure{fig:fig1} GL wrote their torque as a modified form of equation (\\ref{eq:torque}), \\begin{equation} N_{GL} = n(\\omega)\\dot M\\sqrt{GM r_m} \\label{eq:gltorque} \\end{equation} where $n(\\omega)$ is a dimensionless function of the ``fastness parameter'', $\\omega = {\\Omega_*/\\Omega_K(r_m)} \\propto \\dot M^{-3/7}\\Omega_* B_*^{6/7}$. For most observations $\\Omega_*$ can be taken as constant, so that, in their theory, observed torque fluctuations mainly reflect the dependence of $N_{GL}$ on $\\dot M$ Although several functional form has been suggested (e.g. Campbell 1987; Wang 1987, 1995, 1997), for our discussion it is only important that $N_{GL}$ is a smooth and monotonically increasing function of $\\dot M$ that crosses zero at some $\\dot M_{crit}$ corresponding to a critical fastness parameter $\\omega_{crit} \\la 1$. An approximate version of $N_{GL}(\\dot M)$ with $\\omega_{crit}=0.8$ is shown in Figure 1. In particular, the spin-up torque becomes negative at low accretion rates; the magnetospheric radius $r_m \\propto \\dot M^{-2/7}$ moves outwards, close enough to the corotation radius $r_{co}=(GM/\\Omega_*^2)^{1/3}$ that the negative magnetic torques become large. Note, however, that sudden changes in accretion torque require sudden changes in $\\dot M$. ", + "conclusions": "We wish to make two points in this paper. First, the highly-sampled BATSE pulsar timing data is difficult to reconcile with standard explanations of spin-down accretion torques in near-equilibrium rotators, $\\Omega_*=\\Omega_K(r_m)$. According to these theories, the switching between spin up and spin down now seen in the majority of persistent BATSE pulsars would require repeated transitions between two mass accretion rates, with $\\dot M_\\pm$ producing torques of comparable magnitude, but always of opposite sign. Moreover, in one pulsar, the observed anticorrelation between torque and 20-60 keV pulsed luminosity is opposite the predicted effect. Some bistable torque mechanism with a switching time scale much longer than any natural time at the inner disk boundary must be at work. This time scale is more consistent with a global disk phenomenon. Recently, Yi, Wheeler \\& Vishniac (1997) have suggested that the observed torque reversals may be due to a transition from a standard Keplerian disk rotation law to a sub-Keplerian advection-dominated flow (ADF); they write the new rotation law $\\Omega^\\prime(r)=A\\Omega_K(r)$ with $A=0.2$ and assume the star is initially spinning near equilibrium. The sudden transition to ADF decreases the corotation radius, $r^\\prime_{co} = A^{2/3}r_{co}$, while the fastness parameter increases $\\omega^\\prime = \\omega/A \\propto \\dot M^{-3/7}B_*^{6/7}/A$. Assuming a GL-type torque (eq. \\ref{eq:gltorque}) in both states, it is not surprising that Yi \\etal (1997) find acceptable fits to the observed transitions: for a given A, one can always adjust $B_*$ and $\\dot M$ to yield the observed torques before and after the transition. Their fitted parameters, however, do not agree with observational constraints. $\\dot M$ must be near the critical accretion rates required by ADF, $\\dot M_{crit} \\sim 0.1 \\alpha^2 \\dot M_{Edd} \\sim 10^{15}-10^{16}\\rm{gm\\,s^{-1}}$ (Narayan \\& Yi 1995), yet Cen X-3, GX 1+4 and OAO 1657-415 are accreting at much higher rates $\\dot M \\sim 10^{17}-10^{18} \\rm{\\, gm\\,s^{-1}}$ (Nagase 1989; Chakrabarty \\etal 1993, 1997b). Moreover, the fitted magnetic field strengths are at least an order of magnitude smaller than required for these stars to be near spin equilibrium. Nevertheless, this scenario is attractive because it is consistent with the transition time scales discussed above -- the ADF transition itself will occur on a fast thermal time $t_{th} \\sim (\\alpha \\Omega_K)^{-1} \\sim 10^3$ s, while the interval between torque transitions is set by slow changes in the mass transfer rate occurring on a global disk viscous time scale. Our second point is admittedly more speculative. Based on the suggestion by Makishima \\etal (1988) that a reversed disk formed in the binary system containing GX~1+4, the BATSE observations themselves are consistent with accretion from disks having alternating senses of rotation. This interpretation naturally explains the bimodal nature of the torques, the comparable torque magnitudes, the torque-luminosity anticorrelation seen in GX 1+4, and the time scales between torque transitions -- without requiring these systems to be near spin equilibrium. Although we are unaware of any physical mechanism that could produce a reversed disk in a system undergoing Roche-lobe overflow, we suggest that these binaries may be accreting from from a stellar wind, possibly self-excited by the X-rays coming from the accreting neutron star. It may be possible to test the reversed disk hypothesis with further observations. As in GX 1+4, the torque in any system undergoing secular spin-down should spin down faster with increased luminosity. Observations with the X-ray Timing explorer may be able to determine the correlation between torque and luminosity in Cen X-3 while it is in a spin-down state. We estimate that a two day pointing of Cen X-3 with XTE could make of order 5 torque measurements sensitive to 10\\% variations at the $3\\sigma$ level. If the accretion disk is fed from an X-ray excited wind with a flip-flop type instability, changes in column density inferred from low-energy absorption may also reveal transitions in the structure of the accretion flow (Day \\& Stevens 1993). This work was funded in part by NASA grants NAG 5-3119, NAG 5-3109, NAG 5-1458, NAG 5-3293, NAGW-4517 and NGT-51184 and the Alfred P. Sloan Foundation." + }, + "9708/astro-ph9708176_arXiv.txt": { + "abstract": "We use the rotation curves of a sample of dark matter dominated dwarf and low-surface brightness (LSB) late-type galaxies to study their radial mass distributions. We find that the shape of the rotation curves is remarkably similar for all (both dwarf and LSB) galaxies in the sample, suggesting a self-similar density distribution of their dark matter (DM) halos. This shape can be reproduced well by a density profile with a shallow central cusp $[\\rho(r)\\propto 1/r^{\\gamma},{\\ }\\gamma\\approx 0.2-0.4]$ corresponding to a steeply rising velocity curve $[v(r)\\propto r^g,{\\ } g\\approx 0.9-0.8]$. We further show that the observed shapes of the rotation curves are well matched by the average density profiles of dark matter halos formed in very high resolution simulations of the standard cold dark matter model (CDM), the low-density CDM model with cosmological constant ($\\Lambda$CDM), and the cold$+$hot dark matter model with two types of neutrino (CHDM). This is surprising in light of several previous studies, which suggested that the structure of simulated dark matter halos is inconsistent with the dynamics of dwarf galaxies. We discuss possible explanations for this discrepancy and show that it is most likely due to the systematic differences at small radii between the analytic model proposed by Navarro, Frenk, \\& White, with $\\gamma_{\\rm NFW} = 1$, and the actual central density profiles of the dark matter halos. We also show that the mass distributions in the hierarchically formed halos are {\\em on average} consistent with the shape of rotation curves of dark matter dominated galaxies. However, the scatter of the individual profiles around the average is substantial. Finally, we show that the dark matter halos in our hierarchical simulations and the real galaxies in our sample exhibit a similar decrease in their characteristic densities with increasing characteristic radial scales and show increase in their maximum rotation velocities with increase in the radii at which their maximum velocities occur. ", + "introduction": "The amount of luminous matter (stars and gas) in many spiral and irregular galaxies is not sufficient to explain the amplitude and shape of their rotation curves (RCs). This discrepancy is usually interpreted as evidence for the presence of an extended dark matter (DM) halo surrounding the visible regions of galaxies (e.g., Casertano \\& van Gorkom 1991; Persic, Salucci, \\& Stel 1996, and references therein). The extent of the dark matter halos, estimated using satellite dynamics, is\\footnote{We assume that the present-day value of the Hubble constant is $H_0=100h$ km/s/Mpc.}$\\sim 0.2-0.5h^{-1}$ Mpc (Zaritsky \\& White 1994; Carignan et al. 1997; Zaritsky et al. 1997). However, the dynamical contribution of the dark matter can be substantial even in the very inner regions of galaxies: the observed rotation velocities of some dwarf and low-surface brightness (LSB) galaxies imply that DM constitutes a dominant fraction (up to $\\sim 95\\%$) of dynamical mass within the last measured point of their RCs (e.g., Carignan \\& Freeman 1988; Martimbeau, Carignan, \\& Roy 1994; de Blok \\& McGaugh 1997). These {\\em dark matter dominated galaxies} offer a unique opportunity for probing {\\em directly} the density structure of DM halos which can be then compared with predictions of theoretical models. The detailed structure of DM halos formed via dissipationless hierarchical collapse in CDM-like models was recently studied using high-resolution $N$-body simulations (Dubinski \\& Carlberg 1991; Navarro, Frenk, \\& White 1996, 1997, hereafter NFW96 and NFW97). The halo density profiles were found to be cuspy (coreless) and well fitted by the following two-parameter profile (NFW96): \\begin{equation} \\rho_{\\rm{NFW}}(r) = \\frac{\\rho_s}{(r/r_s)\\left(1+r/r_s\\right)^2}. \\end{equation} The characteristic density, $\\rho_s$, and radius, $r_s$, are sensitive to the epoch of halo formation and are tightly correlated with the halo virial mass (NFW 96,97). Therefore, the results of these simulations suggest a coreless and self-similar density structure of DM halos, with the virial mass being the single scaling parameter. The structure of the inner regions of galactic halos was studied by Flores \\& Primack (1994) and Moore (1994), who used high-resolution rotation curve measurements of several dark matter dominated dwarf galaxies. The central density distributions in these galaxies were found to be inconsistent with the singular [$\\rho(r)\\propto 1/r$] behavior predicted by equation (1). The scaling properties of the observed halos were analyzed by Burkert (1995, hereafter B95), who pointed out that shapes of the density profiles of four dwarf galaxies analyzed by Moore (1994) are remarkably similar and are well fitted by the following phenomenological density profile: \\begin{equation} \\rho_{\\rm{B}}(r) = \\frac{\\rho_b}{(1+r/r_b)\\left[1+(r/r_b)^2\\right]}. \\end{equation} Parameters $\\rho_b$ and $r_b$ were found to be strongly correlated, in qualitative agreement with the predictions of hierarchical models (B95). In this paper we study the observed density structure in a sample of dark matter dominated galaxies inferred from their rotation curves. Particularly, we test two predictions of previous simulations of hierarchical halo formation: (1) cuspy central density distribution and (2) self-similarity of the halo density structure. We then use results of high-resolution $N$-body simulations to compare the observed rotation curves with circular velocity profiles of dark matter halos formed in different structure formation models. \\begin{deluxetable}{lcccccccccc} \\tablewidth{0pt} \\tablecaption{The sample of dwarf and LSB galaxies} \\tablehead{ \\colhead{ }&\\colhead{ }&\\colhead{$r_0$}&\\colhead{$V_0$}&\\colhead{$\\rho_0$}&\\colhead{$r_t$}&\\colhead{$V_t$}&\\colhead{$r_{max}$}&\\colhead{$V_{max}$}& Distance& \\nl Galaxy & $M_B$&$h^{-1} {\\rm kpc}$&km s$^{-1}$&$10^8h^3 M_{\\odot} {\\rm kpc^{-3}}$&$h^{-1} {\\rm kpc}$&km s$^{-1}$&$h^{-1} {\\rm kpc}$&km s$^{-1}$&$h^{-1} {\\rm Mpc}$&Reference\\nl (1) &(2) &(3) &(4) &(5) &(6)&(7)&(8)&(9)&(10)&(11) \\nl } \\startdata & & & & & Dwarf & & & & & \\nl DDO 154 &-13.8&2.0&38&0.33&2.9& 76& 5.6& 47&3.0& 1,2 \\nl DDO 170 &-15.2&4.4&52&0.13&6.5&105&12.4& 64&14& 3 \\nl NGC 2915 &-16.8&1.1&69&3.77&1.7&140& 3.2& 86&2.5& 4 \\nl IC 2574 &-15.7&5.1&65&0.15&8.2&138&15.6& 85&2.0& 5 \\nl NGC 5585 &-17.5&2.3&73&0.93&3.4&148& 6.5& 91&4.7& 6 \\nl DDO 236 &-16.8&4.4&59&0.16&6.8&122&13.1& 75&1.4& 7 \\nl DDO 7 &-17.7&3.8&87&0.49&5.5&176&10.5&108&18& 8 \\nl DDO 10 &-16.3&3.5&54&0.23&5.3&112&10.0& 69&7.8& 8 \\nl UGC 2684 &-13.7&1.6&41&0.59&2.6& 86& 4.9& 53&4.1& 8 \\nl DDO 34 &-15.7&1.5&54&1.29&2.1&109& 4.0& 67&5.9& 8 \\nl & & & & & LSB & & & & & \\nl F568-1 &-17.5&3.8&97&0.61& 5.5&197&10.5&121&64& 9 \\nl F568-3 &-17.7&6.1&96&0.23& 9.1&196&17.4&121&58& 9 \\nl F568-V1 &-17.3&4.9&96&0.36& 7.2&194&13.7&119&60& 9 \\nl F571-8 &-17.0&6.8&121&0.30&10.1&248&19.4&153&36& 9 \\nl F574-1 &-17.8&8.8&95&0.11&13.6&197&25.9&121&72 & 9 \\nl F583-1 &-15.9&4.6&71&0.22& 6.7&143&12.8& 88&24& 9 \\nl F583-4 &-16.3&6.0&64&0.11& 9.3&133&17.8& 82&37& 9 \\nl \\tablenotetext{ } {NOTES.-- Col.(2) $M_B$, blue absolute magnitude; col.(3) best fit $r_0$(see eq. [3]; the fitting procedure is described in \\S2.2); col.(4) $V_0=V(r_0)$; col.(5) best fit $\\rho_0$ (see eq.[3]); col.(6) best fit $r_t$ (see eq.[4]); col.(7) best fit $V_t$ (see eq.[4]); col.(8) $r_{max}$; col.(9) $V_{max}=V(r_{max})$; col.(10) distance to galaxy adopted in this study; } \\tablenotetext{ } { REFERENCES.-- (1) Carignan \\& Freeman 1988; (2) Carignan \\& Beaulieu 1989; (3) Lake et al. 1990; (4) Meurer et al. 1996; (5) Martimbeau et al. 1994; (6) C\\^ot\\'e et al. 1991; (7) Jobin M. \\& Carignan C. 1990; (8) van Zee et al. 1997; (9) de Blok et al. 1996. } \\enddata \\end{deluxetable} ", + "conclusions": "To summarize, our results are the following. \\begin{itemize} \\item Rotation curves of DM dominated dwarf and LSB galaxies have a similar shape. This shape is inconsistent with $\\rho_{NFW}(r)$ (eq. [1]), but is well fit by the coreless profile described by eq. (3) which has a shallower slope at small scales: $\\rho(r)\\propto 1/r^{\\gamma},{\\ }\\gamma\\approx 0.2-0.4$, corresponding to a steeply rising velocity curve $[v(r)\\propto r^g,{\\ } g\\approx 0.9-0.8]$. \\item We find that {\\em on average} the velocity profiles of the halos formed in the hierarchical structure formation models analyzed in this paper (CDM, $\\Lambda$CDM, and CHDM) and observed dark matter halos are in reasonably good agreement. We find a substantial amount of scatter in the central density profiles of individual halos around the average. The physical processes which lead to differences in the central density profiles seen in our halos (e.g., environment, dynamical state, etc.) are not clear and will be the subject of a future study. \\item The inner ($r<30 h^{-1}$ kpc) average density profiles of DM halos in all of our simulations are well fit by model (3) with ($\\alpha$, $\\beta$, $\\gamma$)$=$(2,3,0.2), and equivalently the rotation curves are well described by RC (4) with $(a,b,g)=(1.5,0.34,0.9)$. The profiles systematically deviate from the NFW profile (1) at small scales. \\item We find that dark matter dominated dwarf and LSB galaxies show correlations between their characteristic density and radius consistent with the correlations of hierarchically formed DM halos: physically smaller halos are denser. We find a similar correlation between the maximum of the rotation curve, $v_{max}$, and the corresponding radius $r_{max}$. \\end{itemize}" + }, + "9708/astro-ph9708030_arXiv.txt": { + "abstract": "We present evidence for changes in the strength and profile of the iron K$\\alpha$ line in Active Galactic Nuclei (AGN), based on X-ray observations with \\asca. There is a clear decrease in the strength of the line with increasing luminosity. This relation is is not due solely to radio power, as it persists when only radio-quiet AGN are considered and therefore cannot be fully explained by relativistic beaming. In addition to the change in strength, the line profile also appears to be different in higher luminosity sources. We discuss these results in terms of a model where the accretion disk becomes ionized as a function of the accretion rate. ", + "introduction": "\\label{sec:intro} Seyfert 1 galaxies exhibit iron K$\\alpha$ emission lines in their X-ray spectra which are characteristic of relativistic effects in an accretion disk surrounding a central black hole (Tanaka \\etal\\ 1995; Yaqoob \\etal\\ 1995; Nandra \\etal\\ 1997 hereafter N97). These lines can be used as a diagnostic of the innermost regions of AGN, and therefore merit further study in classes other than Seyfert 1s. The iron K$\\alpha$ emission was first studied in detail using the \\gi\\ spectra of Seyfert galaxies (Nandra \\& Pounds 1994 and references therein) and based on these results, Iwasawa \\& Taniguchi (1993, hereafter IT93) suggested that there may be an X-ray ``Baldwin Effect'' whereby the equivalent width (EW) of the emission line reduced with increasing luminosity. However, this result has been disputed (Nandra \\& Pounds 1994) and it was unclear whether the correlation held when emission lines from ``quasars'' were considered (Williams \\etal\\ 1992; IT93). With far greater sensitivity and spectral resolution than \\gi, \\asca\\ (Tanaka, Inoue \\& Holt 1994) can be used to provide a more stringent test of such an hypothesis. \\asca\\ results for individual sources have been suggestive that this trend might hold (e.g., Elvis \\etal\\ 1994; Nandra \\etal\\ 1995). However, a systematic comparison requires consideration of a larger number of sources. The results of an analysis of the dependence of iron line properties with luminosity, based on a sample of 39 AGN with broad optical lines, will be the subject of this {\\it Letter}. ", + "conclusions": "\\label{sec:discuss} Using a (poorly-selected) sample of X-ray observations of broad-line AGN, we have shown clear evidence of an X-ray ``Baldwin'' effect, i.e. a reduction in the strength of the iron K$\\alpha$ line with increasing luminosity. Such an effect was originally suggested based on \\gi\\ data by IT93. Although the effect could be partially due to beaming of the X-rays away from the putative accretion disk in radio-loud sources, there is still strong evidence for differences between high and low luminosity sources when only radio-quiet objects are considered. We therefore conclude that the primary effect is most likely with source luminosity and discuss the impact of our results in that context. The emission line in low-luminosity AGN is thought to arise from an accretion disk, where Doppler and gravitational effects produce the extreme broadening, and especially the red wing. An alternative origin for a line core at 6.4~keV is in the putative molecular torus which may obscure the line-of-sight to Seyfert 2 galaxies (Ghisselini, Haardt \\& Matt 1994; Krolik, Madau \\& Zycki 1994). In principle then, differences in the line profiles could arise from differing contributions from the accretion disk and torus. We observe an effect consistent with this, in that the 6.4~keV peak reduces with increasing luminosity and largely disappearing above $L_{\\rm X} > 10^{45}$~erg s$^{-1}$ (Fig.~\\ref{fig:lum_ew}). The upper limit to the EW of any narrow, 6.4 keV line in the highest luminosity bin is $\\sim 25$~eV, so our data are consistent with a small contribution from the torus in all sources, and if that contribution decreased when the luminosity increased it would account for some of the differences in line profiles. However, we also observe an effect that the red wing reduces with luminosity, implying that the disk-line component also changes. Indeed, the entire effect can be attributed to changes in the disk-line. Nandra \\etal\\ (1995) suggested that the lack of significant iron line emission in high luminosity AGN might be due to the fact that those sources have a high accretion rate, causing the disk to become ionized (Matt, Fabian \\& Ross 1993), with iron being fully stripped. Some support from that hypothesis came with the detection of an emission line consistent with highly-ionized iron in an ``intermediate'' luminosity QSO, PG 1116+215 (Nandra \\etal\\ 1996). Our results are interpretable in this context. For a given black hole mass, higher luminosity sources should have a higher accretion rate, as well as more intense X-ray (ionizing) luminosity. Both would tend to strip atoms in the disk. At some point, iron will begin to be ionized, which should cause more ``blue'' flux to be observed from high-ionization species. In these intermediate ionization states, resonance scattering can also cause a reduction in the line flux (Matt, Fabian \\& Ross 1993, 1996). An increase in the effective fluorescence yield in the He-like and H-like states would cause stronger line emission when those species are dominant, but if the emission comes from a range of radii (and therefore ionization state) in the disk, which appears to be the case (N97), that effect may not be clearly observable. At another transition point, iron atoms in the inner disk will begin to become fully stripped, which would cause a reduction in the ``red wing'' and a shift of the mean energy above 6.4~keV. When iron becomes fully stripped throughout the X-ray illuminated part of the disk, no emission line will be observed from the disk at all. All of these effects are observed in Fig.~\\ref{fig:profs_lumin}. If this model is correct, we should observe associated changes in the Compton reflection component (e.g., Zycki \\& Czerny 1994). As the ionization rises, the disk becomes more reflective in the soft X-ray band, causing a ``soft excess'', with the potential for associated line emission from elements lighter than iron (e.g. O, Ne). Again, for very high ionization states, we see the Compton reflection without it suffering absorption in the disk, making the ``contrast'' with the continuum very low, resulting in an apparently-weak Compton hump. Such effects are only easily testable with instruments with better high-energy efficiency than \\asca. Of course there may be significant differences in the black hole mass when moving from low to high luminosity sources. Estimates of the black hole masses of local AGN have been made based on various arguments such as optical/UV line widths, stellar kinematics and maser observations (e.g., Koratkar \\& Gaskell 1991; Ford \\etal\\ 1994; Miyoshi \\etal\\ 1995). Interestingly, these masses tend to lie in the region $10^{7-8}$~\\Msun, regardless of the technique employed. Larger masses would be expected for the highest-luminosity objects in our sample, since if their emission is isotropic, their luminosities would exceed the Eddington limit unless $M>10^{9}$~\\Msun. However, it is interesting to note that for a $10^{8}$~\\Msun\\ hole, the transition to $\\sim 10$~per cent Eddington accretion occurs at $10^{45}$~erg s$^{-1}$. This is approximately where Matt \\etal\\ (1993) predict that the disk will start to become significantly ionized. and where we begin to see a change in the line profile. Super-Eddington accretion occurs at $10^{46}$~erg $s^{-1}$, above which luminosity the emission line disappears. Thus we speculate that the AGN in our sample cover a relatively small mass range, and that the differences in the source luminosities are due to differences in accretion rate, which then affect the line profiles. The implication of this is that the quasar phenomenon is short-lived. Further exploration of these models awaits the formation of larger, and preferably complete, well-selected samples within the \\asca\\ archive, which we anticipate within the next few years." + }, + "9708/astro-ph9708206_arXiv.txt": { + "abstract": "We present a technique which employs artificial neural networks to produce physical parameters for stellar spectra. A neural network is trained on a set of synthetic optical stellar spectra to give physical parameters (e.g.\\ \\teff, \\logg, \\met). The network is then used to produce physical parameters for real, observed spectra. Our neural networks are trained on a set of 155 synthetic spectra, generated using the \\sp\\ program written by Gray (Gray \\& Corbally 1994, Gray \\& Arlt 1996)\\nocite{gray_94b}\\nocite{gray_96a}. Once trained, the neural network is used to yield \\teff\\ for over 5000 B--K spectra extracted from a set of photographic objective prism plates (Bailer-Jones, Irwin \\& von Hippel 1997a)\\nocite{bailerjones_97a}. Using the MK classifications for these spectra assigned by Houk (1975, 1978, 1982, 1988)\\nocite{houk_75a}\\nocite{houk_78a}\\nocite{houk_82a}\\nocite{houk_88a}, we have produced a temperature calibration of the MK system based on this set of 5000 spectra. It is demonstrated through the metallicity dependence of the derived temperature calibration that the neural networks are sensitive to the metallicity signature in the real spectra. With further work it is likely that neural networks will be able to yield reliable metallicity measurements for stellar spectra. ", + "introduction": "The MK classification system was first proposed in its current form in 1943 by Morgan, Keenan \\& Kellman (1943)\\nocite{morgan_43a}, and has since undergone a number of revisions (e.g.\\ Keenan \\& McNeil (1976)\\nocite{keenan_76a}, Morgan, Abt \\& Tapscott (1978)\\nocite{Morgan_78a}). MK classification is the only widely used system for stellar spectral classification. Over its history it has contributed towards a number of important developments in astronomy, such as the further development of the now-famous HR diagram (Hertzprung 1911, Russell 1914) and the identification of anomalous stars. Currently, MK classification is largely used as a tool in the preliminary analysis of unusual stars, and in selecting stellar samples for further study. An often-stated advantage of the MK system is that its classifications, often based upon the visual inspection of spectra, are static because they are based on a set of standards. However, a given spectrum may be classified differently by different people, and any one person may also classify a given spectrum differently at different times. These problems of subjectivity could be partially alleviated through the use of automated classifiers (von Hippel et~al.\\ 1994\\nocite{vonhippel_94a}, Bailer-Jones et~al.\\ 1997a\\nocite{bailerjones_97a}). Automated classifiers could also produce quantitative errors associated with their classifications. Another problem with the MK system is that it lacks a well-defined metallicity parameter, whereas metallicity variations are known to have a significant effect on the appearance of high ($\\sim$ 1\\,\\AA) resolution spectra. This limits the system to classifications of bright, nearby stars which show only limited metallicity variations. Attempts to extend and revise MK classification (e.g.\\ Corbally, Gray \\& Garrison 1994)\\nocite{corbally_94a} may well prove valuable, but as our understanding of stellar spectra grows, particularly from computational work with model atmospheres and synthetic spectra, it becomes increasingly desirable to obtain reliable physical parameterizations of stars. Any classification system is a compromise between retaining the full information in the spectrum and the need for a compact summary of it. The optimal `summary' is of course given by the physical parameters. Advances in computational power and data storage since the inception of MK classification mean that it is now practicable to process and store large numbers of spectra. The development of fast, automated classifiers will mean that it is feasible to `classify' large numbers of stellar spectra in terms of their physical parameters and to re-classify them rapidly whenever stellar models are improved. Physical parameters should be obtained from an original spectrum, rather than an empirical classification, as the latter may well disregard certain spectral features which later turn out to be important. One of the advantages of the MK classification system is that it is an empirical system based on unchanging standards, whereas any direct parameterization of a spectrum depends upon the quality of stellar models and will change as these improve. MK classifications could remain as useful labels giving a rough `feel' for a spectrum, and the work of von Hippel et~al.\\ (1994)\\nocite{vonhippel_94a} and Bailer-Jones et~al.\\ (1997a)\\nocite{bailerjones_97a} has shown that reliable automated MK classification is possible. In this paper we demonstrate how we extend our automated techniques to the determination of physical parameters directly from an observed spectrum. ", + "conclusions": "We have shown that neural networks trained on synthetic spectra provide low-error predictions for the effective temperature, \\teff, of a star based on the star's optical spectrum. By applying these neural networks to spectra with exisiting MK classifications, we have obtained a calibration between the MK spectral type parameter and \\teff. This calibration shows a good agreement with a number of calibrations from the literature. The calibration was obtained statistically from a number of optical spectra at each spectral type with a resolution of $\\approx$ 2\\,\\AA\\,$pix^{-1}$. The precision of this calibration is largely limited by the cosmic scatter in temperature for a given MK class and by limitations of the stellar models. It has also been shown that metallicity effects have to be considered when trying to determine \\teff. Further work is required before neural networks can be used to accurately quantify metallicities. In particular, further processing of the synthetic spectra into the format of the MHD spectra (e.g.\\ by the addition of noise and calibration of the flux scale) may be required. Our work nonetheless demonstrates that our networks are sensitive to metallicity." + }, + "9708/astro-ph9708212_arXiv.txt": { + "abstract": "We investigate X-ray isophote twists created by triaxiality differences between the luminous stellar distributions and the dark halos in elliptical galaxies. For a typically oblate luminous galaxy embedded in a more prolate halo formed by dissipationless collapse, the triaxiality difference of $\\Delta T \\simeq 0.7$ leads to typical isophote twists of $\\langle \\Delta \\psi_{\\rm X} \\rangle \\simeq 16^\\circ \\pm 19^\\circ$ at 3 stellar effective radii. In a model which includes baryonic dissipation the effect is smaller, with $\\Delta T \\simeq 0.3$ and $\\langle \\Delta \\psi_{\\rm X} \\rangle \\simeq 5^\\circ \\pm 8^\\circ$. Thus, accurate measurements of X-ray isophote twists may be able to set constraints on the interactions between baryons and dissipationless dark matter during galaxy formation. The 30$^\\circ$ X-ray isophote twist in the E4 galaxy NGC 720 cannot be reproduced by our model, suggesting an intrinsic misalignment between the halo and the stars rather than a projection effect. ", + "introduction": "Determining the intrinsic shapes of the luminosity and mass distributions of elliptical galaxies is an important unresolved problem. Both observational constraints from individual systems and predictions from theory demonstrate that the stellar and dark matter components of ellipticals are well described as triaxial ellipsoids. Parametric (e.g., Ryden 1992; Lambas, Maddox, \\& Loveday 1992) and non-parametric (Fasano \\& Vio 1991; Tremblay \\& Merritt 1995; Ryden 1996) inversions of the ellipticity distribution of the stellar components rule out the hypotheses that elliptical galaxies are all oblate or all prolate. Because of the information lost in projection, inversions that include triaxiality cannot determine a unique shape distribution, but the ellipticals seem to fall into two classes: bright ($M_B < -20$), triaxial, boxy galaxies and faint, roughly oblate, disky galaxies (Tremblay \\& Merritt 1996). This division is consistent with other structural, dynamical, and evolutionary evidence for two classes of ellipticals (e.g., Fasano 1991; Fasano \\& Vio 1991; Busarello, Longo, \\& Feoli 1992; Capaccioli, Caon, \\& D'Onofrio 1992; Nieto, Poulain, \\& Davoust 1994; J{\\o}rgensen \\& Franx 1994; Kormendy \\& Bender 1996). Binney (1978, 1985; see also Contopoulos 1956) first noted that a triaxial galaxy with its net angular momentum vector aligned with its short axis would show a ``kinematic misalignment'' between the projected axes of the rotation and the light, and surveys of elliptical galaxy kinematics have found that such kinematic misalignments are common (Davies \\& Birkinshaw 1988; Franx, Illingworth, \\& Heckman 1989; Jedrzejewski \\& Schechter 1989). However, a kinematic misalignment may also be caused by an {\\it intrinsic} misalignment between the rotational and the short axes because the angular momentum vector of an equilibrium triaxial system can lie anywhere in the plane of the long and short axes (Heiligman \\& Schwarzschild 1979). Alternatively, the galaxy may not be in equilibrium (due to a recent accretion event or to tidal forces from a satellite galaxy), or it may be a misidentified S0 with a bar (see Merritt 1997a for a review of examples). A statistical analysis of kinematic misalignments in bright ellipticals (Franx, Illingworth, \\& de Zeeuw 1991) found limits on the mean stellar triaxiality and the mean intrinsic misalignment angle of $\\langle T_* \\rangle \\leq$ 0.7 and $\\langle \\psi_{\\rm int} \\rangle \\leq$ 45$^\\circ$, where the upper bounds for both parameters occurred only for solutions with two widely-separated peaks. Although degeneracies prevented strong constraints on $T_*$, there were indications for a bimodal population of galaxies, with a large fraction of nearly-oblate, short-axis rotators, and a small fraction of nearly-prolate, long-axis rotators. Tenjes et al. (1993) and Statler (1994a,b, Statler \\& Fry 1994) have devised methods which invert the projected streamlines of stellar orbits to determine the intrinsic shape of a galaxy, and their initial results also indicate the existence of a bimodal population (Statler, Dutta, \\& Bak 1997; Bak \\& Statler 1997). Theoretical models of halo formation generally produce flattened, prolate-triaxial halos (e.g., Bardeen et al. 1986; Frenk et al. 1988; White \\& Ostriker 1990; Aguilar \\& Merritt 1990; Katz 1991; Cannizzo \\& Hollister 1992; Cole \\& Lacey 1996). For example, $N$-body simulations of a small sample of halos by Dubinski (1991, 1992) and Dubinski \\& Carlberg (1991) found that at 25 kpc they are very flat (mean short axis ratio of $\\langle c_{\\rm d} \\rangle \\simeq 0.42 \\pm 0.06$) and nearly-prolate ($\\langle T_{\\rm d} \\rangle \\simeq 0.8 \\pm 0.2$). The final halo shapes and angular momenta were very sensitive to the initial cosmological tidal field, and the angular momentum vectors were well-aligned with the short axes, with an average intrinsic misalignment of $\\langle \\psi_{\\rm int} \\rangle \\simeq 26^\\circ \\pm 22^\\circ$. Warren et al. (1992) found similar results ($\\langle c_{\\rm d} \\rangle \\simeq 0.62 \\pm 0.13$ and $\\langle T_{\\rm d} \\rangle \\simeq 0.68 \\pm 0.24$ at 40 kpc), and found that the more massive halos were slightly flatter and more prolate. Again, the angular momentum vector was most commonly aligned with the short axis. We expect dissipative processes to produce baryonic galaxies with shapes that are drastically different from those of their parent halos (e.g., Katz \\& Gunn 1991; Udry 1993; Navarro \\& White 1994), but the concomitant effects on the shapes of the halos are unclear. Preliminary simulations suggest that halos become rounder and more oblate through interactions with baryons (Evrard, Summers, \\& Davis 1994; Dubinski 1994), probably due to the destruction and damping of halo box orbits. Dubinski (1994) found that the flatness and prolateness are reduced relative to a dissipationless model ($c_{\\rm d} \\sim$ 0.6 versus 0.4, and $T_{\\rm d} \\sim$ 0.4-0.5 versus 0.8, at 20 kpc). Late, major mergers of disk galaxies in compact groups may have produced some fraction of the present population of ellipticals (e.g., Heyl, Hernquist, \\& Spergel 1994; Weil \\& Hernquist 1996; Barnes \\& Hernquist 1996). Although the simulations have not carefully explored the effects of the mergers on the shapes of the dark halos, Weil \\& Hernquist (1996) suggest that the shapes of remnant halos are rather round and oblate-triaxial, having no correlation with the shapes of the remnant stellar distributions. Current theoretical and observational results thus suggest that the modestly triaxial-oblate luminous parts of elliptical galaxies are embedded in halos which are more prolate, so that there should be a strong gradient with radius in the triaxiality of the mass distribution. Direct measurements of the shapes of halos are far more difficult than measurements of the shapes of the stellar distributions. Careful examination of the kinematics of polar rings (Arnaboldi et al. 1993), gas disks (Lees 1991; Bertola et al. 1991; Franx, van Gorkom, \\& de Zeeuw 1994; Plana \\& Boulesteix 1996; Bureau \\& Freeman 1997), and planetary nebulae (Hui et al. 1995; Mathieu, Dejonghe, \\& Hui 1996) have been used to constrain the shapes of both stellar and dark matter potentials. Unfortunately, such tracers are rarely found at large enough galactocentric radii to be helpful for probing the dark matter potential, and it is likely that there are systematic correlations between the shapes of halos and the presence of rings and disks. Gravitational lenses can also be used to directly measure the shapes of mass distributions, and the data appear to require very flat mass distributions. However, the number of lenses is small, and the quantitative effects of external tidal perturbations are not yet understood (King \\& Browne 1996; Kochanek 1996; Keeton, Kochanek, \\& Seljak 1997; Witt \\& Mao 1997). The most promising candidate for directly measuring the intrinsic shapes of nearby halos is high-resolution mapping of the X-ray emission from hot gas in the halo potential. For example, Buote \\& Canizares (1994, 1996b, 1997; hereafter BC94, BC96b, BC97) examined the radial profiles, position angles, and axis ratios of the X-ray isophotes of the bright, isolated E4 galaxy NGC 720, and found that the shape and the position angle of the potential cannot be produced by the stars. Thus, by using only geometric evidence and the assumption that the gas is in quasi-hydrostatic equilibrium, they showed that the stars must be embedded in a more massive, {\\it flatter}, dark matter halo. For the S0 galaxy NGC 1332, Buote \\& Canizares (1996a) reached a qualitatively similar conclusion. X-ray observations can also be used to constrain the triaxialities of halos by searching for X-ray isophote twists. The projection of a triaxial light distribution whose triaxiality varies with radius produces isophotes whose axis orientations vary with radius (see Mihalas \\& Binney 1981). Isophote twists in the optical surface brightness are common (e.g., Williams \\& Schwarzschild 1979; Leach 1981), and can be used to constrain the shape of an individual galaxy (Fasano 1995), although care must be taken to rule out intrinsic axis twist caused by other factors (Fasano \\& Bonoli 1989; Nieto et al. 1992). X-ray emission from a triaxial system can show an analogous isophote twist --- as pointed out by Binney (1978) for the case of a galaxy cluster --- allowing the triaxiality of an individual galaxy's potential to be probed. In the simplest case, if a galaxy's luminous and dark matter distributions are intrinsically aligned and have {\\it constant} but {\\it different} triaxialities, then the projected axes of its X-ray isophotes will twist from small radii, where the potential is dominated by the stellar core, to large radii, where it is dominated by the dark halo. Similarly, as discussed by BC96b, the radial triaxiality gradients produced in halos by dissipative processes can result in X-ray isophote twists. Here we employ simple models to study the behavior of X-ray twists, and to determine how well they constrain the triaxiality of halos. In \\S 2.1, we use an analytic approximation to find the amplitudes of twists expected for different assumptions about the shapes of galaxies and halos. In \\S 2.2, we numerically calculate and project the X-ray emission for a small sample of models to verify the predictions of the analytic model, and to examine the radial behavior of the twist more closely. In \\S 3, we attempt to model the large twist observed in NGC 720 (BC94, BC96b), and we present our conclusions in \\S 4. ", + "conclusions": "Current observational evidence indicates that the stellar parts of elliptical galaxies are modestly flattened and close to oblate, while current theories for the formation of galaxies suggest that the halos are more flattened and more prolate. The extent of this shape difference depends on the interactions between baryons and dissipationless dark matter during galaxy formation, but all current models imply that the mass distribution of an early type galaxy increases in flatness and prolateness with radius. There is already some evidence from the flatness of X-ray isophotes (Buote \\& Canizares 1994, 1996a) and from gravitational lens models (King \\& Browne 1996; Kochanek 1996) that the typical mass distribution may be flatter than the luminosity distribution. A difference in triaxiality between luminous galaxies and their dark halos is detectable through misalignments of gravitational lenses with their stellar distributions, and is a natural explanation for the second shear axis that is necessary to fit lensing models, although intrinsic misalignments and external tidal shear sources are also viable explanations (Keeton, Kochanek, \\& Seljak 1997). Such a triaxiality difference also produces X-ray isophote twists (Binney 1978; Buote \\& Canizares 1996b). We examined the behavior of X-ray isophote twists using several simple models of a luminous galaxy of constant triaxiality embedded in an intrinsically-aligned dark halo of a different triaxiality. A simple analytic approximation gives accurate estimates of the asymptotic position angles of the X-ray isophotes, but predicting the detailed radial behavior of the twist requires numerical simulations. For a reasonable model of a galaxy and halo, we find that the X-ray isophote position angle makes a gradual transition from the center of the galaxy, where it is aligned with the position angle of the optical isophotes, to the periphery, where it is aligned with the position angle of the projected halo mass. The ``half-way'' point of the twist occurs well inside the stellar effective radius ($\\sim$~0.1-0.6 $R_{\\rm eff}$), and is not detectable given present X-ray resolution limits, but the misalignment of the X-ray isophotes at large radii with the stellar isophotes is more easily observable. By examining the amplitude of the misalignment for a large population of galaxies, we can in principle distinguish between alternative models for the halo shapes. The very prolate halos predicted by dissipationless collapse simulations should produce mean misalignments of $\\langle \\Delta \\psi_{\\rm X} \\rangle \\simeq 16^\\circ \\pm 19^\\circ$ at $\\sim 3 R_{\\rm eff}$, while the more oblate halos predicted by simulations which include baryonic dissipation should produce smaller misalignments ($\\simeq 5^\\circ \\pm 8^\\circ$). In practice, measurements of these twists will be difficult because the twist angles are small, the X-ray isophotes are fairly round ($\\langle q_{\\rm X} \\rangle \\sim 0.9$), and the number of isolated, nearby X-ray galaxies is limited. Our results are not very sensitive to the parameters of the model, but the presence of a discrete source component can modify the detailed properties of the twist and should be included in any statistical test. Note that while we scaled our models to match the halo shapes predicted by simulations, {\\it any} mass distribution with a strong radial triaxiality gradient will produce an X-ray isophote twist when observed from the proper angles. For example, the disruption of box orbits by a central stellar cusp can cause the central parts of an elliptical galaxy to become more oblate than the outer parts due to the longer time scales for scattering and phase mixing of the outer orbits (e.g. Merritt 1997b), leaving a twisting signature in the X-ray isophote position angles. We attempted to model the observed X-ray twist of the bright, isolated E4 galaxy NGC 720 (BC94, BC96b), whose X-ray isophotes are strongly flattened, and show a large, abrupt position angle twist of $\\sim 30^\\circ$ at $\\sim$ 1-2 $R_{\\rm eff}$. While a model halo from the dissipationless scenario fits the data better than does a halo from the dissipational scenario, it is in both scenarios difficult to reproduce both the twist behavior and the isophote flatness using reasonable parameters for the galaxy and halo. The simplest explanation for NGC 720 is that the halo and the stars are intrinsically misaligned. Such an intrinsic misalignment may arise naturally for a halo which forms by dissipationless collapse, or it may be caused by late-history major mergers. Whether caused by triaxiality or by intrinsic misalignment, the shapes and orientations of stellar and X-ray isophotes are important fossil clues to the formation history of galaxies. \\vspace{1.5cm} We thank David Buote and Christine Jones for their helpful comments. CSK is supported by NSF grant AST-9401722 and NASA ATP grant NAG5-4062. \\newpage" + }, + "9708/astro-ph9708162_arXiv.txt": { + "abstract": "We use the new results of the HEGRA detector on the TeV $\\gamma$--ray emission from MKN 501 to set upper limits on the energy density of the cosmic infrared background (CIRB). Contrary to previous interpretations of the $\\gamma$--ray spectrum of MKN 421 as showing an intergalactic absorption cutoff at 5 TeV, the observed spectrum of MKN 501 extends beyond 10 TeV and appears to be unattenuated by $\\gamma\\gamma$ collisions with the low-energy CIRB photons. The upper limits on the CIRB intensity -- derived both assuming an {\\it a priori} shape for the CIRB spectrum and without model-dependent assumptions -- are thus quite strong and come almost in conflict with the observational evaluations based on deep surveys of extragalactic sources in the near- and mid-IR. If spectra at TeV energies for extragalactic $gamma$-ray sources like this for MKN 501 will be confirmed with improved statistics, we may be forced to conclude that the process of $\\gamma\\gamma$ interaction in the intergalactic space is more complex than expected and the average intergalactic magnetic field extremely weak ($B<10^{-11}\\ G$). ", + "introduction": "Cosmic history from the decoupling ($z=1500$) to the epoch of lighting of the first luminous sources (at redshifts $z\\sim 3$ to 5) is one of the biggest unknowns of present--day observational cosmology. High redshifts and dust extinction during early active phases both degrade the energetic optical--UV photons emitted by massive stars, decaying particles, or more exotic energy sources, to the infrared wavelengths. A fundamental information on the total energy budget associated with astrophysical processes occurring at high redshifts is then provided by observations of the cosmic background at infrared wavelengths (CIRB). Unfortunately, the infrared domain presents various levels of difficulty to the observational astronomer, because of the huge backgrounds from the Earth's atmosphere, the Interplanetary dust (IPD), and diffuse dust in the Milky Way, in addition to the background produced by the telescope itself. Also the sensitivity and stability of infrared detectors are far poorer than those used in the optical. Because of all this, the detection and characterization of the diffuse background flux of low-energy photons coming from primeval structures has been exceedingly difficult so far. Even dedicated experiments exploiting cooled platforms outside the atmosphere, among which the most important is DIRBE on COBE (\\cite{Hauser96}), have failed so far to detect significant signals from the CIRB above the intense foregrounds. The upper limits allowed by the foreground emission deconvolution of DIRBE maps are significantly higher than expectations over a substantial -- and crucial -- $\\lambda$--range from a few to $\\simeq 100$ $\\mu$m. The situation, at these wavelengths in particular, is not likely to improve in the future, until a mission flying to the outer Solar System will get rid of the fundamental limitation set by the IPD (both the scattered light and dust re--radiation). Under these circumstances, a very interesting alternative to the direct detection of CIRB photons has been suggested by Stecker and de Jager (1993) soon after the discovery of high-energy photon fluxes coming from distant Blazars (with GRO, \\cite{Hartman92} and with the Whipple Observatory, \\cite{Punchetal92}). The idea is to infer the CIRB spectral intensity from combined GeV and TeV observations of a set of active galactic nuclei (AGN), by exploiting the $\\gamma-\\gamma$ interactions and pair production between the AGN high--energy photons and low--energy background photons in the line--of--sight to the source. The interaction is expected to produce an absorption feature, testable in principle, in the source TeV spectrum. Interesting limits have been discussed by Stecker and de Jager (1993), de Jager, Stecker and Salamon (1994), and Dwek and Slavin (1994), {\\it all based on TeV observations of the Blazar MKN 421}. To summarize, the best current upper limits on the CIRB in the spectral range from 10 to 40 micron are those reported by de Jager, Stecker \\& Salamon (1994), with a $2\\sigma$ upper limit of $\\lambda I_\\lambda < 2\\ 10^{-8}\\ W/m^2/sr$. In the same large waveband interval, marginal detections, at levels of $2\\ 10^{-9} <\\lambda I_\\lambda < 2\\ 10^{-8} W/m^2/sr$ were reported by de Jager, Stecker \\& Salamon (1994) and Dwek \\& Slavin (1994). At shorter wavelengths, 1 to 10 $\\mu$m, upper limits have been obtained by Stecker \\& de Jager (1993), Stecker (1996), Dwek \\& Slavin (1994) and Biller et al. (1995). The most conservative bounds, accounting for the precise spectral shape of the CIRB, given by Dwek \\& Slavin (1994, $\\lambda I_\\lambda < 10^{-7}\\ W/m^2/sr$), still keep a substantial factor ($>$10) higher than the expected contribution of known sources. While all previous analyses relied uniquely on TeV observations of the blazar MKN 421, we exploit here a new high--quality dataset of TeV gamma--ray observations by HEGRA of MKN 501 during a state of high activity (Aharonian et al 1997a) to further constrain the intensity of the diffuse IR background. Two sets of constraints on the CIRB spectral intensity are derived in Section 2. One is based on the assumption that the background spectrum is dominated by the contribution of distant galaxies, and thus reflects the average galactic IR spectrum. The other constraint is free of model--dependent assumptions and treats the CIRB as a combination of twelve bins wherein the background spectrum is flat (in $\\lambda I_\\lambda$) with independent arbitrary normalizations. Extremely tight constraints on the CIRB ensue from this analysis, reflecting the missing evidence of any absorption in the TeV spectrum of MKN 501. A discussion is given in Section 3 in terms of a low average past emissivity both of galaxies and of primeval energy sources. Indeed the limits are so severe that they begin to conflict with the integrated IR flux of distant galaxies, recently detected in large numbers by ground-based and space observatories. We finally emphasize the alternative possibility that the interaction of high-energy gamma-rays with low-energy photons is more complex than previously supposed. $H_0=75\\ km/s/Mpc$ is assumed throughout the paper. ", + "conclusions": "In either case, both assuming the CIRB shape and relaxing it, the constraints on the CIRB intensity appear dramatic. Essentially the MKN 501 gamma-ray spectrum does not display the expected effect of absorption, it rather shows a simple $E^{-2.5}$ power-law spectrum. How reliable are these limits in view of the possible systematic errors of $\\sim$25\\% in the energy estimates of the HEGRA telescopes (\\cite{aharon97c})? This is very easy to estimate in the case of a flat binned $\\lambda I_\\lambda$ CIRB spectrum. The limits in Figure 3 would move upward and towards shorter wavelengths with the fractional amount of energy overestimate. Similar relaxation would occur also in the case of a more specialized model CIRB spectrum. One could use the optical depths from Table~1 to estimate the amount of relaxation. Similarly, a higher normalization of the source spectrum would relax the model--independent limits by the ratio of the two normalizations. Is this featureless spectrum of MKN 501 inconsistent with that observed for MKN 421, which is almost at the same distance? The latter has been interpreted by some authors (e.g. \\cite{Stecker96}) as showing a turnover at $\\epsilon \\simeq 3-5\\ TeV$, which was attributed to $\\gamma\\gamma$ absorption with the CIRB. In fact, new observations of this source during a high activity state do not appear to confirm the presence of absorption (Krennrich et al. 1997), and show significant counting rate above 5 TeV. So, at least during this high state, MKN 421 seems to show a power law spectrum similar to the spectrum discussed here for MKN 501. The constraints on the CIRB intensity from TeV observations of MKN 501 start to approach the \"measured\" lower limits at 2.2, 6.7 and 15 $\\mu$m given by the integrated emission of galaxies already resolved in deep integrations at those wavelengths. Deep surveys have been performed from ground in the K-band and from space by the mid--IR camera (ISOCAM, see~\\cite{Cesarskyetal96}) on the ISO satellite in the two latter bands. A summary of these \"direct\" determinations of the galaxy contribution to the CIRB, and a discussion of the related uncertainties, are given by Franceschini et al (1997) and Oliver et al. (1997). In any case, the CIRB cannot be lower than reported at these three wavelengths. Few possibilities are left. The first one is that the CIRB is very close to the limits allowed by the gamma--ray spectrum of MKN 501 observed by Aharonian et al. (1997a). This would imply a very strong constraint on any signals unrelated to the emission of distant galaxies (see e.g. \\cite{RRC88}, for a review). But, in view of the fact that the same power-law spectral shape as show in Fig. 1 for MKN 501 has been confirmed by later integrations on this source (Aharonian et al. 1997a; Protheroe et al. 1997), that apparently a similar shape is also suggested for MKN 421, and because an appreciable CIRB flux has already been detected and resolved into discrete sources, we find more likely that {\\it we have to revise our concepts about the propagation of TeV gamma-rays into the intergalactic space, and that something complicates the process}. The question is: why the propagation of TeV gamma-rays in intergalactic space should not produce the expected absorption in high energy spectra of distant sources? A possible solution could be that part of the source spectrum is regenerated in $\\gamma$--ray cascading (pair production + Inverse Compton). In such a cascading process, the $\\gamma$--ray spectrum of the source is depleted around the region of maximum absorption. If the $\\gamma$--ray emission of MKN 501 at the source extends above $10^{14}$ eV, the spectrum would be depleted in collisions with microwave background photons. The $e^+e^-$ pairs generated on the microwave background would Inverse Compton scatter on the microwave background to regenerate photons of lower (TeV) energy, thus generating `bumps' on a power-law production spectrum (\\cite{ProthSta93}). The resulting $\\gamma$--ray spectrum may then appear unattenuated at observation. This would however require not only a $\\gamma$--ray spectrum extending to very high energy, but also a very low ($\\sim 10^{-11}$ Gauss) value for the extragalactic magnetic field in the direction of MKN501. Otherwise the $e^+e^-$ pairs would deflect in the magnetic field and form a halo around the source, well outside of the angular resolution of the HEGRA detector. A good deal of constraints useful to disentangle between these two possibilities are soon expected by improved observations of the MKN501 outburst (which has been observed by the Whipple and CAT Cherenkov telescopes, \\cite{Breslinetal97}) with different energy thresholds and wavelength bands and by refined forthcoming data on MKN 421. \\\\[2truemm] {\\bf Acknowledgements.} The authors appreciate the contribution of an anonymous referee to the improvement of the paper. TS thanks J. Buckley and T.K.~Gaisser for the careful reading of the manuscript and E. Dwek for comments. The research of TS is funded in part by NASA grant NAG5--5106. \\vfill\\eject" + }, + "9708/astro-ph9708097_arXiv.txt": { + "abstract": "This paper reports the results of a near infrared spectroscopic survey of LINER galaxies undertaken with a new infrared spectrograph at the 5~m Hale telescope. The galaxy sample includes 11 LINERs with spectra covering the [FeII] (1.2567~$\\mu$m), Pa$\\beta$ (1.2818~$\\mu$m), H$_2$ ( 1-0 S(1), 2.1218~$\\mu$m) and Br$\\gamma$ (2.1655~$\\mu$m) near infrared emission lines, and one additional galaxy with only [FeII] and Pa$\\beta$ line coverage. All of the LINERs with infrared line detections have strong [FeII] and/or H$_2$ emission, with about half (4 out of 9) having extremely high ratios ($>$2) of [FeII] to Pa$\\beta$. The strength of the H$_2$ and [FeII] lines is well correlated with the optical [OI] line, with many LINERs having higher ratios of [FeII]/Pa$\\beta$, H$_2$/Br$\\gamma$ and [OI]/H$\\alpha$ than other galaxy types. The LINERs with the highest [FeII]/Pa$\\beta$ ratios (termed ``strong'' [FeII] LINERs) show evidence for recent star formation. Shocks from compact supernova remnants may enhance the [FeII] emission in these ``strong'' [FeII] LINERs. The LINERs with lower [FeII]/Pa$\\beta$ ratios (termed ``weak'' [FeII] LINERs) are more consistent with Seyfert-like activity, including higher ionization states, some strong x-ray sources and some broad H$\\alpha$ detections. The [FeII] luminosity and the [FeII]/Pa$\\beta$ ratio in these objects are more easily explained by hard x-ray excitation than in the ``strong'' [FeII] LINERs. These ``weak'' [FeII] LINERs are considered prime candidates for being low luminosity Seyfert nuclei. ", + "introduction": "LINERs (Low Ionization Nuclear Emission-line Region galaxies) are the most common and lowest energy examples of active galaxies known (Heckman 1980). The main LINER characteristic is unusually strong forbidden line transitions from low ionization states such as [OII]($\\lambda$=3727~\\AA), [NII]($\\lambda$=6583~\\AA), and [SII]($\\lambda$=6717,6731) relative to lines from higher ionization states. The original definition used by Heckman (1980) was [OII]($\\lambda$=3727~\\AA) / [OIII]($\\lambda$=5007~\\AA) $>$ 1 and [OI]($\\lambda$=6300~\\AA) / [OIII]($\\lambda$=5007~\\AA) $>$ 1/3. Classical LINER galaxies also have much less total energy in the spectral lines as compared to Seyfert galaxies and other types of active galactic nuclei (AGN); often down by a factor of a 100 in comparison to Seyfert's. The weakness of the spectral lines makes LINERs difficult to study, particularly when the galaxies often have strong stellar absorption features. Fast (v$\\sim$100~km~s$^{-1}$) shocks (Heckman, 1980), photoionization from a power-law source (e.g. Ho, Filippenko, \\& Sargent 1993) or from a cluster of very hot stars (e.g. Terlevich \\& Melnick 1985; Shields 1992), especially in very dense environments, can all duplicate the observed emission line flux ratios in LINERs. An important aspect of the star formation models, involves compact supernova remnants which are confined by the high densities within the nuclear region and which produce strong shocks. X-ray heating is particularly attractive for photoionization because hard x-rays are able to penetrate deeply into molecular clouds creating large partially ionized regions where low ionization species will dominate. An intriguing possible variation on the photoionization models recently proposed by Eracleous, Livio \\& Binette (1995), is that LINERs have a compact object (probably a black hole) in the nucleus which periodically disrupts a star during a close orbital approach. As the stellar material accretes onto the central source, high energy photons are produced and a Seyfert-like broad line region appears. As the material is consumed, the ionizing flux drops and the high ionization states weaken quickly. Low ionization lines, however, remain strong for much longer since the decay time is longer and the light crossing time in this region is much larger than that of the broad line region. This ``Duty Cycle Hypothesis'' was motivated by the observation that about 20\\% of LINERs had detectable 2300{\\AA} emission with HST (Maoz et al. 1995). These UV bright galaxies, have no other obvious difference from the UV dark LINERs. Under this theory, the UV bright galaxies would still have ongoing accretion, while the others are in the quiescent phase. In support of this theory, Eracleous et al. (1995) point to NGC 1097 which was observed to make a sudden transition from a LINER to a Seyfert 1 (Storchi-Bergmann, Baldwin \\& Wilson 1993). It is also possible that LINERs represent a heterogeneous class of objects. Some LINERs may have a central ``monster'' like Seyfert galaxies, while others have one or more dense clusters of hot young stars. Shocks may play a role in enhancing the forbidden lines in either of these two groups. Whatever the case, extending the number of observed spectral diagnostics into the infrared gives greater leverage on the problem of understanding the ultimate origin of the LINER phenomena. A significant number of important infrared spectral features exist which can provide useful astrophysical insights. Among these are the Br$\\gamma$ (2.1655$\\mu$m) and Pa$\\beta$ (1.2818$\\mu$m) hydrogen recombination lines. These lines trace ionizing photons and can be directly related to other recombination lines and to the strength of the UV continuum, while suffering much less from dust extinction than the Balmer lines. Another series of important infrared lines are the H$_2$ rotation-vibration transitions which have no analog within the optical. The strongest H$_2$ line is the 1-0 S(1) transition at 2.1218 $\\mu$m. H$_2$ emission is ubiquitous in starburst and Seyfert galaxies and is thought to originate in slow shocks, UV fluorescence and X-ray heating. The forbidden transitions of singly ionized iron ([FeII]), are also strong in the infrared. Iron is believed to play an an important role as a coolant in shocked environments (Nussbaumer and Storey 1988) but in the general interstellar medium it is often depleted compared to other elements since most iron is locked up in dust grains. The strongest near infrared [FeII] line is the 1.2567$\\mu$m transition in the J band. This line is often very strong in Seyfert galaxies and in most cases is thought to trace faster shocks ($\\sim$100 km sec$^{-1}$) than the H$_2$ lines (Graham et al. 1990). [FeII] is also thought to be strong in x-ray heated environments where dust grains have been evaporated. In this paper, we report new near infrared spectroscopy of a sample of bright LINERs, and we attempt to use these data to constrain the possible excitation mechanisms. ", + "conclusions": "This paper has described an infrared spectroscopic survey of 12 ``classical'' LINER galaxies. The spectra have concentrated on the [FeII](1.2567$\\mu$m), Pa$\\beta$, H$_2$(2.1218$\\mu$m) and Br$\\gamma$ infrared lines. The major results are: \\begin{enumerate} \\item {[FeII] and H$_2$ are the strongest infrared lines in classical LINERs. Using extrapolated H$^+$ line strengths from the optical, approximately half of the classical LINERs have ratios of [FeII]/Pa$\\beta$ and/or H$_2$/Br$\\gamma$ a factor of two or more higher than typical Seyfert galaxies and a factor of five or more higher than typical starburst galaxies.} \\item {A natural subdivision between the LINERs occurs at [FeII]/Pa$\\beta$ = 2. The four ``strong'' [FeII] LINERs exhibit evidence for recent or ongoing star formation. As a group, the [FeII] emission in these LINERs is consistent with shock excitation from compact supernova remnants. The five ``weak'' [FeII] LINERs have more in common with Seyfert galaxies, including the only two objects in the sample with broad H$\\alpha$ (NGC~3998 and NGC~4258). The lower [FeII]/Pa$\\beta$ and H$_2$/Br$\\gamma$ ratios of ``weak'' [FeII] LINERs are consistent with hard x-ray heating from a power-law source. Neither excitation mechanism is ruled out for either type of object, however, and it is possible that it is just the relative strengths of the two mechanisms that are different in the two groups.} \\item {In most of the LINERs, the estimated amount of Pa$\\beta$ absorption from the stellar population can account for the lack of Pa$\\beta$ detected in emission. Several of the LINERs have unusually strong Pa$\\beta$ absorption ( $>$2$\\AA$ EQW) indicative of younger stellar populations. The five galaxies with the strongest absorption all have strong [FeII] lines and supernova remnants from recent star formation are a plausible explanation for their enhanced [FeII] strengths.} \\item {A strong linear correlation is observed between H$_2$/Br$\\gamma$, [FeII]/Pa$\\beta$ and [OI]/H$\\alpha$ for a range of 100 in all ratios and for all of the observed galaxy types. However, the ``strong'' [FeII] LINERs are stronger in [FeII] than the correlation with [OI] predicts. Both the H$_2$/Br$\\gamma$ and to a lesser degree the [FeII]/Pa$\\beta$ ratio, appear to separate the three galaxy classes: LINERs, Seyferts and starbursts. Only six objects have been observed to have H$_2$/Br$\\gamma$ higher than 3, and all are LINERs. In conjunction with the [OIII]/H$\\beta$ ratio, the class distinctions are even clearer.} \\item{The shallow far infrared spectral slopes of some of the ``weak'' [FeII] LINERs in comparison to the strengths of H$_2$ and [OI], appear inconsistent with the correlations observed for starburst and ULIRG galaxies. This argues very strongly that a nonthermal heat source is providing much of the dust heating and line excitation for these objects.} \\item {LINERs with x-ray detections appear to have sufficient x-ray luminosity to power the observed infrared lines. Starbursts often have [FeII] and H$_2$ lines many times stronger than the x-rays luminosity should be able to produce. Again the strong x-ray LINERs do not appear consistent with star formation, but instead behave like mini-Seyferts.} \\end{enumerate}" + }, + "9708/astro-ph9708005_arXiv.txt": { + "abstract": "We have found the first convincing evidence for spiral structure in the accretion disc of a close binary. The eclipsing dwarf nova binary IP Peg, observed during the end phase of a rise to outburst, shows strong Balmer and Helium emission lines in its spectra, with asymmetric double peaked velocity profiles produced in the accretion disc around the white dwarf. To reveal the two armed spiral on the accretion disc, we de-project the observed emission line profiles onto a Doppler coordinate frame, a technique known as Doppler tomography. The two armed spiral structure we see in the Doppler tomograms is expected to form when the disc becomes sufficiently large in outburst so that the tides induced by the secondary star can excite waves in the outer disc. Such spiral waves have been predicted in studies of tidal effects in discs and are fundamental in understanding the angular momentum budget of accretion discs. ", + "introduction": "IP Pegasi is an interacting binary system containing a white dwarf receiving mass through an accretion disc from a Roche lobe filling late type star. These accretion disc fed systems called cataclysmic variables (see Warner (1995) for an excellent overview) provide one of the best laboratories for accretion physics due to their proximity and convenient time scales. The strong emission lines in their spectra originate in the accretion flow and are powerful observational probes of the local gas conditions. The picture of a viscous disc, transporting angular momentum outwards as material slowly spirals inwards, forms the basis of our understanding of accretion flows in X-ray binaries and AGNs as well. One of the main longstanding problems of accretion discs is their angular momentum transport mechanisms. In order to sustain the observed mass transfer rates highly efficient viscous processes must be available to transport the angular momentum outwards. Although the famous $\\alpha$ prescription (Shakura \\& Sunyaev 1973), which scales the effective viscosity by a dimensionless parameter $\\alpha$, has been very succesfull it also shows how poorly these processes are understood. Turbulent magnetic fields (Tout \\& Pringle 1992, Schramkowski \\& Torkelsson 1996) and spiral shocks (Spruit et al. 1987) are two promising mechanisms even though the effective $\\alpha$ expected from such models is still low. A second issue that has received less attention is the removal of the angular momentum at the outer disc via a tidal torque between disc and companion star (e.g. Papaloizou \\& Pringle 1977). IP Peg is a member of the subclass of CVs called dwarf novae that display semi-periodic outbursts during which the system brightens by several magnitudes as more mass is suddenly transfered through the disc. These systems provide a great test case for accretion disc models. IP Peg is one of the few eclipsing dwarf novae, where the inclination of the orbital plane ($\\sim$80$^{\\circ}$) is large enough for the 0.5 M$_{\\odot}$ companion star to cover the 1.02 M$_{\\odot}$ white dwarf and most of the accretion disc as it passes in front every 3.8 hours. IP Peg's outbursts have an amplitude of about 2 magnitudes and recur roughly every 3 months during which the accretion disc is the dominant light source. We present spectrophotometric observations of the dwarf nova IP Peg at the late stages of a rise to outburst and use Doppler imaging to map the accretion disc. Observations are presented in section 2 followed by the analysis of the tomograms in section 3. The tidal origin of the spirals is discussed in section 4. ", + "conclusions": "The tidal interaction manifestated in the spiral pattern turns out to be an important factor for outburst discs. Work is now in progress to use different observations of this phenomenon in different emission lines and at different epochs to sample the physical conditions of the disc material. Observing high ionization lines like HeII can show the presence of shocks and will indicate the implication for the angular momentum budget. Furthermore future observations of disc structure in different objects (with different mass ratios and disc sizes) will provide us with a new insight in tidal theory and perhaps the outburst mechanism. In this way dwarf novae disc provide an excellent laboratory for tides in astrophysical discs, since the time scales of the outbursts lasting a week and recurring every couple of months, allows one to study the dynamical behaviour of the disc and its tidal response. Tidal spirals in galaxies for example, thought to be generated in the same manner by a companion galaxy, have very long dynamical time scales making it impossible to study their evolution." + }, + "9708/hep-ex9708039_arXiv.txt": { + "abstract": "The Alpha Magnetic Spectrometer (AMS) is a state of the art detector for the extraterrestrial study of antimatter, matter and missing matter. After a precursor flight on STS91 in may 1998, AMS will be installed on the International Space Station where it will operate for three years. In this paper the AMS experiment is described and its physics potential reviewed. ", + "introduction": "The disappearance of cosmological antimatter and the pervasive presence of dark matter are two of the greatest puzzles in the current understanding of the universe. The Big Bang model assumes that, at its very beginning, half of the universe was made out of antimatter. The validity of this model is based on three key experimental observations: the recession of galaxies (Hubble expansion), the highly isotropic cosmic microwave background and the relative abundances of light isotopes. However, a fourth basic observation, the presence of cosmological antimatter somewhere in the universe, is missing. Indeed measurements of the intensity of gamma ray flux in the MeV region exclude the presence of a significant amount of antimatter up to the scale of the local supercluster of galaxies (tens of Megaparsecs). It follows that, either antimatter has been destroyed immediately after the Big Bang by some unknown mechanism, or matter and antimatter were separated (by some other unknown mechanism) into different region of space, at scales larger than superclusters. All efforts to reconcile the the absence of antimatter with cosmological models that do not require new physics failed (see \\cite{Steigmann,Kolb,Peebles}, for a review of these theories). We are then currently unable to explain the fate of half of the baryonic matter present at the beginning of our universe. Rotational velocities in spiral galaxies and dynamical effects in galactic clusters provide us convincing evidence that, either Newton laws break down at scales of galaxies or, more likely, most (up to 99\\%) of our universe consists of non-luminous (dark) matter. There are several dark matter candidates. They are commonly classified as ''hot'' and ''cold'' dark matter, depending on their relativistic properties at the time of decoupling from normal matter in the early universe. As an example, light neutrinos are obvious candidates for ''hot'' dark matter while Weakly Interacting Massive Particles (WIMP's) are often considered as ''cold'' dark matter candidates \\cite{Ellis,Turner}. \\begin{figure}[htb] \\begin{center} \\mbox{\\epsfig{file=y96192aUSShutCut,width=7 cm}} \\caption{\\em {AMS on STS 91 (Discovery)}} \\end{center} \\end{figure} \\begin{figure*}[htb] \\begin{center} \\mbox{\\epsfig{file=Blank,width=14cm}} \\caption{\\em {The International Space Station Alpha; AMS will be installed on the left side of the main truss}} \\end{center} \\end{figure*} In either cases we are currently unable to explain the origin of most of the mass of our universe. To address these two fundamental questions in astroparticle physics a state of the art detector, the Alpha Magnetic Spectrometer (AMS) \\cite{AMS} has been recently approved by NASA to operate on the International Space Station Alpha (ISSA). AMS is manifested for a precursor flight with STS91, (Discovery, may 1998, Figure 1), and for a three year long exposure on the International Space Station (ISS) (Figure 2), after its installation during Utilization Flight n.4 (Discovery, january 2002). AMS has been proposed and is being built by an international collaboration involving China, Finland, France, Germany, Italy, Rumenia, Russia, Switzerland, Taiwan and US. \\begin{figure}[htb] \\begin{center} \\mbox{\\epsfig{file=y97087dabisBattistonGray-EPSF,width=5.5cm}} \\vspace{0.1cm} \\caption{\\em {Astroparticle physics}} \\label{fig:astroparticle} \\end{center} \\end{figure} \\clearpage \\cleardoublepage ", + "conclusions": "" + }, + "9708/astro-ph9708139_arXiv.txt": { + "abstract": "We use ROSAT PSPC data to study the X-ray properties of a sample of twelve poor groups that have extensive membership information (Zabludoff \\& Mulchaey 1997; Paper I). Diffuse X-ray emission is detected in nine of these groups. In all but one of the X-ray detected groups, the X-ray emission is centered on a luminous elliptical galaxy. Fits to the surface brightness profiles of the X-ray emission suggest the presence of {\\it two} X-ray components in these groups. The first component is centered on the central elliptical galaxy and is extended on scales of 20--40 h$_{\\rm 100}$$^{-1}$ kpc. The location and extent of this component, combined with its X-ray temperature ($\\sim$ 0.7--0.9 keV) and luminosity ($\\sim$ 10$^{41-42}$ h$_{\\rm 100}$$^{-2}$ erg s$^{-1}$), favor an origin in the interstellar medium of the central galaxy. Alternatively, the central component may be the result of a large-scale cooling flow. The second X-ray component is detected out to a radius of at least $\\sim$ 100--300 h$_{\\rm 100}$$^{-1}$ kpc. This component follows the same relationships found among the X-ray temperature (T), X-ray luminosity (L$_{\\rm X}$) and optical velocity dispersion($\\sigma$$_{\\rm r}$) of rich clusters. This result suggests that the X-ray detected groups are low-mass versions of clusters and that the extended gas component can properly be called the intragroup medium, in analogy to the intracluster medium in clusters. The failure to detect an intragroup medium in the three groups with very low velocity dispersions is consistent with their predicted X-ray luminosities and temperatures based on the relationships derived for clusters and X-ray detected groups. The best-fit value of $\\beta$ derived from the $\\sigma$$_{\\rm r}$-T relationship for groups and clusters is $\\sim$ 0.99$\\pm{0.08}$, implying that the galaxies and hot gas trace the same potential with equal energy per unit mass and that the groups are dynamically relaxed. We also find a trend for the position angle of the optical light in the central elliptical galaxy to align with the position angle of the large-scale X-ray emission. This trend is consistent with that found for some rich clusters containing cD galaxies (Rhee, van Haarlem \\& Katgert 1992; Sarazin et al. 1995; Allen et al. 1995). The alignment of the central galaxy with the extended X-ray emission suggests that the formation and/or evolution of the central galaxy is linked to the shape of the global group potential. One possible scenario is that the central galaxy formed via galaxy-galaxy mergers early in the lifetime of the group and has not been subject to significant dynamical evolution recently. ", + "introduction": "Because most galaxies occur in small groups, understanding the physical nature of these systems is critical for cosmology. An outstanding question is whether poor groups are simply low-mass versions of richer clusters or physically different systems. Establishing the nature of groups has proved difficult because these systems typically contain only three or four bright galaxies. Thus, it is not even known if the majority of the cataloged groups are real, bound systems or if they are simply chance superpositions (e.g., Ramella et al. 1989; Hernquist et al. 1995). The presence of diffuse X-ray emission provides evidence for a common potential in some poor groups (e.g., Mulchaey et al. 1993, Ponman \\& Bertram 1993, David et al. 1994, Ebeling et al. 1995, Pildis et al. 1995, Henry et al. 1995, Mulchaey et al. 1996a, Ponman et al. 1996, Burns et al. 1996). Many ROSAT studies suggest there is a correlation between the morphological composition of the group and the existence of X-ray emission (e.g., Ebeling et al. 1995, Pildis et al. 1995, Henry et al. 1995, Mulchaey et al. 1996a), providing further evidence for the reality of the X-ray detected groups. However, it is still possible that the X-ray gas is merely a projection of unbound gas in filaments along the line-of-sight (Hernquist et al. 1995). Perhaps the best way to establish the physical reality of groups is to extend the kinematic properties of these systems to include much fainter members. We have recently completed a fiber spectroscopy study of a sample of twelve poor groups (Zabludoff \\& Mulchaey 1997; Paper I) and have shown that the X-ray detected groups are indeed bound systems. Even if the reality of some groups has been demonstrated, there is still controversy about the origin of the X-ray emission in these systems. For example, both Dell'Antonio et al. (1994) and Mahdavi et al. (1997) find a much flatter relationship between X-ray luminosity and velocity dispersion ($\\sigma$$_{\\rm r}$) for groups than is found for rich clusters. These authors argue that the X-ray emission in groups is dominated by emission from individual galaxies and not from a global group potential. In contrast, Ponman et al. (1996) remove the X-ray emission from individual galaxies and find reasonable agreement between the L$_{\\rm X}$--$\\sigma$$_{\\rm r}$ relationship for Hickson Compact Groups (HCGs) and rich clusters. Henry et al. (1995) and Burns et al. (1996) also find that the X-ray luminosity function is a natural extension of the relationship for clusters, implying a similar physical mechanism for the X-ray emission. However, both the Henry et al. (1995) and Burns et al. (1996) studies studies only include the most X-ray luminous groups and not the more typical, low luminosity systems where the differences in the L$_{\\rm X}$--$\\sigma$$_{\\rm r}$ relationship have been reported. Some of the uncertainty in the L$_{\\rm X}$--$\\sigma$$_{\\rm r}$ relationship and in other comparisons of optical and X-ray properties may result from poorly determined optical properties. For most groups, velocity dispersions have been estimated from as few as three or four galaxy velocity measurements. However, in Paper I we showed that calculating $\\sigma$$_{\\rm r}$ from only a few galaxies can result in large uncertainties in the measurement, often underestimating $\\sigma$$_{\\rm r}$ by a factor of 1.5 or more. Thus, the values of $\\sigma$$_{\\rm r}$ used in many previous comparisons of the optical and X-ray properties of groups may be flawed. The X-ray observations of poor groups are also subject to large uncertainties that in general do not plague similar measurements in rich clusters. For example, diffuse emission from the global potential dominates the cluster X-ray emission. However, in poor groups, the contribution of X-rays from individual galaxies may be substantial. The presence of luminous X-ray emission in poor groups is found predominantly in systems with a central elliptical galaxy (e.g., Mulchaey et al. 1996a). The expected X-ray emission from such a galaxy can be comparable to the diffuse, extended emission observed in many groups. Thus, it is not clear whether the observed emission is associated with an extended halo of the central galaxy (e.g., Trinchieri et al. 1997) or represents hot gas in the group's global potential. If the observed emission is dominated by a component from the central galaxy's interstellar medium, comparisons with global group properties such as velocity dispersion may not be particularly meaningful. In Paper I, we showed that X-ray detected groups typically contain at least $\\sim$ 20-50 members to absolute magnitudes of M$_{\\rm B}$ $\\sim$ -14 to -16 + 5 log$_{\\rm 10}$h$_{\\rm 100}$. With this many galaxy velocities available, we are able to calculate robust velocity dispersions for poor groups. Here we use deep ROSAT observations of the twelve groups studied in Paper I to examine the nature of the X-ray emission in these systems. In the next section, we describe the ROSAT observations and data reduction techniques. In \\S3, we study the spatial distribution of the X-ray emission and discuss the evidence for two X-ray components in these groups. Spectral analysis of the ROSAT data, including temperature and luminosity measurements, are presented in \\S4. Comparisons of the derived optical and X-ray properties of groups are given in \\S5, and our conclusions are summarized in \\S6. All distance-dependent quantities in this paper are calculated assuming H$_{\\rm o}$ = 100 km s$^{-1}$ Mpc$^{-1}$ (h=H$_{\\rm o}$/100). ", + "conclusions": "We have presented ROSAT PSPC observations of twelve groups with detailed galaxy membership (Paper I). Diffuse X-ray emission is found in nine groups. The most luminous galaxy in each of the X-ray detected groups is an elliptical, whose position is coincident with the peak of the X-ray emission in all but one case. Surface brightness profiles of the X-ray emission strongly suggest the presence of two components in these groups: one on scales of $\\sim$ 20--40 h$^{-1}$ kpc and one on much larger scales (of at least $\\sim$ 100--300 h$^{-1}$ kpc). The temperatures of the central and extended components are significantly different in some groups, consistent with the interpretation that the components are distinct. The extent, temperature ($\\sim$ 0.7--0.9 keV) and luminosity ($\\sim$ 10$^{41-42}$ h$^{-2}$ erg s$^{-1}$) of the first component is consistent with that observed in elliptical galaxies in other environments, suggesting that this component most likely originates in the interstellar medium of the central galaxy. Alternatively, the central component may be the result of a cooling flow. The extended X-ray component follows the extrapolation of the relationships found among velocity dispersion, X-ray temperature and X-ray luminosity for rich clusters. This suggests X-ray detected poor groups can be thought of as scaled-down versions of clusters, with the extended X-ray component in groups representing the intragroup medium, in analogy to the intracluster medium in clusters. The best fit to the $\\sigma$$_{\\rm r}$-T relationship for X-ray groups and clusters gives a mean $\\beta$ value of 0.99$\\pm{0.08}$, suggesting the galaxies and hot gas trace the same potential and that the energy per unit mass in the gas and galaxies is equal. The values of $\\beta$ derived independently from fits to the surface brightness profiles are consistent with $\\beta$ $\\sim$ 1 in many groups. The lower values of $\\beta$ implied from other studies may be due to significant underestimates of the group velocity dispersion and to contamination of the surface brightness profiles by the central galaxy X-ray emission. While the ROSAT data suggest the presence of two independent X-ray components in these groups, it is clear that the formation and/or evolution of the central galaxy is somehow linked to the extended group potential. In particular, there is a strong trend for the optical isophotes of the central galaxy to align with the X-ray emission isophotes on large scales. This result might be expected if the central galaxy formed via galaxy-galaxy mergers early in the lifetime of the group and has not been recently disturbed. The fact that the central galaxy is at rest in the center of the group's potential (Paper I) is consistent with this scenario. A similar phenomenon has been observed in some clusters (e.g., Rhee, van Haarlem \\& Katgert 1992; Sarazin et al. 1995; Allen et al. 1995), offering further evidence of the similarities between X-ray luminous, poor groups and rich clusters. The results of this paper demonstrate the insight that can be gained by combining detailed group membership information with quality X-ray observations. Still, there are many outstanding questions that can be addressed with further work. Our sample contains few non-X-ray detected or low temperature groups. We are in the process of extending our spectroscopy program to include more of these systems. Significant improvements in the X-ray observations can also be expected. Many groups have been observed with ASCA, which should help determine the metallicity of the intragroup medium, providing further constraints on the origin of this component. Higher spatial resolution observations will also allow the central galaxy-group interface to be studied in much more detail." + }, + "9708/astro-ph9708249_arXiv.txt": { + "abstract": "We present an improved calibration of photometric metallicity indicators, derived from the new metallicity scale for Globular Clusters presented by Carretta \\& Gratton (1997) and based on direct high resolution spectroscopy of 160 stars in 24 globular clusters. We have carefully recalibrated the traditional abundance indices based upon the red giant branch (RGB) morphology, both in the $V,B-V$ and $V,V-I$ planes, namely the dereddened colour at the luminosity level of the horizontal branch (HB), and the magnitude difference between the HB and the RBG at a given dereddened colour. Finally, we give new accurate relations to employ in the Simultaneous Metallicity Reddening method by Sarajedini (1994), also tied to the Carretta \\& Gratton (1997) abundance scale. ", + "introduction": "Since globular clusters (GCs) are among the oldest objects in galaxies, they are widely recognized as very useful tracers of the chemical and dynamical evolution of their parent hosts. The accurate knowledge of their global metal content, measured by the [Fe/H] ratio, is critical for many astrophysical problems. In particular, being very massive and luminous systems of coeval stars that show, to a first approximation, a similar (initial) chemical composition, globular clusters represent the cornerstones in establishing the existence of an age-metallicity relation and/or a metallicity-galactocentric distance gradient, up to the most distant regions of the galactic halo. This in turn provides strong constraints on models of galactic formation. Moreover, variations in the [Fe/H] content among globular clusters can be interpreted as a fossil record of the global processes of chemical enrichment occurred through the history of the Galaxy. Finally, precise metallicities are one of the basic ingredients in deriving accurate ages using parallaxes measured by the Hipparcos astrometry satellite (see Gratton et al. 1997; Reid 1997). Even if the best way to get a quantitatively accurate estimate of the metal abundance of any star is detailed abundance analysis of high resolution spectra, there are unfortunately some shortcomings that limit the application of this technique to the study of GCs. Due to their large distances, reliable high resolution, high signal-to-noise spectra can be obtained with the present day instrumentation only for the brightest giants. Only an handful of stars near the main sequence turn-off (hence reflecting the initial chemical composition, undisturbed by mixing in later evolutionary phases) have been observed yet. Moreover, high-resolution spectroscopy is a very time consuming observing technique. Therefore, in the past years, a number of indirect metallicity indicators have been devised to overcome these problems. Almost all of them are based on integrated parameters that bypass the distance limit also for very far clusters, but they require a very accurate calibration in order to provide the true content in [Fe/H]. A direct calibration, able to tie the observed photometric indices to the actual number of iron atoms as measured from high resolution spectral line profiles is henceforth strongly needed. In Section 2 we discuss the philosophy of our approach; Section 3 and 4 are devoted to the presentation and discussion of new calibrations for several metallicity indicators; a short summary is presented in Section 5. \\begin{table*} \\caption{GCs used as calibrators. Data are taken from GC97, ZW, SL (tab. 5), Sarajedini (1994, tab.1), and Sarajedini \\& Milone (1995, for NGC5053 and NGC4590). Y or N indicate in columns 3 and 9 whether the GC is a primary calibrator for GC97 and SL respectively.} \\begin{tabular}{lcccccccccc} \\hline &&&&&&&&&&\\\\ GC &[Fe/H] &Direct &[Fe/H] &E(B-V) &(B-V)$_{0,g}$ &$\\Delta V_{1.2}$ &$\\Delta V_{1.1}$ & calib. &(V-I)$_{0,g}$ &$\\Delta V_{1.2}$ \\\\ & CG97 &spectr.? &ZW & & &$V,B-V$ &$V,B-V$ &in SL? & &$V,V-I$ \\\\ &&&&&&&&&&\\\\ \\hline &&&&&&&&&&\\\\ NGC104 (47Tuc) &-0.70 &Y &-0.71 &0.04 &0.958 &1.275 &0.798 &Y &1.032 &1.028 \\\\ NGC288 &-1.07 &Y &-1.40 &0.02 &0.852 &1.884 &1.503 &N && \\\\ NGC362 &-1.15 &Y &-1.28 &0.03 &0.832 &2.050 &1.741 &N && \\\\ NGC1261 &-1.09 &N &-1.31 &0.00 &0.860 &2.095 &1.735 &N && \\\\ NGC1851 &-1.08 &N &-1.29 &0.02 &0.873 &1.806 &1.433 &Y &0.953 &1.609 \\\\ NGC1904 &-1.37 &Y &-1.69 &0.01 &0.801 &2.332 &2.001 &N && \\\\ NGC4590 (M68) &-1.99 &Y &-2.09 &0.07 &0.694 &2.810 &2.508 &Y &0.885 &2.470 \\\\ NGC5053 &-2.43 &N &-2.41 &0.06 &0.647 &3.101 &2.741 &Y &0.847 &2.770 \\\\ NGC6352 &-0.64 &Y &-0.60 &0.21 &0.994 &0.953 &0.591 &N && \\\\ NGC6397 &-1.82 &Y &-1.91 &0.18 &0.717 &2.842 &2.483 &N &0.904 &2.330 \\\\ NGC6535 &-1.53 &N &-1.75 &0.44 &0.745 &2.537 &2.134 &N && \\\\ NGC6752 &-1.42 &Y &-1.54 &0.04 &0.781 &2.264 &1.873 &Y &0.949 &1.935 \\\\ NGC7078 (M15) &-2.12 &Y &-2.17 &0.10 &0.691 &2.972 &2.601 &Y &0.882 &2.538 \\\\ NGC7089 (M2) &-1.34 &N &-1.58 &&&&& & 0.934 &2.039\\\\ Eridanus &-1.18 &N &-1.41 &0.03 &0.838 &1.897 &1.596 &N && \\\\ ESO121 &-0.83 &N &-0.93 &0.03 &0.907 &1.507 &1.101 &N && \\\\ Lindsay1 &-0.94 &N &-1.10 &0.04 &0.864 &1.888 &1.522 &N && \\\\ Pal14 &-1.36 &N &-1.60 &0.05 &0.803 &2.288 &1.939 &N && \\\\ &&&&&&&&&&\\\\ \\hline \\end{tabular} \\end{table*} \\begin{figure*} \\hspace{3cm}\\resizebox{12cm}{17cm}{\\includegraphics{6618f1.ps}} \\caption{Calibration of the $(B-V)_{0,g}$ parameter using the CG97 metallicity scale and the cluster sample by SL. The upper panel shows the calibrations obtained using only the 6 SL primary calibrators (solid line) and all the 17 clusters (dotted line). In the lower panel the calibration is based only upon the 9 clusters that have metallicities derived by CG97 from direct analysis (CG97 reference clusters). In both panels filled symbols represent clusters with [Fe/H]'s derived in CG97, while open symbols represent clusters with ZW metallicities corrected to the CG97 scale. Triangles, filled or open, indicate the 6 SL primary calibrators.} \\end{figure*} \\begin{figure*} \\hspace{3cm}\\resizebox{12cm}{17cm}{\\includegraphics{6618f2.ps}} \\caption{Calibration of the $\\Delta V_{1.2}$ parameter of SL using the CG97 metallicities. The meaning of symbols is as in Figure 1.} \\end{figure*} ", + "conclusions": "We used the new and homogeneous metallicity scale, derived by CG97 from updated model atmospheres and direct detailed abundance analysis of high resolution spectra of globular cluster giants, to calibrate the traditional RGB photometric indicators in terms of the [Fe/H] ratio. Especially in the $V,B-V$ plane, the relations found provide very good relative determinations of [Fe/H] with an uncertainty on a single measurement of 0.08~dex on average, and as low as 0.04~dex. This goes towards lessening the disagreement existing in the past between photometric and spectroscopic determination of the metal abundance in globular clusters." + }, + "9708/astro-ph9708157_arXiv.txt": { + "abstract": "{Hubble Space Telescope Observations of Black Hole X-Ray Transients are discussed in the context of the disk instability outburst model. We focus on the multiwavelength campaign following GRO~J1655-40 through the summer 1996 outburst.} ", + "introduction": "As soon as the class was discovered, the obvious similarities between the Black Hole X-Ray Transients (BHXRTs) and their white dwarf analogues, dwarf novae (DN), guided investigations into the mechanisms responsible for the dramatic outbursts exhibited by the former. The outbursts in DN have been successfully explained as the result of temperature-dependent viscosity in the accretion disk: the Disk Instability Model (DIM) (Cannizzo 1993). The longer recurrence timescales for BHXRTs and the shapes and durations of their outburst lightcurves, however, provide a challenge to the DIM (Lasota 1996). The DIM makes definite quantitative predictions for the temperature distribution, and hence the expected broad band spectrum, throughout the outburst cycle ({\\it e.g.} Cannizzo, Chen, \\& Livio 1995). Accretion disk emission is likely to dominate in the UV, so one of the primary motivations for spectroscopic observations of BHXRTs with HST is, therefore, to observe the broad band spectral evolution, and hence address the question of the driving mechanism for the transient outbursts. This paper reviews the UV-optical spectra of BHXRTs obtained with HST, and describes the consequent deductions about the outburst mechanisms. ", + "conclusions": "The observations described herein have provided some intriguing challenges to the theoretical models for transient outbursts. As we collect detailed multi-epoch data on more systems it is becoming clear that the BHXRTs exhibit complex and diverse behavior, and it seems we need to consider a variety of mechanisms in order to properly understand the wealth of observational phenomena." + }, + "9708/astro-ph9708227_arXiv.txt": { + "abstract": "We present deep K-band imaging at the positions of four very faint X-ray sources found in the UK {\\rosat} Deep Survey \\cite{mn_nov} to have no optical counterpart brighter than R$\\sim$23. Likely identifications are found within the {\\rosat} error circle in all four fields with R$-$K colours of between 3.2$\\pm$0.4 and 6.4$\\pm$0.6. From a consideration of the R$-$K colours and X-ray to optical luminosity ratios of the candidate identifications, we tentatively classify two of the X-ray sources as very distant ($z\\sim 1$) clusters of galaxies, one as a narrow emission line galaxy and one as an obscured QSO. ", + "introduction": "Optical identification of sources detected in deep X-ray images, particularly from the {\\rosat} satellite, have, in recent years, proved a highly successful probe of the origins of the soft Cosmic X-ray background (XRB) \\cite{Shanks+91,Boyle+94,mn_nov}. In particular, it has been seen that at bright fluxes, the major class of contributors is QSOs. However, they are unlikely to contribute more than $\\sim$50\\% of the X-ray background at 1keV \\cite{Boyle+94,Jones+96,mn_nov} and, even were their contribution to increase greatly beyond the limits of current surveys, their characteristic X-ray spectrum is too soft to match the residual XRB \\cite{R-C+96}. Deeper X-ray surveys, however, are showing that at fainter fluxes, there is a new class of galaxies with narrow optical emission lines (NELGs) which have suitably hard X-ray spectra \\cite{McHardy+95/6,R-C+96,Almaini+96}. The deepest X-ray survey yet optically identified is the UK {\\rosat} Deep Survey. This reaches a depth of $2 \\times 10^{-15}$ erg cm$^{-2}$ s$^{-1}$ (0.5-2keV) and is described in detail in \\scite{mn_nov}. Within the ``complete'' sample area defined by \\scite{mn_nov}, there are eleven unidentified sources (out of a total of 70 sources). Of the unidentified sources, most (7) remain unidentified because adequate optical spectra have not been obtained for all likely candidates. However, for four of the X-ray sources there are no optical candidates brighter than ${\\rm R}\\sim 23$. This is a puzzle since they are not primarily the faintest X-ray objects in the survey. One possibility is that the source of the X-rays are obscured AGN~--~long proposed as a contributor to the XRB (see for example \\pcite{Shanks+96}). The discovery of such objects would have important implications for the origin of the hard X-ray background. The X-rays could also be coming from distant clusters of galaxies, which can have high X-ray to optical luminosity ratios (see, for example, \\pcite{Stocke+91}). Other possibilities include further NELG objects or very high redshift QSOs, where the bright Ly$\\alpha$ emission has been redshifted out of the R-band. All of these classes of objects, particularly the obscured AGN, will have red optical/infra-red colours, so in 1996 we performed deep K-band imaging at UKIRT of the four `blank' Deep Survey fields to identify candidate objects and try to determine their nature. The results of this imaging are presented in this paper. In section~\\ref{sec:survey} we briefly describe the X-ray and optical data that make up the deep survey, in section~\\ref{sec:irdata} we present the K-band images and in~\\ref{sec:ident}, we consider possible scenarios. Finally in section~\\ref{sec:conclusions} we discuss the possible significance of the results. ", + "conclusions": "\\label{sec:conclusions} We have used deep K-band imaging to attempt to identify the optical candidates of four faint X-ray sources with `blank' optical fields (ie nothing within $\\sim 15$ arcsec of the X-ray centroid with R$\\la$23). Of the four fields, firm candidates have been found in three. The remaining field (object r112; figure~\\ref{fig:r112_piccies}) shows a collection of similar, red objects in the vicinity of the X-ray centroid. No firm conclusions about the nature of the candidates can be achieved with just the current photometric data but we have, nevertheless, been able to draw tentative conclusions. Of the four X-ray sources, r73 has the least red candidate. In fact, both its R$-$K colour (3.2) and X-ray to optical luminosity ratio are entirely consistent with the Narrow Emission Line Galaxies (NELGs) found in significant quantities at the fainter end of the UK {\\rosat} Deep Survey \\cite{mn_nov}. On the other hand, both r112 and r130 have candidates that are redder than in r73 (r112 has an R$-$K of 4.2, r130 has R$-$K of 6.1). These objects are probably better described by high redshift ($z\\sim 1$) galaxy clusters, the candidates having absolute magnitudes, R$-$K colours and {\\lxlopt} ratios consistent with typical $z\\sim 1$ brightest cluster galaxys (BCGs). The collection of red objects in r112 also supports this conclusion. The candidate object for source r82 also is also consistent with being a distant brightest cluster galaxy. However, in this case both the deep R and K-band images show an apparently point-like object. It may therefore also be explained as a distant ($z\\sim 2$) intrinsically reddened AGN. If this is the case, with this deep {\\rosat} observation, and others like it, we are just starting to see the tip of the iceberg which has been proposed as a significant contributor to the cosmic X-ray background, not just in the soft band, but over a wide range of energies (eg~\\pcite{Madau+94}). Future observations will be required to confirm these tentative identifications. In particular, infra-red spectroscopy over a range of wavelengths would enable one to search for broad, redshifted emission lines characteristic of QSOs (a QSO with a redshift of $z\\sim 2$ would have H$\\alpha$ redshifted into the H or K band). Further IR photometry in other bands to give IR colours may also prove a useful diagnostic and deeper imaging with good seeing would allow spatial extension to be seen if any of the objects are NELGs or moderate redshift clusters (ie~$z\\ll 1$). Deep IR imaging, therefore, has shown itself to be a powerful tool in the study of the faintest X-ray sources. In all four of the optically blank fields containing X-ray sources which we have observed in the K-band, firm candidates have been found. In addition, consideration of the photometric and X-ray properties of the candidates has enabled us to make preliminary identifications of the types of objects producing the X-ray emission, and possibly identified an obscured QSO~---~a much hypothesised but little observed contributor to the cosmic X-ray background." + }, + "9708/astro-ph9708011_arXiv.txt": { + "abstract": " ", + "introduction": "The Magellanic Clouds (MCs) are an ideal testing ground for theories of the late stages of stellar evolution. They are nearby and yet far enough that to first order the depth of the MCs may be neglected, and all stars can be considered to be at the same distance (approximately 50 and 63 kpc for the LMC and SMC, respectively). One of the primary observables is the luminosity function of oxygen-rich and carbon-rich AGB stars. Until the late-eighties AGB stars were searched for in the optical, either spectroscopically (using grisms or direct spectroscopy of red stars) or by identifying variable stars in the appropriate period range. Infrared observations were usually done as follow-up. When the IRAS Point and Faint source catalogues became available people started identifying AGB stars from those catalogues. These searches find a different population of AGB stars that are more luminous and more redder and that have been missed by the previous optical searches. To illustrate this, I have used a dust radiative transfer model to calculate the expected magnitudes of a short period, low luminosity and a long period, high luminosity carbon Mira at the distance of the LMC (see Table~1). The estimates of the luminosities and mass loss rates are from the period-luminosity relation of Groenewegen \\& Whitelock (1996) and the period-mass loss rate relation of (Groenewegen et al. 1997). It is clear that low and intermediate luminous carbon stars (but similarly for O-rich AGB stars) are bright in the optical but are far below the detection limit of IRAS, which was approximately 200 mJy at 12 $\\mu$m. Stars more luminous than approximately 10~000 \\lsol, or less luminous stars with a very high mass loss rate could have been detected by IRAS. In addition one should not forget that most AGB stars are large-amplitude LPV's (long period variables). The full-amplitude of Miras are \\more 2.5 mag in $V$, \\more 0.9 mag in $I$ and \\more 0.4 mag in $K$ and at 12$\\mu$m. IRAS could therefore also have detected less luminous stars that happened to be near maximum light at the time of the IRAS observations. It is clear that both the optical surveys and IRAS must be biased. From Table~1 it appear that the near-infrared (NIR) region may be the most unbiased way to search for AGB stars. In fact, there are two projects underway that will do so, namely the NIR sky-surveys DENIS and 2MASS. In addition ISO provides the opportunity the make detailed observations of individual objects with unprecedented sensitivity and at wavelengths not accessible from the ground. In Sect.~2 the surveys that have been conducted to identify AGB stars are described. In Sect.~ 3 and 4 first results from DENIS and ISO are presented. I conclude in Sect.~5. \\begin{table}[htb] \\vspace{-0.5cm} \\begin{center} \\caption{Expected fluxes for carbon Miras in the LMC} \\begin{tabular}{llllllllllll} \\hline Period & Lum. & \\mdot & $V$ & $I$ & $J$ & $H$ & $K$ & $[12]$ & $[25]$ \\\\ (days)& (\\lsol) & (\\msolyr) & (mag) & & & & & (mJy) & (mJy) \\\\ \\hline 320 & 3000 & 1.7 $\\times$ 10$^{-8}$& 17.2 & 14.4 & 12.5 & 11.7 & 11.0 & 6.0 & 1.6 \\\\ 680 & 10000 & 8.4 $\\times$ 10$^{-6}$& 26.8 & 20.9 & 16.1 & 13.9 & 11.7 & 177 & 76 \\\\ \\hline \\end{tabular} \\end{center} \\vspace{-1.0cm} \\end{table} ", + "conclusions": "ISO provides the opportunity to observe AGB stars in the MCs with unprecedented sensitivity and at wavelengths that are inaccessible from the ground. It is in a way unfortunate that the majority of ISO time is spent on follow-up observations of IRAS detected stars, as we know that this sample is biased. Only one program uses the survey capability of CAM to observe field stars, but even that is restricted to 0.5 sq. degree in the LMC only. There is life after ISO: the near-infrared surveys DENIS and 2MASS will observe the whole of the MCs, and in this respect will provide the most unbiased view of the AGB population. In addition, the multi-wavelength photometric and spectroscopic capabilities of the VLT, the multi-object optical spectroscopic capability of de 2dF instrument on the AAT, and the MACHO/EROS projects that identify LPVs ensure that extra-galactic AGB research will be an interesting research topic in the future. \\vspace{-0.3cm} \\subsection*{Acknowledgements} I would like to thank the people who have feeded me with data and for useful discussions and without whom this review would not have been possible: Joris Blommaert, Maria Rosa Cioni, Eric Josselin, Jacco van Loon, Cecile Loup and Norman Trams. \\vspace{-0.3cm}" + }, + "9708/astro-ph9708132_arXiv.txt": { + "abstract": "{ We use multi-fiber spectroscopy of 12 poor groups of galaxies to address: (1) whether the groups are bound systems or chance projections of galaxies along the line-of-sight, (2) why the members of each group have not already merged to form a single galaxy, despite the groups' high galaxy densities, short crossing times, and likely environments for galaxy-galaxy mergers, and (3) how galaxies might evolve in these groups, where the collisional effects of the intra-group gas and the tidal influences of the global potential are weaker than in rich clusters. Each of the 12 groups has fewer than $\\sim$ five cataloged members in the literature. Our sample consists of 1002 galaxy velocities, 280 of which are group members. The groups have mean recessional velocities between 1600 and 7600 \\ks. Nine groups, including three Hickson compact groups, have the extended X-ray emission characteristic of an intra-group medium (Mulchaey \\& Zabludoff 1997 (Paper II)). We conclude the following: {\\it (a) The nine poor groups with diffuse X-ray emission are bound systems with at least $\\sim $20-50 group members to $M_B \\sim -14$ to $-16 + 5$log$_{10}$ \\lith}. The large number of group members, the significant early-type population (up to $\\sim 55\\%$ of the membership) and its concentration in the group center, and the correspondence of the central, giant elliptical with the optical and X-ray group centroids argue that the X-ray groups are not radial superpositions of unbound galaxies. The velocity dispersions of the X-ray groups range from 190 to 460 \\ks. We are unable to determine if the three non-X-ray groups, which have lower velocity dispersions ($< 130$ \\ks) and early-type fractions ($= 0$\\%), are also bound. {\\it (b) Galaxies in each X-ray-detected group have not all merged together, because a significant fraction of the group mass lies outside of the galaxies and in a common halo.} The velocity dispersion of the combined group sample is constant as a function of radius out to the virial radius of the system (typically $\\sim 0.5$\\lith\\inv\\ Mpc). The virial mass of each group ($\\sim 0.5$-$1\\times 10^{14}h^{-1} \\Mdot$) is large compared with the mass in the X-ray gas and in the galaxies ({\\it e.g.}, $\\sim 1 \\times 10^{12} h^{-5/2} \\Mdot$ and $\\sim 1 \\times 10^{13} h^{-1} \\ \\Mdot$, respectively, in NGC 533). These results imply that most of the group mass is in a common, extended halo. The small fraction ($\\sim 10$-$20\\%$) of group mass associated with individual galaxies suggests that the rate of galaxy-galaxy interactions is lower than for a galaxy-dominated system (Governato \\etal 1991; Bode \\etal 1993; Athanassoula 1997), allowing these groups to virialize before all of their galaxies merge and to survive for more than a few crossing times. {\\it (c) The position of the giant, brightest elliptical in each X-ray group is indistinguishable from the center of the group potential, as defined by the mean velocity and the projected spatial centroid of the group galaxies.} This result suggests that dominant cluster ellipticals, such as ``cD\" galaxies (Matthews, Morgan, \\& Schmidt 1965), may form via the merging of galaxies in the centers of poor group-like environments. Groups with a central, dominant elliptical may then fall into richer clusters (Merritt 1985). This scenario explains why ``cD\"s do not always lie in the spatial and kinematic center of rich clusters (Zabludoff \\etal 1990; Dunn 1991; Zabludoff \\etal 1993), but instead occupy the centers of subclusters in non-virialized clusters (Geller \\& Beers 1983; Bird 1994; Beers \\etal 1995). {\\it (d) The fraction of early-type galaxies in our poor groups varies significantly, ranging from that characteristic of the field ($\\simless 25$\\%) to that of rich clusters ($\\sim 55$\\%).} The high early type fractions are particularly surprising, because all of the groups in this sample have substantially lower velocity dispersions (a factor of $\\sim 2$-5) and galaxy number densities (a factor of $\\sim 5$-20) than are typical of rich clusters. Hence, the effects of disruptive mechanisms like galaxy harassment (Moore \\etal 1996) on the morphology of poor group galaxies are weaker than in cluster environments. In contrast, the kinematics of poor groups make them preferred sites for galaxy-galaxy mergers (Barnes 1985; Aarseth \\& Fall 1980; Merritt 1985), which may alter the morphologies and star formation histories of some group members. If galaxy-galaxy interactions are not responsible for the high early type fractions, it is possible that the effects of environment are relatively unimportant at the current epoch and that the similarity of the galaxy populations of rich clusters and some poor groups reflects conditions at the time of galaxy formation. {\\it (e) The fraction of early-type group members that have experienced star formation within the last $\\sim$ 2\\lith\\inv\\ Gyr is consistent with that in rich clusters with significant substructure ($\\sim 15\\%$; Caldwell \\& Rose 1997).} If some of the subclusters in these rich, complex clusters are groups that have recently fallen into the cluster environment, the similarity between the star formation histories of the early types in the subclusters and of those in our sample of field groups indicates that the cluster environment and associated mechanisms like ram pressure stripping (Gunn \\& Gott 1972) are not required to enhance and/or quench star formation in these particular galaxies. If the recent star formation is tied to the external environment of the galaxies and not to internal instabilities, it is more likely that galaxy-galaxy encounters have altered the star formation histories of some early type galaxies in groups and in subclusters. \\bigskip \\bigskip \\noindent{\\it Subject headings}: galaxies: clustering --- galaxies: distances and redshifts --- galaxies: elliptical and lenticular, cD --- galaxies: evolution --- galaxies: interactions --- cosmology: dark matter --- cosmology: large-scale structure of Universe } \\vfill\\eject ", + "introduction": "Most galaxies in the local universe, including our own Galaxy, belong to poor groups of galaxies. Despite the ubiquity of the group environment, we know little about the matter content of groups and the evolution of group galaxies outside of the Local Group. Because poor groups typically contain fewer than five bright ($\\simless M^*$) galaxies, studies to date have been hampered by small number statistics. Some of the critical, unanswered questions are (1) whether poor groups are in fact bound systems with significant populations of fainter members, (2) why many poor groups, with their high galaxy densities, short crossing times, and favorable environments for galaxy-galaxy mergers, survive long enough to be cataloged, and (3) how galaxies might evolve in an environment where the influences of the intra-group medium and the global potential are weak compared with those in rich clusters. The advent of multi-object spectroscopy now makes it possible to address these questions in unprecedented detail. In this paper, we present the first results from a fiber spectroscopic survey of 12 poor groups of galaxies. The issue of whether many poor groups, even those identified from redshift surveys, are bound systems instead of chance superpositions of galaxies along the line-of-sight has been a puzzle. The existence of one poor group, our Local Group, is unchallenged. In contrast, Ramella \\etal (1989) show that $\\sim 30\\%$ of groups of three or four galaxies in the CfA Redshift Survey (Huchra \\etal 1995) are probably unbound, geometric projections. One useful approach in finding bound systems is to identify those, like certain Hickson compact groups, with apparently interacting members (Rose 1977; Hickson 1982; de Oliveira \\& Hickson 1994). Yet these systems are also subject to projection effects ({\\it i.e.,} a pair of interacting galaxies with two interlopers; Mamon 1992) and, because they constitute only a small fraction of cataloged groups, may not be dynamically representative. Another strategy is to search for poor groups with diffuse X-ray emission (Mulchaey \\etal 1993; Ponman \\& Bertram 1993; Pildis \\etal 1995), in which the existence of a common gravitational potential is suggested by the intra-group gas ({\\it e.g.,} Ostriker \\etal 1995). ROSAT images reveal that at least $25\\%$ (22 of 85) of Hickson compact groups (Ponman \\etal 1996) have such an intra-group medium. However, the potentials of some poor groups may be too shallow for emission from hydrostatic gas to reach detectable levels (Mulchaey \\etal 1996a). It is even possible that the gas, like the galaxies in some cases, is merely a projection of unbound material in a filament alone the line-of-sight (Hernquist \\etal 1995). To determine whether a poor group is a bound system and to explore its kinematics in detail, we must spectroscopically identify more members. Furthermore, if fainter populations of galaxies do exist in these groups, then we will have the statistics necessary to quantify the mass associated with the galaxy, X-ray gas, and dark matter components and to better understand how galaxies may evolve in such environments. If some poor groups are bound systems, then another critical question is why they exist at all. Poor groups have higher galaxy densities than the field and lower velocity dispersions than cluster cores, making them favorable sites for galaxy-galaxy mergers (Barnes 1985). Galaxies are tidally interacting or merging in many Hickson compact groups (de Oliveira \\etal 1994; Longo \\etal 1994; Hunsberger \\etal 1996; Yun \\etal 1997). The likelihood of mergers and the short group crossing times ($\\simless 0.05$ of a Hubble time) suggest that most groups should have already merged into one object. Therefore, either bound groups are collapsing for the first time or only a small fraction of the group mass is tied to the galaxies, lowering the rate of galaxy-galaxy interactions relative to a galaxy-dominated system and allowing the group to survive many crossing times (Governato \\etal 1991; Bode 1993; Athanassoula 1997). To resolve this issue by measuring the underlying mass distribution of poor groups, we need to improve the statistics of group membership with an extensive spectroscopic survey. Once we know if a group is real and how much mass is associated with its galaxies, we can investigate the influences of group environment on galaxy evolution. For example, because ``cD\" galaxies in clusters lie in regions of high local density (Beers \\& Geller 1983; Zabludoff \\etal 1990; Beers \\etal 1995), but not always in the center of the global cluster potential, these galaxies may evolve first in poor group-like environments prior to the final collapse and virialization of the cluster as a whole (Merritt 1985). If so, then ``cD\"s are likely to form via galaxy-galaxy mergers in the center of a collapsing group, where the conditions for mergers are most favorable (Merritt 1984; Tremaine 1990). Over the lifetime of the group, dynamical friction or radial orbits may bring in more galaxies to merge with the ``cD.\" X-ray-detected poor groups, which almost always contain a giant ($\\simless \\Mstar - 1$) elliptical near the peak of the X-ray emission, are ideal laboratories for testing this picture of ``cD\" formation. If the group environment is the birthplace of ``cD\"s, then the giant elliptical will be coincident with the centroid of the projected spatial distribution of galaxies and will have little peculiar velocity with respect to the mean of the system. An extensive spectroscopic survey of group members will allow us to test this formation hypothesis directly in the poor group environment. The factors that might affect the evolution of galaxies in poor groups are different from those present in rich clusters. Some of the proposed cluster-based processes, such as ram pressure stripping (Gunn \\& Gott 1972) and galaxy harassment (Moore \\etal 1996), are less effective in group environments, where the number density of bright galaxies and the global velocity dispersion are small compared with clusters. The lower velocity dispersions of poor groups suggest instead that galaxy-galaxy interactions ({\\it e.g.}, close tidal encounters or mergers) are likely to dominate any environmentally-dependent evolution of galaxies in groups. If clusters evolve hierarchically by accreting poor groups of galaxies (subclusters), members of an infalling group have recently experienced the hot, dense cluster environment for the first time. Therefore, galaxies in poor groups in the field are a control sample for understanding the factors that influence the evolution of their counterparts in subclusters. For example, we can compare the morphologies and recent star formation histories of galaxies in the substructures of complex clusters like Coma (Caldwell \\& Rose 1997) with those of galaxies in poor field groups. Differences between the samples would argue that cluster environment is important in transforming galaxies at the present epoch. On the other hand, the lack of such differences would suggest, as the simplest explanation, either that star formation and morphology are influenced by mechanisms present in both field groups and subclusters, such as galaxy-galaxy encounters, or that the effects of environment on galaxies are insignificant compared with conditions at the time of galaxy formation. This paper is organized in four sections that address whether poor groups are bound systems, why some groups have not yet disappeared, and how the group environment may affect galaxy evolution. Section $\\S2$ describes the the sample of 12 poor groups, the fiber spectra, and the classification of group members as early or late type. We discuss our results in $\\S3$, including the membership and global velocity dispersion of each group, the composite group velocity dispersion profile, the location of the dominant elliptical relative to the kinematic and projected spatial centers of groups, the fraction of early type galaxies in groups, and the fraction of these early types that show evidence for recent or on-going star formation. Section 4 summarizes our conclusions. ", + "conclusions": "\\subsection{Group Membership} Are poor groups bound systems or chance superpositions of galaxies along the line-of-sight? Our spectroscopic survey samples the fields of groups whose membership previously totaled less than five bright ($\\simless \\Mstar$) galaxies. The detection of a significant population of fainter members would be a first step in demonstrating that these poor groups are not just geometric projections of unbound galaxies. We determine the galaxy membership of each group from a pessimistic, $3\\sigma$-sampling algorithm (Yahil \\& Vidal 1977). We use the statistical bi-weight estimators of location (mean velocity) and scale (velocity dispersion) to identify $3\\sigma$ outliers in the distribution of galaxy velocities within $\\pm 3000$ \\ks\\ of the center of the main peak (see Beers \\etal 1990 for a description of the bi-weight estimators). On each successive iteration, the $3\\sigma$ outliers are removed and the location and scale of the peak re-calculated. We halt the procedure prior to the removal of the last set of outliers. The resulting membership of each of the 12 groups is indicated by the shaded histograms in Figure 2, which shows the galaxy velocity distributions from 0 to 30000 \\ks. The width of the velocity bins is 250 \\ks, roughly $3\\times$ the typical external error. In total, there are 280 group members. The first nine groups are detected by the ROSAT PSPC, the last three are not (see Paper II). The histograms of the X-ray groups reveal large populations of group members down to absolute magnitudes of $M_B \\sim -14$ to $-16 + 5$log$_{10}$ \\lith. Because the membership algorithm can not be applied to the small number of galaxies in the peak of each non-X-ray group, we accept all the galaxies within contiguous bins as group members. As a result, the membership for the non-X-ray groups may be overestimated. We show the projected spatial distributions of the group members in Figure 3. The angular size of each plot is $1.62\\times 1.62$ degrees (the fiber spectrograph field is $1.5\\times 1.5$ degrees). Each tickmark corresponds to 5.7 arcmin. Digitized scans of Palomar Sky Survey or UK Schmidt plates from the STScI/DSS are not complete for every group. Hence, we mark the boundary of the unsampled regions with a dashed line: the westernmost fifth of the NGC 741 field, in the northernmost fifth of the NGC 5129 field, in the northernmost fifth of the HCG 90 field, and the northernmost tenth of the NGC 7582 field. The morphological types of galaxies with apparent magnitudes of $m_B \\sim 17$ or brighter are indicated by ``0\" for early and ``S\" for late. The filled circles mark the untyped group members. The scale bar below each group name is 0.3\\lith\\inv\\ Mpc. In the X-ray groups, as in rich clusters, the early types concentrate more in the group centers than do the late types. There are no early types in the three non-X-ray groups (see $\\S3.5$). The number counts of group members in the velocity histograms in Figure 2 are not directly comparable, because each group field is sampled over a different physical radius and to a different absolute magnitude. To compare the galaxy number densities of the groups, we subsample each system within a projected radius of 0.3\\lith\\inv\\ Mpc and to an absolute magnitude of $M_B \\sim -17 + 5$log$_{10}$ \\lith. Figure 4 shows the observed number counts of group members within these limits (shaded). To roughly compensate for incomplete sampling down to the magnitude limit (completeness is indicated by the fraction above each histogram bar), we assume that the fraction of all unobserved galaxies that are group members is the same as the fraction of all observed galaxies that are members. The white histogram shows these ``corrected\" group galaxy counts. Our rough calibration of the FOCAS magnitudes introduces more uncertainty into the ``corrected\" galaxy counts, so a small difference between the counts of two groups is not significant. Also note that the ``corrected\" counts are lower limits for two groups, NGC 7582 and NGC 5846, for which the $1.5\\times 1.5$ degree fiber field corresponds to less than 0.3\\lith\\inv\\ Mpc (0.21 and 0.24\\lith\\inv Mpc, respectively). Nevertheless, as this plot shows, non-X-ray-detected groups have lower galaxy densities than are typical of X-ray groups. For the same radial and magnitude cuts, the core of the Coma cluster (NED) has 83 galaxies (a lower limit because we make no correction for incompleteness). Thus, the galaxy density of the Coma core is $\\sim$5-$20\\times$ that of the poor group cores. By sampling to deeper magnitudes and to larger radii than past studies, we find that the physical extents of the poor groups in Figure 3, even those of HCG 42, HCG 62, and HCG 90, are larger than the typical values in the literature for Hickson compact groups (Hickson \\etal 1992). The number densities of galaxies within 0.3\\lith\\inv\\ Mpc and with $M_B \\simless -17 + 5$log$_{10}$ \\lith\\ in HCG 42, HCG 62, and HCG 90 (Figure 4) are a factor of $\\sim 70$-700 times lower than those inferred from the brightest four members (Hickson, Kindl, \\& Huchra 1988). These discrepancies result from the method of selecting Hickson groups, which are defined as concentrations of four or five bright galaxies within a projected radius of $\\simless 0.1$\\lith\\inv\\ Mpc (see also de Carvalho \\etal 1994). The galaxy number densities of the Hickson compact groups are not significantly different from those of the other poor groups in our sample. In general, the presence or absence of diffuse X-ray emission better differentiates between poor groups of high and low galaxy densities (and velocity dispersions ($\\S3.2$) and early type fractions ($\\S3.5$)). Could the X-ray groups be superpositions of three or four unbound, field galaxies and their fainter satellites? This explanation is unlikely, because the central, giant elliptical is not typical of field galaxies and the early type galaxies tend to concentrate in the group core. In addition, studies of the satellite populations of $\\simless \\Mstar$ field galaxies with the same fiber spectrograph setup (Zaritsky \\etal 1997), within similar projected radii from the primary galaxy, and to comparable absolute magnitude limits as our group sample find on average one satellite for each primary. The ratios of faint to giant galaxies in the X-ray groups exceed that expected from satellite statistics alone. The large populations of fainter galaxies and the central concentration of early types in the X-ray groups suggest that these groups are bound systems. The dynamical state of the non-X-ray groups is less obvious. In $\\S3.2$, 3.3, and 3.4, we use the global kinematic properties, the velocity dispersion profile, and the position of the central, giant elliptical relative to the kinematic and projected spatial group centroids as additional tests of whether the X-ray groups are bound and possibly virialized. \\subsection{Group Velocity Dispersions} \\subsubsection{Results} To date, the line-of-sight velocity dispersions (\\sigmar) of poor groups have been uncertain due to their determination from only a few galaxies. This uncertainty has translated into a large uncertainty in cosmologically important properties like the underlying mass distribution and the global baryon fraction. Because we increase the membership by a factor of 10 in many of our groups, we can, for the first time, determine poor group l-o-s velocity dispersions with sufficient precision ($\\simless 20\\%$) that the differences among them are statistically meaningful. The calculation of the mean velocity \\meanz\\ and \\sigmar\\ for each group is based on the bi-weight estimators of location and scale (Beers \\etal 1990) corrected for cosmological effects. These estimators are more robust than the standard mean and velocity dispersion ({\\it e.g.}, Danese \\etal 1980), but, for this sample, the standard velocity dispersion is always within the $68\\%$ confidence limits of \\sigmar. The velocity dispersions of the X-ray groups range over more than a factor of two, from 190 \\ks\\ in HCG 90 and 210 \\ks\\ in HCG 42 to 430 \\ks\\ in NGC 741 and 460 \\ks\\ in NGC 533. The results of $\\S3.3$ argue that the group velocity dispersion remains fairly constant with radius out to at least 0.5\\lith\\inv\\ Mpc. However, to confirm that the range of velocity dispersions above does not result from the different physical radius to which each group is sampled by the fixed angular size of spectrograph field, we determine the membership and the velocity dispersions within a 0.3\\lith\\inv\\ Mpc radius. The resulting \\sigmar 's are indistinguishable within the $68\\%$ confidence errors from those determined from all the data. Therefore, we use the \\meanz\\ and \\sigmar\\ determination from all the group members in subsequent analyses. Table 2 lists the group, the group optical projected centroid in J2000 (unweighted by luminosity), the total number of galaxies with measured redshifts in the fiber field $N_{tot}$, the number of group members $N_{grp}$, the bi-weight estimators of the mean heliocentric velocity \\meanz\\ and the line-of-sight velocity dispersion \\sigmar, and the physical radius of the group sampled by the fiber field $r_{samp}$. For the groups where $r_{samp} \\geq 0.67$ (the median pairwise radius $r_p$ for CfA Redshift Survey groups (Ramella \\etal 1989)), we tabulate $r_p$, the mean harmonic (virial) radius $r_h$, the virial mass $M_{vir}$, and the ratio of the crossing time to a Hubble time $t_c/t_H$. These last four kinematic quantities are calculated as in Ramella \\etal (1989). For the five groups with calculated masses and crossing times, $r_{samp}$ is $\\simgreat 20\\%$ larger than $r_p$, suggesting that the limited size of the fiber field does not artificially reduce our estimate of the group extent and bias the kinematic determinations. The median value of the harmonic radius of the four X-ray groups is 0.41\\lith\\inv\\ Mpc, consistent with the median $r_h = 0.5$\\lith\\inv\\ Mpc for groups in the CfA Redshift Survey (Ramella \\etal 1989). However, $r_h$ depends on how the distribution of galaxies is biased with respect to that of the dark matter in the system. To check that the harmonic radius is an accurate estimate of the virial radius, we calculate the radius $r_{500}$ that corresponds to an overdensity of 500, where N-body simulations show that the dynamical equilibrium hypothesis is satisfied (Evrard, Metzler, \\& Navarro 1996). For a poor group with an X-ray temperature typical of our X-ray groups (about 1 keV, Paper II), $ r_{500} = 1.24 \\sqrt{(kT / 10 \\ {\\rm keV}}) \\ h^{-1} \\ {\\rm Mpc} = 0.4 \\ h^{-1} \\ {\\rm Mpc} $. Therefore, we conclude that biasing between the galaxies and dark matter does not significantly affect the estimates of the kinematic quantities in Table 2. The velocity dispersions of poor groups in the literature are usually estimated from the four or five brightest members. With our deeper redshift samples, we can determine the accuracy of these past estimates. The velocity dispersions calculated from only the five brightest galaxies in each group $\\sigma_5$ tend to underestimate \\sigmar, because the tails of the galaxy velocity distribution are not well-sampled. In the worst case in our sample, the NGC 2563 group, $\\sigma_5$ underestimates \\sigmar\\ by a factor of three. In total, $\\sigma_5$ underestimates \\sigmar\\ by more than a factor of 1.5 for five of the nine X-ray groups. The difference between $\\sigma_5$ and \\sigmar\\ for the non-X-ray groups is small, because we calculate \\sigmar\\ itself from only $\\sim 5$ galaxies. \\subsubsection{Implications} The $\\sim 20$-50 group members, the central concentration of early type galaxies, and the short crossing times ($\\simless 0.05$ of a Hubble time) of the X-ray groups suggest that they are bound systems, not geometric superpositions of galaxies, and that the group cores are close to virialization or virialized. Because we do not detect diffuse X-ray emission or a significant fainter population in the three non-X-ray groups, we are unable to determine their dynamical state. The non-X-ray groups, which consist of one or two $L^*$ or brighter spirals with several fainter galaxies that may be satellites, are morphologically akin to the Local Group (although our samples are not sufficiently deep to ascertain whether any group has a dwarf spheroidal population like that of the Local Group; van den Bergh 1992). If the non-X-ray groups are dynamically similar to the Local Group, they are bound (see Zaritsky 1994). Our current data do not exclude this possibility --- the velocity dispersions of the non-X-ray groups are consistent with the upper limits on their X-ray luminosities (Paper II). If the non-X-ray groups are bound, but are collapsing for the first time like the Local Group (cf. Zaritsky 1994), the virial mass of systems like NGC 664 in Table 2 is uncertain by a factor of two (as the system is not yet virialized). If the non-X-ray groups are just chance superpositions, their kinematic quantities in Table 2 do not represent the properties of bound groups. \\subsection{Velocity Dispersion Profile} \\subsubsection{Results} The ratio of the mass associated with group galaxies to the mass associated with the common group halo determines the timescale of galaxy-galaxy interactions and thus the group's ability to survive for more than a few crossing times (Governato \\etal 1991; Bode 1993; Athanassoula 1997). The statistics of our sample make analyzing group kinematics on an individual basis difficult. To obtain an understanding of the underlying mass distribution in poor groups, we combine the velocity and projected spatial distributions of all of the X-ray group members. Figure 5 shows the velocity offset vs. the projected radial offset of 204 X-ray group members from the central, giant elliptical (the brightest group galaxy, hereafter BGG). The velocity difference is normalized with the internal velocity dispersion of the BGG ($\\sigma_{BGG}$). We include only the seven X-ray groups for which $\\sigma_{BGG}$ is known (HCG 42, HCG 90, NGC 2563, NGC 5129, NGC 5846, NGC 533, NGC 741; McElroy 1995; Trager 1997). We calculate the normalized velocity dispersion $\\delta$ (i.e., the {\\it rms} deviation in the normalized velocity offset) within each radial bin (dashed lines). The velocity dispersion of the combined group sample does not decrease significantly with radius from the central $\\sim 0.1$\\lith\\inv\\ Mpc to at least $\\sim 0.5$\\lith\\inv\\ Mpc, in contrast to the more than factor of two decrease in $\\delta$ that would be observed if the entire mass were concentrated within the first bin. The extended mass is either in the galaxies, in a common halo through which the galaxies move, or in both the galaxies and a diffuse halo. If all the mass were tied to the galaxies, most of the mass would be associated with the bright, central elliptical in those groups in which the BGG dominates the light. For groups with a few galaxies that have luminosities comparable to the BGG, $\\delta$ would be increased at large radii by subgroups consisting of a massive galaxy and the subgroup members that are orbiting it. If this picture were accurate, we would expect the $\\delta$ profiles of groups with several comparably bright galaxies to be shallower than those in which the BGG is dominant. However, the combined velocity dispersion profile of a subsample of two groups (HCG 42 and NGC 741) in which the BGG dominates the light ({\\it i.e.}, the BGG luminosity exceeds the combined luminosity of the other $L^*$ or brighter galaxies) is indistinguishable from that of the entire sample. Therefore, we conclude that most of the group mass lies in a smooth, extended dark halo. Is this halo associated with the group as a whole or with the BGG? If a galaxy forms inside a dark halo through dissipative collapse (cf. Blumenthal \\etal 1986), baryons concentrate in the center, deepening the gravitational well and increasing the velocity dispersion of the galaxy relative to that of the halo. In all the radial bins in Figure 5, $\\delta > 1$, indicating that the BGG is on average dynamically {\\it cooler} than the surrounding group. Thus, it is likely that the dark halo belongs to the group and not to the central, giant elliptical (note that the group velocity dispersion profile does not decrease with radius when the velocity and projected radial offsets are defined relative to the group mean velocity and projected spatial centroid, as in Figure 6a). We reach the same conclusion in Paper II by examining the X-ray images and spectra of the hot gas in these groups. It is possible that contamination from galaxies not bound to the group artifically increases the velocity dispersion of the combined group sample, especially at large radii from the BGG. We can conservatively estimate the degree of contamination by defining the 12 galaxies in the fifth (last) bin in Figure 5 with peculiar velocities $> 1.33\\sigma_{BGG}$ and $< -1.33\\sigma_{BGG}$ as interlopers. Because the interloper fraction should be constant with peculiar velocity, we predict that there are a total of about $12(3/2) = 18$ interlopers in the fifth bin. The ratio of the areas sampled by the fourth and fifth bins is 0.24. Therefore, only about four of the galaxies in the fourth bin and one or none of the galaxies in each the inner three bins are likely to be interlopers. Even excluding the four galaxies with the most extreme peculiar motions from the fourth bin only decreases the velocity dispersion in this bin by $15\\%$ to $\\delta = 1.1$. We conclude that the shape of the combined velocity dispersion profile within $0.5$\\lith\\inv\\ Mpc is not significantly affected by outliers. \\subsubsection{The Mass of Group Halos} \\renewcommand{\\thefootnote}{\\fnsymbol{footnote}} The virial masses of the poor X-ray groups are $M_{vir} \\sim 0.5$-$1 \\times 10^{14} h^{-1} \\Mdot$. What fraction of this mass is in the common group halo and what fraction is associated with group galaxies? We expect that early in a group's evolution, the individual galaxy halos are tidally-limited by the global potential of the group (cf. Peebles 1970; Gunn 1977). We must first estimate the tidal extent of the halos of X-ray group members before we can estimate their masses. Using Merritt's (1984) approach, we assume that the group's potential is in the form of an analytic King model (King 1972)\\footnote[2]{Although cosmological simulations predict a somewhat different form for the density profile (Navarro \\etal 1996), we use a King profile because (1) it is a good fit to the data (cf. Paper II) and (2) we wish to be consistent with past analyses of rich clusters.}.\\renewcommand{\\thefootnote}{\\arabic{footnote}} The maximum tidal radius of a galaxy near the core of a group is $ r_T \\approx {1 \\over 2} (r_c \\sigma_g / \\sigma_r)$, where we adopt $r_c \\sim 125$\\lith\\inv\\ kpc as the group's core radius (equivalent to the cluster value; Bahcall 1975) and $\\sigma_g$ as the galaxy's internal, line-of-sight velocity dispersion. For $L^*$ galaxies ($\\sigma_g \\approx 225$ \\ks; Tonry \\& Davis 1981) in groups with velocity dispersions of \\sigmar\\ $= 200$ and 450 \\ks, the maximum tidal radii are $\\sim 75$ and 35\\lith\\inv\\ kpc, respectively, which are larger than the luminous radii typical of such galaxies. These values of $r_T$ are less than the $> 150$\\lith\\inv\\ kpc dark halos of isolated disk galaxies (Zaritsky \\& White 1994), but greater than the predicted, tidally-limited extent of galaxies in the cores of rich clusters ($\\sim 15$\\lith\\inv\\ kpc, Merritt 1984). What is the fraction of the group mass associated with its galaxies? We estimate the tidally-limited mass of an $L^*$ galaxy from the virial theorem: $m \\approx {1 \\over 2} (r_T \\sigma_g^2/G)$ for a King model (1966). For $L^*$ galaxies in a \\sigmar\\ $= 200$ and 450 \\ks\\ group, we obtain $m^* \\sim 4\\times 10^{11}$ and $2\\times 10^{11} h^{-1} \\ \\Mdot$, respectively. To estimate the mass of the brighter, central elliptical, we assume that its internal velocity dispersion is $\\sigma_{BGG} \\sim 300$ \\ks\\ (typical of the central, giant elliptical in our X-ray groups) and that its radius is $\\sim 100$\\lith\\inv\\ kpc (which is between $r_T$ and $r_c$, a range consistent with the effective optical radii of ``cD\" galaxy envelopes (Schombert 1988), with the maximum radius to which the X-ray emitting gas in a BGG is detected (Paper II), and with the possibility that the BGG is less tidally-limited than other group members, because of its position in the group center ($\\S3.4$) where the tidal forces are symmetric and relatively weak (Merritt 1984)). We obtain $\\sim 10^{12} h^{-1} \\ \\Mdot$ as an estimate of the mass of the BGG's in the sample groups. For a group like NGC 741, the mass associated with the four $L^*$ or brighter galaxies within the virial radius $r_h = 400$\\lith\\inv\\ kpc is then $\\sim 10^{12} + 3(2 \\times 10^{11}) h^{-1} \\ \\Mdot \\approx 2 \\times 10^{12} h^{-1} \\ \\Mdot$. We can also estimate the fraction of mass tied to fainter group galaxies. Assuming that the luminosity spectrum of group galaxies is a Schechter function (1976) with $\\gamma = 1.25$ and a lower cut-off at $0.05L^*$ ($L^* \\sim 1 \\times 10^{10} \\ L_{\\odot}$; Kirshner \\etal 1983), we use the empirical luminosity-velocity dispersion relationship ($L \\propto \\sigma_g^4$; Faber \\& Jackson 1976), the equation for tidal radius ($r_T \\propto \\sigma_g$), and the virial theorem ($m \\propto \\sigma_g^2 r_T$) to convert the luminosity function into a tidally-limited mass function (see Merritt 1984 for a similar discussion). We obtain $$dn = {4 \\over 3} e^{-{\\left(m \\over m^*\\right)^{4/3}}} {\\left(m \\over m^*\\right)}^{\\left(1 - 4\\gamma\\right)/3} d\\left({m \\over m^*}\\right)$$ for galaxies with tidally-limited masses greater than $0.1m^*$, where we assume that the halos of all galaxies are tidally-limited during a close passage to the group core and that only non-luminous matter is removed. The mass of the $L^*$ or brighter members is $\\sim 20\\%$ of the total galaxy mass. Therefore, the total mass in galaxies in NGC 741 is $\\sim 1 \\times 10^{13} h^{-1} \\ \\Mdot$, in contrast to the the virial mass of $\\sim 1 \\times 10^{14} h^{-1} \\ \\Mdot$. The fraction of mass in the galaxies within the virial radius is then roughly $10\\%$. In lower velocity dispersion groups such as NGC 4325, there are fewer bright members, but the galaxy halos are less tidally-limited and the virial mass is smaller. Hence, galaxies comprise a greater fraction of the total mass within the virial radius, $\\sim 20\\%$. The dominant source of uncertainty in the estimation of the fraction of group mass associated with individual galaxies is the galaxy halo mass. The degree to which group galaxy halos are tidally-limited depends on such key unknowns as the details of galaxy halo formation, the initial orbits of the group members, and the evolution of the group's mass density. It is also possible that additional mass could be collisionally-stripped from galaxies later in the evolution of the group (see $\\S3.5$). The exploration of these effects awaits improved simulations. \\subsubsection{Implications} The small fraction of the group mass associated with its galaxies will increase the timescale for galaxy-galaxy interactions, such as mergers, close tidal encounters, and dynamical friction, by decreasing the cross-sections of the galaxies (Governato \\etal 1991; Bode 1993; Athanassoula 1997). This result argues that poor groups survive longer than predicted by models in which all the mass is tied to individual galaxies and may explain why so many poor groups are observed in lieu of single merger remnants. \\subsection{The Giant Elliptical vs. the Group Centroid} \\subsubsection{Results} The poor groups with diffuse X-ray emission in our sample have a giant ($\\simless \\Mstar - 1$) elliptical that lies within $\\sim 5$-10\\lith\\inv\\ kpc of the X-ray peak (Paper II). The coincidence of the BGG position and the X-ray peak suggests that BGG's lie in the center of the group potential. However, it is possible that X-ray emission in the core is dominated by light from the BGG, hiding an offset between the galaxy and the center of the potential as defined by peak of the intra-group medium. Therefore, the optical data are required to determine whether the BGG is at rest with respect to the group as a whole. In this section, we test this hypothesis in two ways, asking (1) whether BGG's are more concentrated in the core than other group members and (2) whether the radial velocity and projected position of the BGG's are consistent with the kinematic and projected spatial centroids of the groups. How are the BGG's distributed in projected radius and peculiar velocity compared with other group galaxies? Figure 6a shows the kinematic ($y$) and projected spatial ($x$) offsets of the BGG (filled squares) and other group members (filled circles) from the group centroid for the nine X-ray groups. Neither the projected spatial centroid ($\\langle r \\rangle$) nor the mean velocity ($\\langle \\upsilon \\rangle$) of the group are weighted by galaxy luminosity, so the velocity and projected position of the BGG do not bias the centroid calculations. The velocity offset of each galaxy from the mean velocity of the group is divided by the group velocity dispersion \\sigmar\\ to compensate for differences in the size of the group global potentials. We define a statistic $R^2 = (x/\\delta_x)^2 + (|y|/\\delta_{|y|})^2$, where $\\delta_x$ and $\\delta_{|y|}$ are the $rms$ deviations in $x$ and $|y|$ for the entire sample. Thus, a galaxy that has a large peculiar motion and/or that lies outside the projected group core will have a larger $R$ value than a galaxy at rest in the center of the group potential. Figure 6b shows the distributions of $R$ for all of the group members (solid) and for the subset of BGGs (shaded). A Student's t-test gives $<3 \\times 10^{-5}$ as the probability that the means of the non-BGG and BGG distributions are consistent. Therefore, the BGG's are significantly more concentrated in the core, suggesting that BGG's occupy the center of the potential on orbits different from those of other group members. Are the peculiar velocities and projected positions of the BGG's consistent with the kinematic and projected spatial centers of the groups? The $y$ errorbars in Figure 6a represent the $68\\%$ confidence limits on $y$ ($\\epsilon_y$) obtained from adding the errors in the group mean velocity and the BGG velocity in quadrature. To estimate the $68\\%$ error in $x$ for a BGG in a group with $N_{grp}$ members, we use a statistical jackknife test in which samples of $N_{grp}$ galaxies are drawn from HCG 62, the group with the most members. For each BGG, we adopt the $rms$ deviation in the distribution of $x$ ($\\epsilon_x$) for these samples as the $x$ error. (For the BGG in HCG 62, we use the smallest $x$ error calculated among the other groups). The heavy line in Figure 6b shows the distribution of $R$ obtained by assuming that all of the BGG's lie in the centers of their groups. For each BGG, we model the distributions of errors about the center in $|y|$ and in $x$ as Gaussian with $rms$ deviations equal to $\\epsilon_y / \\delta_{|y|}$ and $\\epsilon_x / \\delta_x$, respectively. The heavy line is the distribution in $R$ that results from 1000 random draws from the model distributions of $x$ and $|y|$ for each BGG. A t-test fails to differentiate between the means of the model $R$ distribution and that of the BGGs (shaded) at better than the $95\\%$ confidence level. We conclude that, to within the errors, the BGG's occupy the center of the group potential. \\subsubsection{Timescales for BGG Evolution} The presence of a giant elliptical at rest in the center of each X-ray group is consistent with the picture in which BGG's form from galaxy mergers early in the group's evolution (Merritt 1985; Bode 1994). BGG formation may occur during the initial collapse of the group, before galaxy halos are tidally-truncated and merger rates decline (Merritt 1985). The absolute magnitudes of the BGG's place them in the class of bright ellipticals whose morphologies and kinematics are consistent with merger evolution ($M_B \\simless -20 +5$log$_{10}$ \\lith; Kormendy \\& Bender 1996; Merritt \\& Tremblay 1996). Such a galaxy may be forming in HCG 90. The total luminosity of the several merging galaxies in HCG 90's core is comparable to the absolute magnitudes of the BGG's in the other X-ray groups. Do BGG's continue to experience mergers after their formation? The BGG's in our sample show no signatures of later mergers with massive galaxies; their spectra do not indicate recent star formation, their morphologies are not obviously disturbed, and their position in the center of the X-ray emission argues against recent disruptions of the surrounding gas by mergers. For example, one BGG (NGC 5846) is classified as an intermediate age elliptical (the last star formation event occurred more than $\\sim 6$-7\\lith\\inv\\ Gyr ago; Trager 1997). However, NGC 2300, the central, giant elliptical in the first poor X-ray group discovered (Mulchaey \\etal 1993), has shells (Schweizer \\& Seitzer 1992; Forbes \\& Thomson 1992) and a spectrum that includes a young stellar population ($\\simless 2$\\lith\\inv\\ Gyr old; Trager 1997). Thus, there is some indirect evidence that at least one BGG has experienced a late merger. We can estimate whether the central elliptical is likely to grow via subsequent mergers with other group members. One possible mechanism is the accumulation of slow-moving, massive galaxies in the group core by dynamical friction. To determine the effects of dynamical friction, we roughly estimate the largest radius from which an $L^*$ galaxy could fall into the center of a group in a Hubble time. We use the dynamical friction timescale (Binney \\& Tremaine 1987; their eq. 7-26) and the Coulomb logarithm definition (7-13b) for an $L^*$ galaxy on an initially circular orbit to solve for the maximum radius. We use the tidally-limited masses of an $L^*$ galaxy in a \\sigmar\\ $= 450$ \\ks\\ and 200 \\ks\\ group from $\\S{3.3}$. For a galaxy moving with circular velocity $\\upsilon_c = \\sqrt 2$\\sigmar, the maximum radius ranges from $\\sim 250$\\lith\\inv\\ kpc for the high velocity dispersion group to 400\\lith\\inv\\ kpc for the low \\sigmar\\ group. This timescale estimate is consistent with simulations (Merritt 1984), which predict that dynamical friction can cause 1-$2L^*$ worth of merger candidates to accumulate in the core of a \\sigmar\\ $= 500$ \\ks\\ group within a Hubble time. If these candidates do merge with the BGG, some enhancement of the BGG's luminosity is possible, and the luminosity segregation resulting from dynamical friction might then be hidden by the mergers. In contrast, the same equations and simulations predict that an $L^*$ galaxy in a rich cluster ($\\sigma_r \\sim 1000$ \\ks) requires an improbable $\\sim 40$\\lith\\inv\\ Gyr to fall to the center from an initial radius of 250\\lith\\inv\\ kpc. \\subsubsection{Implications} The existence of a central, giant elliptical in the X-ray groups has implications for the formation of ``cD\" cluster galaxies (Matthews, Morgan, \\& Schmidt 1965). The method by which ``cD\"s evolve is unknown, but there is indirect evidence that these galaxies form outside the cores of rich clusters. For example, ``cD\"s are not always at rest with respect to the cluster potential, but lie instead in local overdensities (Beers \\& Geller 1983; Beers \\etal 1995). In addition, the presence of a ``cD\" does not depend on the global kinematic properties of the cluster ({\\it i.e.,} clusters that contain a ``cD\" have similar velocity dispersions to those that do not, Zabludoff \\etal 1990). Our observations suggest that ``cD\"s form from galaxy-galaxy mergers in poor groups. Because poor group velocity dispersions are comparable to the internal velocity dispersions of their galaxies, the probability that a galaxy-galaxy collision will lead to a merger is higher in groups than in clusters (see $\\S3.5$). Poor groups with central, massive ellipticals may later merge to form clusters or fall into existing clusters (Merritt 1985; Beers \\etal 1995). In this picture, the BGG would be displaced for a time from the center of the cluster potential. We note that the NGC 2563 X-ray group is one of the groups in Cancer (Bothun \\etal 1983), an association of groups likely to evolve into a rich cluster after its virialization. In future work, we will use deep optical surface photometry to determine whether the BGG's in the X-ray groups have the extended halos of ``cD\" galaxies in rich clusters. \\subsection{Group Early-Type Fractions} \\subsubsection{Results} The effects of environment on galaxy evolution are poorly understood. The study of poor groups provides a critical link between the evolution of isolated galaxies in the field and of galaxies subject to the hot, dense environments of clusters. Despite the usefulness of group galaxies as a control sample, their properties, especially at faint magnitudes, are not well-known. One possible test of environmental influences is to compare the morphologies of group members with those of field and cluster galaxies. Previous work suggests that the fraction of early type (E and S0) galaxies in X-ray detected groups varies widely and that some X-ray groups have no late types among their brightest members (Ebeling \\etal 1994; Pildis \\etal 1995; Henry \\etal 1995; Mulchaey \\etal 1996b). However, these studies typically include only the four or five brightest galaxies, which biases samples toward ellipticals, and target only the central $\\simless 0.3$\\lith\\inv\\ Mpc, where early types concentrate ({\\it e.g.}, Figure 3). Therefore, to compare the morphologies of group and cluster members, we must sample both environments to similar physical radii and absolute magnitude limits. We plot the distribution of the early type (E, S0) fractions $f$ of the 12 sample groups in Figure 7. In the X-ray groups, $f$ ranges widely from $\\sim 55\\%$ (HCG 62, NGC 741, and NGC 533) to $\\sim 25\\%$ ({\\it e.g.}, NGC 2563). The lower value of $f$ for the X-ray groups is characteristic of the field ($\\sim 30\\%$; Oemler 1992). We find no early types among the 6-8 galaxies in each of the three non-X-ray groups (shaded). We can test whether the groups with the highest early type fractions, such as NGC 533 (14 of 25) and NGC 741 (10 of 18), are statistically different from the field. The likelihood of at least $k$ successes in $n$ Bernoulli trials with success probability $f_F$ is $\\sum_j {n! \\over {j!(n-j)!}} {f_F}^j (1 - f_F)^{n-j}$ for all $j$ such that $n \\geq j \\geq k$. The fraction of field galaxies that are E's and S0's ($f_F$) is $\\sim 30\\%$ (Oemler 1992). If we assume that the E and S0 fraction in NGC 533 is actually consistent with $f_F$, the probability of finding 14 or more early types in a sample of 25 group members is very small, $6 \\times 10^{-3}$. The probability that the early type fraction in NGC 741 is consistent with the field is also small, $2 \\times 10^{-2}$. When these probabilities are considered together, the hypothesis that the early type fractions of both groups are consistent with the field is rejected at the $\\sim 4 \\sigma$ level. NGC 533, NGC 741, and HCG 62 are sampled to physical radii and absolute magnitude limits similar to NGC 2563, an X-ray group in which $f$ is consistent with the field, and to NGC 644, a non-X-ray group in which $f$ is much lower than the field ($0\\%$). We conclude that the spread in the distribution of $f$ in Figure 7 is statistically significant. How do the early type fractions of $\\sim 55\\%$ in NGC 533, NGC 741, and HCG 62 (16 of 30) compare with values typical of rich clusters? These three groups are sampled to physical radii of $\\sim 0.6$-0.8\\lith\\inv\\ Mpc, and their members are typed to absolute magnitude limits of $M_B \\sim -16$ to $-17 + 5$log$_{10}$ \\lith. Within this range of radii, the early type fractions of rich clusters are $\\sim 0.55$-0.65 (Whitmore \\etal 1993). Because the absolute magnitude of the Whitmore \\etal sample ($M_B \\sim -18 + 5$log$_{10}$ \\lith) is brighter than ours, the cluster early type fractions are relatively biased toward ellipticals. Cutting the group data at the brighter $M_B$ limit worsens the statistics, but increases $f$. For example, six of the seven galaxies in NGC 533 brighter than the Whitmore \\etal limit are E's or S0's. Likewise, five of the remaining six members of NGC 741 are early types. Therefore, it is conservative to conclude that the early type fractions in some poor groups are comparable to those in rich clusters. \\subsubsection{The Early Type Fraction-Velocity Dispersion Relation} The inset in Figure 7 shows the correlation between early type fraction and group velocity dispersion for the 12 sample groups. The solid line is an unweighted fit to the data, the dashed line is a fit weighted by the velocity dispersion errors. In both cases, the correlation between morphology and velocity dispersion is significant at the $>0.999$ level. The relation is surprisingly robust, given that the groups are sampled over a range of physical radii ($\\sim 0.2$-1.0\\lith\\inv\\ Mpc) and absolute magnitude limits ($M_B \\sim -14$ to $-17 + 5$log$_{10}$ \\lith). It is possible that the relation is artificially strengthened by the three non-X-ray groups, which may not be bound. We note, however, that the Local Group would similarly anchor the tail of the correlation. The form of the relation cannot be the same for rich clusters --- our fit to the group data predicts that the early type fraction in a $\\sigma_r \\sim 1000$ \\ks\\ cluster is an unphysical $f = 124\\%$. Therefore, the relation must turn up between the poor group and rich cluster regimes. The group $f - \\sigma_r$ relation implies either that galaxy morphology is set by the local potential size at the time of galaxy formation (Hickson, Huchra, \\& Kindl 1988) and/or that \\sigmar\\ and $f$ increase as a group evolves (Diaferio \\etal 1995). \\subsubsection{Galaxy-Galaxy Collision Timescales} If environment does alter the morphologies of group galaxies after the formation of the group, which environmental influences are most important? Proposed mechanisms to disturb the distribution of stars in cluster galaxies include the tidal shocking of stellar disks by the global potential (B. Moore 1997, priv. comm.), galaxy-galaxy harassment (Moore \\etal 1996), collisional stripping ({\\it e.g.}, Richstone 1975), and mergers ({\\it e.g.}, Barnes 1989; Weil \\& Hernquist 1996). Strong tidal shocks could strip the stellar disks from Sc or Sd galaxies passing through the strong gravitational field of the group core, producing remnants consistent with the mass profiles of dwarf elliptical galaxies (B. Moore 1997, priv. comm.); however, only four of the 72 galaxies in our early type sample are as faint as dE/dSph's ({\\it i.e.,} $M_B \\simgreat -16 + 5$log$_{10}$ \\lith), so we do not consider this mechanism in interpreting the group data. Galaxy harassment is not likely to operate effectively in groups, where the number density of bright galaxies and the global velocity dispersion are much smaller than in clusters (B. Moore 1997, priv. comm.). In the following discussion, we calculate the likelihood of galaxy-galaxy encounters (collisions, mergers) by adapting simple models applied in the past to rich clusters (cf. Richstone \\& Malumuth 1983; Merritt 1984). Galaxy-galaxy interactions are known to significantly affect galaxy morphology ({\\it e.g.}, Schweizer 1986, Hibbard \\etal 1994). Observations of poor compact groups suggest that close tidal interactions and accretion events play a role in the evolution of some group galaxies. For example, the isophotes of compact group ellipticals are typically more irregular than those of ellipticals in cluster environments (Zepf \\& Whitmore 1993), and the fraction of group members with disturbed morphologies and peculiar kinematics is larger than for field galaxies (Rubin \\etal 1991; de Oliveira \\& Hickson 1994). If poor groups are longer-lived than previously supposed, and merged galaxies have had more time to relax, the fraction of merger remnants in groups may be even higher than indicated by short-lived merger signatures. What is the timescale for galaxy-galaxy encounters in our poor groups? We now examine the timescales in four cases: (a) an $L^*$ galaxy in a \\sigmar\\ = 450 \\ks\\ group, (b) a galaxy of typical luminosity $L_{typ}$ in a \\sigmar\\ = 450 \\ks\\ group, (c) an $L^*$ galaxy in a \\sigmar\\ = 200 \\ks\\ group, and (d) an $L_{typ}$ galaxy in a \\sigmar\\ = 200 \\ks\\ group. {\\bf Definition of Collision Timescale.} Following Richstone \\& Malumuth (1983), we define the collision time as $T_{col} = [{\\overline n} \\pi {r_T}^{2} \\sqrt{2} \\sigma_r]^{-1}$, where ${\\overline n}$ is the mean number density of galaxies within the half-mass radius, ${r_T}$ is the tidal radius of the galaxy, and $\\sqrt{2} \\sigma_r$ is the typical encounter velocity in an isotropic system. We assume that the group mass distribution inside the virial radius is described by a King model (1966) with a core radius of $r_c \\sim 125$\\lith\\inv\\ kpc and with a shape defined by the ratio of the group tidal radius to $r_c$, $R_T/r_c = 20.2$ (see Bahcall 1977 and Dressler 1978 for discussions of the shapes and sizes of clusters). The half-mass radius of this model is $2.83 r_c = 354$\\lith\\inv\\ kpc, comparable to the typical group virial radius. Therefore, we calculate ${\\overline n}$ using the mass within the virial radius. The mean mass of a group member for the tidally-limited mass function defined in $\\S3.3$ is ${\\overline m} = 0.35m^*$, where $m^*$ is the mass of an $L^*$ galaxy. {\\bf (a) $L^*$ Galaxy in a \\sigmar\\ = 450 \\ks\\ Group.} To calculate $T_{col}$ for an $L^*$ galaxy in a \\sigmar\\ = 450 \\ks\\ group, we use the NGC 533 group as an example. In $\\S3.3$, we found $r_T \\sim 35$\\lith\\inv\\ kpc for an $L^*$ galaxy in this group. What is ${\\overline n}$? The total mass of NGC 533 within $r_h = 400$\\lith\\inv\\ kpc is $M_{vir} \\sim 1 \\times 10^{14} h^{-1} \\ \\Mdot$. Thus, if we assume that NGC 533 is spherical and that $f_{halo} \\sim 90\\%$ is the fraction of the total mass not tied to the galaxies, then the average galaxy number density within $r_h$ is ${\\overline n} = (3 M_{vir} / 4 \\pi {r_h}^3) (1 - f_{halo}) / {\\overline m} \\sim 529 h^3$ Mpc$^{-3}$. This calculation shows that we would overestimate ${\\overline n}$ and thus underestimate the collision timescale by a factor of 10 by assuming that all of the group mass is associated with the galaxies. Note also that $f_{halo}$ is assumed to be constant over time, as would be the case if the tidal limitation of the galaxies occurred early in the group's evolution. If $f_{halo}$ increases in time, then the collision timescales below are somewhat overestimated. With these values of ${\\overline n}$ and $r_T$, we obtain $T_{col} \\sim 0.07\\times$ a Hubble time --- an $L^*$ galaxy experiences a collision in every $\\sim$ four crossing times. In comparison, a rich cluster like Coma has at least $\\sim 5 \\times$ the galaxy number density and more than twice the velocity dispersion of NGC 533. If cluster galaxies have halos truncated to about half the extent of those in NGC 533 (Merritt 1984), then the collision time for an $L^*$ galaxy in a rich cluster is $\\sim 0.04$ of a Hubble time. Therefore, over the same period of time, an $L^*$ galaxy in a rich cluster is nearly twice as likely to experience a collision as an $L^*$ galaxy in NGC 533, the poor group with the shortest estimated $T_{col}$ in the sample. {\\bf (b) $L_{typ}$ Galaxy in a \\sigmar\\ = 450 \\ks\\ Group.} To understand more about how the collision history of a galaxy depends on its luminosity and on the velocity dispersion of the group, we estimate the collision time for a typically bright galaxy in NGC 533. The value of ${\\overline n}$ is the same as in case (a), but what is $r_T$? The luminosity of a typically bright galaxy is $L_{typ} \\approx 0.15L^*$ for the Schechter luminosity function assumed in $\\S3.3$ (Richstone \\& Malumuth 1983). $L_{typ}$ is two magnitudes fainter than $L^*$ and is also the approximate limit at which the completeness of our group samples is at least $\\sim 60\\%$ ($\\S3.1$). We use the Faber-Jackson (1976) relation and the internal velocity dispersion of an $L^*$ galaxy ($\\approx 225$ \\ks; Tonry \\& Davis 1981) to estimate $\\sigma_g$ for an $L_{typ}$ galaxy ($\\sim 140$ \\ks). The tidal radius of an $L_{typ}$ galaxy is then $\\sim 20$\\lith\\inv\\ kpc, and the collision time is about three times longer than for an $L^*$ galaxy in NGC 533. {\\bf (c) $L^*$ Galaxy in a \\sigmar\\ = 200 \\ks\\ Group.} An $L^*$ galaxy in a lower velocity dispersion group will collide less frequently than will its counterpart in NGC 533. The tidal radius of this galaxy is $\\sim 75$\\lith\\inv\\ kpc ($\\S3.3$). What is ${\\overline n}$? For a \\sigmar\\ = 200 \\ks\\ group, we estimate the virial mass within 400\\lith\\inv\\ kpc by assuming that the virial radii are the same and scaling NGC 533's $M_{vir}$ by the square of the ratio of the velocity dispersions of the two groups. Because (1) the virial mass is smaller, (2) the average galaxy mass is larger, and (3) $f_{halo} \\sim 80\\%$ is smaller in the \\sigmar\\ = 200 \\ks\\ group, ${\\overline n}$ is a factor of $\\sim$five times smaller than in NGC 533. As a consequence, the collision timescale for an $L^*$ galaxy in a \\sigmar\\ = 200 \\ks\\ group is $\\sim 3\\times$ longer than for an $L^*$ galaxy in NGC 533. {\\bf (d) $L_{typ}$ Galaxy in a \\sigmar\\ = 200 \\ks\\ Group.} The tidal radius of a typically bright galaxy in a \\sigmar\\ = 200 \\ks\\ group is $\\sim 45$\\lith\\inv\\ kpc, where we assume that the galaxy's velocity dispersion is the same as in case (b). The value of ${\\overline n}$ is the same as in case (c). Therefore, $T_{col}$ is $\\sim 8\\times$ longer in this case than for an $L^*$ galaxy in NGC 533. {\\bf Summary.} Collision timescales are shortest for bright cluster members, whose cross-sections are large because of the high relative velocities and galaxy number density. $T_{col}$ is increasingly longer for an $L^*$ galaxy in a \\sigmar\\ = 450 \\ks\\ group, an $L^*$ galaxy in a \\sigmar\\ = 200 \\ks\\ group, an $L_{typ}$ galaxy in a \\sigmar\\ = 450 \\ks\\ group, and an $L_{typ}$ galaxy in the lower \\sigmar\\ group. However, as we will see, the merger timescales do not follow this hierarchy, because they depend not only on the collision timescale, but also on the fraction of collisions that lead to mergers. \\subsubsection{Galaxy-Galaxy Merger Timescales} To roughly calculate the merger timescale, we assume that galaxies in collision merge if the ratio of the collision speed to their internal velocity dispersion ($\\sigma_g$) is $\\leq 3$ (Tremaine 1980; Richstone \\& Malumuth 1983). In other words, we determine the fraction of collisions $\\zeta$ that are slower than $3\\sigma_g$. As discussed in Richstone \\& Malumuth (1983), the fraction of relative velocities in one direction varies as $e^{-\\upsilon/4{\\sigma_r}^2}$. The merger fraction is then $\\zeta = \\int_{0}^{3\\sigma_g} e^{-\\upsilon^2/4{\\sigma_r}^2}\\upsilon^2 d\\upsilon / \\int_0^{\\infty} e^{-\\upsilon^2/4{\\sigma_r}^2}\\upsilon^2 d\\upsilon$. The merger timescale is thus $T_{merg} = T_{col} \\zeta^{-1}$. For an $L^*$ galaxy in NGC 533, $T_{merg} \\sim 4T_{col} \\sim 0.3t_H$. In contrast, the merger timescale in a rich cluster is $\\sim 100\\times$ its collision time (Richstone \\& Malumuth 1983), or about four Hubble times. Thus, $T_{merg}$ for an $L^*$ galaxy is $\\sim 13\\times$ shorter in NGC 533 than in rich clusters. For an $L_{typ}$ galaxy in NGC 533, the ratio of the velocity dispersion of the group to that of the galaxy is larger. Therefore, the fraction of collisions that lead to mergers is smaller, and $T_{merg}$ is $11\\times$ longer, than for an $L^*$ galaxy. The efficiency of mergers is higher in a \\sigmar\\ = 200 \\ks\\ group, because the internal velocity dispersion of an $L^*$ and an $L_{typ}$ galaxy is comparable to \\sigmar. This effect is somewhat offset by fewer collisions in the lower velocity dispersion group. Therefore, the merger timescale for an $L^*$ and an $L_{typ}$ galaxy is $0.7\\times$ and $4\\times$ that for an $L^*$ galaxy in NGC 533, respectively. In summary, the merger timescales are increasingly longer for an $L^*$ galaxy in a \\sigmar\\ = 200 \\ks\\ group, an $L^*$ galaxy in a \\sigmar\\ = 450 \\ks\\ group, an $L_{typ}$ galaxy in the lower \\sigmar\\ group, an $L_{typ}$ galaxy in the higher \\sigmar\\ group, and an $L^*$ galaxy in a rich cluster. \\subsubsection{Collisional Stripping Timescale} Even if two colliding galaxies do not merge, the encounter may disrupt their morphologies. The timescale over which mass is removed from a colliding galaxy is shorter in poor groups than in clusters, because the slower relative velocities of the colliding galaxies makes stripping more effective. Numerical experiments (Richstone 1975; Dekel, Lecar, \\& Shaham 1980; Richstone \\& Malumuth 1983) suggest that only 1-10\\% of a galaxy's mass is lost during a collision at velocities typical of galaxies in rich cluster. If we employ the upper bound on the fraction of mass lost, then $T_{strip} \\sim 10 T_{col}$ is a very rough estimate of the collisional stripping timescale in our highest velocity dispersion groups. Therefore, a collision that is effective at stripping mass from an $L^*$ galaxy in NGC 533 might occur within a Hubble time. More sophisticated simulations are required before this estimate can be improved. \\subsubsection{Implications} What do the merger timescales imply for galaxy evolution in poor groups? If \\sigmar\\ = 450 \\ks\\ groups like NGC 533 form from the mergers of lower \\sigmar\\ groups, galaxy-galaxy mergers occur less frequently as the group evolves. In this case, we would expect to observe more on-going galaxy mergers in lower \\sigmar\\ groups and more merger remnants in higher \\sigmar\\ groups. In a sample of compact groups with velocity dispersions of order 200 \\ks, a high fraction of galaxies are interacting ($\\sim 43\\%$; de Oliveira \\& Hickson 1994). The early type fractions in our $\\sigma_r \\simgreat 400$ \\ks\\ groups are as high as in rich clusters. If some early type galaxies are evolved merger remnants, this highly-simplified model is qualitatively consistent with the observed relationship between early type fraction and velocity dispersion, {\\it i.e.,} the galaxy populations of higher velocity dispersion groups are more evolved on average. At some point in the group's evolution, perhaps at a velocity dispersion near 400 \\ks, any morphological evolution resulting from galaxy mergers ceases, and the fraction of merger remnants in poor groups and rich clusters is comparable. The implicit upturn in our $f - \\sigma_r$ relation suggests such a saturation point. Alternatively, the similarity of some group and cluster early type fractions, and the steepening of the $f - \\sigma_r$ relationship at high \\sigmar, might arise from conditions at the time of galaxy formation. For example, it is possible that poor groups such as NGC 533 and rich clusters like Coma begin as similar mass density perturbations with correspondingly similar galaxy populations. In this simple picture, NGC 533 does not develop a cluster-size potential, because its field lacks the surrounding groups that Coma later accretes. In summary, the cluster-like fraction of early type galaxies in NGC 533, NGC 741, and HCG 62 suggests that the cluster environment is not required to produce copious quantities of E and S0 galaxies. The simplest explanation is either that fluctuations in the initial conditions permitted early types to form in these comparatively low velocity dispersion, low galaxy density environments or that the galaxies' subsequent evolution was the product of a mechanism, such as galaxy-galaxy interactions, common to both groups in the field and groups that become subclusters. Although additional environmental mechanisms may affect the evolution of cluster galaxies, such cluster-specific processes are not required to explain the current data. A cluster that evolves hierarchically from subclusters with the properties of NGC 533, NGC 741, and HCG 62 will have, at least initially, a similar galaxy population. As we demonstrate below, another test of the importance of cluster environment on galaxy evolution at the current epoch is to compare the recent star formation histories of galaxies in poor groups and in subclusters. \\vfill\\eject \\subsection{Early-Type Group Galaxies with Young Stellar Populations} \\subsubsection{Results} The star formation histories of galaxies in poor groups provides additional insight into the environmental factors that may influence the evolution of galaxies. One approach is to examine the spectra of the early types for evidence of on-going star formation or of a young stellar population. We can then compare the fraction of E and S0 group members that have recently formed stars with a sample from rich clusters with complex structure. If early types in poor field groups have different recent star formation histories than those in infalling subclusters, which have only just been exposed to the cluster environment for the first time, then cluster-based galaxy evolutionary mechanisms are clearly potent at the current epoch. On the other hand, if the two samples are similar, then the mechanisms that operate exclusively in clusters or that are much more effective in clusters than in poor groups are not required to enhance and/or quench star formation in these galaxies. In the latter case, recent star formation may be the product of galaxy-galaxy encounters, which are present in both groups and subclusters, or of evolution that is independent of the galaxies' present environment. To determine which of the E and S0 group members have spectra that indicate on-going or recent star formation, and to compare this population with that in rich, complex clusters, we first automate the calculation of the [O II]$\\ \\lambda 3727$ equivalent width in the manner of Zabludoff \\etal (1996). The equivalent width uncertainties, which are typically less than 1 \\AA, are calculated using counting statistics (the detector is a photon counter with approximately zero read noise), the local noise in the continuum, and standard propagation of errors. Figure 8a shows the unfluxed spectra of the central regions of four early type group members with significant [OII] emission ($> 5$\\AA, excessive for a normal Sa type galaxy (Kennicutt 1992ab; Caldwell \\& Rose 1997)). This criterion is the same as that used by Caldwell \\& Rose (1997) to identify star formation in the spectra of early type galaxies in rich clusters with substructure. However, the ratios of [OIII]$\\ \\lambda 5007$ to H$\\beta \\ \\lambda 4861$ and [OII]$\\ \\lambda3727$ to [OIII]$\\ \\lambda 5007$ suggest that H$90\\_{017}$ is an AGN. Two other spectra in Figure 8a, H$42\\_{023}$ and H$62\\_{008}$, are either star forming or weaker AGN. To be conservative, we include only N$741\\_{020}$ as a star forming galaxy in the subsequent analysis. Are there early-type group members that have experienced star formation in the last few Gyr and are now quiescent? The ratio of the strength of the Ca II H $\\lambda 3968 \\ +$ Balmer H$\\epsilon \\ \\lambda 3970$ line to the Ca II K $\\lambda 3934$ line is a strong indicator of recent star formation (cf. Leonardi \\& Rose 1996), especially when there is some filling of the other Balmer absorption lines. We visually identify the spectra of group galaxies with early type morphologies that have H+$H\\epsilon$ to K ratios $> 1$ and signal-to-noise ratios $S/N > 6$. These seven spectra are shown in Figure 8b. To further test whether the galaxies in Figure 8ab have a young stellar component, we calculate the 4000 \\AA\\ break $D_{4000}$ for the sample spectra as in Zabludoff \\etal (1996). The break strength is a measure of the galaxy's color, with lower $D_{4000}$ indicating a bluer stellar population. The eight star forming and post-star forming galaxies in Figure 8ab lie in the blue tail of the $D_{4000}$ distribution (Figure 9, shaded), implying that they have a younger stellar population than is typical for the other early types in the sample (white). The unshaded galaxy blueward of the peak of the shaded subsample is the probable AGN discussed above, H$90\\_{017}$. The case for recent star formation in the shaded subsample is bolstered by stellar population synthesis models (Bruzual \\& Charlot 1995), which indicate that for a double burst model of star formation in which the second burst lasts a Gyr or less, a galaxy with $D_{4000} \\simless 2$ has experienced the second burst within the last $\\sim 2$ Gyr, regardless of the assumed initial mass function, burst strength, or progenitor type. \\subsubsection{Comparison with Rich Clusters} The star forming and post-star forming galaxies whose spectra are in Figure 8ab are $12\\%$ of the 64 early type group members for which we have spectra. How does this fraction compare with that in clusters with infalling groups? Caldwell and Rose (1997) examined galaxy spectra in clusters with significant substructure and found that $\\sim 15\\%$ of galaxies with early type morphologies have signatures of recent star formation. The resolution of our spectra is $\\sim 3\\times$ poorer than for that cluster sample, preventing us from precisely duplicating the young stellar population index used by Caldwell and Rose ({\\it e.g.}, their continuum definitions correspond to one pixel or less in our spectra). However, like all of our post-star forming early types, most of the galaxies in Caldwell and Rose's post-starburst sample have sufficiently strong Balmer $H\\epsilon$ so that H+$H\\epsilon$ EW $>$ K EW (note that Caldwell and Rose include a few possible AGN and transitional S0/a type galaxies of the kind that we exclude from the group sample). Therefore, the fractions of early types with ``abnormal\" Balmer absorption and/or [OII] emission in poor groups and in subclusters are roughly consistent. There are two potential problems in comparing the poor group and cluster samples. First, although the spectroscopy of both samples was conducted with fibers of similar size, the physical radius to which the group and cluster galaxies are sampled by the fibers differs, because the Caldwell and Rose systems are generally more distant than ours. This ``aperture bias\" might alter the apparent contributions of the young and old stellar populations in our our spectra relative to the Caldwell and Rose data. For example, our fibers will tend to miss [OII] emission or Balmer absorption in the outer parts of some E and S0's and to oversample nuclear light in comparison with the Caldwell and Rose spectra. Second, our group sample includes early type galaxies that are several magnitudes fainter than those in the cluster sample. Although a Kolmogorov-Smirnov test fails to distinguish (at better than the $95\\%$ level) between the distributions of estimated absolute magnitude $M_B$ for the early types with and without young stellar populations in our groups, an increase in the fraction of recently star forming galaxies at fainter absolute magnitudes would complicate the comparison of the group and cluster samples. One way to roughly address these two issues is to compare the Caldwell and Rose data in Coma with our data in NGC 533, because these systems are at similar distances. We cut the NGC 533 sample at the absolute magnitude limits of the Coma sample, $M_B \\simless -16.6 + 5$log$_{10}$ \\lith. Three of the 16 galaxies that that we classify as early types have young stellar populations, but no on-going star formation. This fraction of $23$\\% is not significantly different from the fraction of early types with post-starburst spectra in the NGC 4839 subcluster ($11/38 = 29\\%$; Caldwell \\& Rose 1997), whose non-coincident galaxy and X-ray gas distributions suggest that the group is falling into the Coma cluster for the first time (Caldwell \\etal 1993; White \\etal 1993a; Zabludoff \\& Zaritsky 1995). Although the statistics of the samples are poor, it is suggestive that the early type galaxies in NGC 533 have a recent star formation history similar to those in the NGC 4839 subcluster, which has the highest fraction of post-starburst early types of any subclump in a nearby cluster. The spectra in Figure 8b indicate that some of the early type galaxies in poor groups experienced recent episodes of star formation and are now quiescent. What processes might induce and/or quench star formation in a group galaxy? Mechanisms proposed to deplete gas in rich cluster galaxies include ram pressure stripping (Gunn \\& Gott 1972), transport processes like viscous stripping and thermal conduction (Nulsen 1982; Cayatte \\etal 1994), and the expulsion of gas via supernovae-driven winds (cf. Larson \\& Dinerstein 1975). The tidal limitation of a galaxy's halo by the group potential ($\\S3.3$) might also remove gas from a reservoir outside the optical radius. By removing gas from galaxies, these processes are likely to reduce subsequent star formation. Although there are no models at present, it is also conceivable that these mechanisms compress or shock inter-stellar gas, producing an increase in star formation. Because the metallicity of the intra-group gas is poorly known, it is difficult to place strong constraints on the contribution of supernovae ejecta to the intra-group medium. Therefore, we consider only the effects of ram pressure stripping and gas transport processes below. \\subsubsection{The Effects of Ram Pressure} It is possible to roughly estimate the ram pressure of the intra-group medium on the inter-stellar medium of group members. Specifically, we estimate the radius $r_s$ at which the disk of an $L^*$ galaxy is subject to stripping. The condition for stripping for a galaxy that moves with uniform velocity through a uniform intra-group medium is $\\rho_{IGM} {\\upsilon_{\\perp}}^2 > 2 \\pi G \\sigma_{tot} \\sigma_{gas}$, where $\\rho_{IGM}$ is the density of the intra-group medium, $\\upsilon_{\\perp}$ is the component of galaxy velocity relative to the intra-group medium that is perpendicular to the disk, $\\sigma_{tot}$ is the surface density of the disk at radius $r_s$, and $\\sigma_{gas}$ is the surface density of gas in the disk at $r_s$ (Gunn \\& Gott 1972). This simple condition is consistent with the predictions of n-body/hydrodynamic simulations (Kundi\\'c \\etal 1997). Because NGC 533 has the highest $\\rho_{IGM} {\\upsilon_{\\perp}}^2$ of the 12 groups, we estimate the ram pressure experienced by one of its late type members to obtain a conservative upper limit for the sample. First we calculate the ram pressure term for a galaxy in the core of NGC 533. We adopt $\\upsilon_{\\perp} \\approx \\sqrt 2 \\sigma_r =$ 650 \\ks. We integrate a King profile (King 1966) of NGC 533's intra-group medium that excludes the central elliptical (Paper II) to obtain an estimate of the central gas density, $\\rho_{0,IGM} \\sim 2 \\times 10^{-3} h$ cm$^{-3}$. The average gas density within 300\\lith\\inv\\ kpc is then $\\overline {\\rho_{IGM}} \\sim 9 \\times 10^{-4} h$ cm$^{-3}$. Therefore, the ram pressure on a galaxy in the core of NGC 533 is $\\overline {\\rho_{IGM}} {\\upsilon_{\\perp}}^2 \\sim 4 \\times 10^{12} h$ cm$^{-1}$ s$^{-2}$. To determine the radius in a disk galaxy at which ram pressure stripping is effective, we calculate where the restoring force pressure $2 \\pi G \\sigma_{tot} \\sigma_{gas}$ exceeds the ram pressure. We use the averaged total mass and HI surface density profiles of six nearby, $L^*$ or brighter galaxies from Giovanelli \\etal (1981) to obtain estimates of $\\sigma_{tot}$ and $\\sigma_{gas}$ at different radii (note that the units in their Figure 7 are mislabelled $\\Mdot$ kpc$^{-2}$, instead of $\\Mdot$ pc$^{-2}$). The exponential scale length of the total mass in the composite disk is $\\sim 4$ kpc. The ram pressure condition is satisfied at $r_s \\sim 15$ kpc, or about four scale lengths. This radius is comparable to the furthest extent of optical light in disk galaxies ($\\sim 3$-5 scale lengths; van der Kruit \\& Searle 1981). As a result, ram pressure will probably not strip gas inside the optical disk of a $L^*$ group member, but may deplete gas if a reservoir exists at larger radii, limiting subsequent star formation by preventing replenishment by infalling gas. It is less likely that ram pressure will affect the morphology of group galaxies, {\\it e.g.}, transforming the bulge-to-disk ratio of a late type spiral into that of an S0. We note that the estimated central gas density in NGC 533 is comparable to that in Coma, where $\\rho_{0,ICM} \\sim 5 \\times 10^{-3} h^{-1/2}$ cm$^{-3}$ (Hughes 1989). If we assume conservatively that the {\\it average} gas density is the same in groups and clusters, and thus that differences in ram pressure depend only on global velocity dispersion, then galaxies in a \\sigmar\\ $\\sim 1000$ \\ks\\ rich cluster experience at least $\\sim 5$-25$\\times$ more ram pressure on average than do the members of our groups. There are several limitations of this estimate of the ram pressure in poor groups. First, the calculation does not include the replenishment by supernovae of gas removed by ram pressure from the outer disk and halo. Second, we assume that gas is uniformly distributed in the disk --- any clumping makes the gas harder to strip away because the stripping force is reduced in proportion to the area of the gas clump, while the gravitational force is unaffected (T. Kundi\\'c 1997, priv. comm.). Third, the effects of ram pressure in groups could be enhanced by the tidal tails produced during the collision of two galaxies (see Hibbard \\etal 1994). The intra-group medium will strip the rarefied gas in the tails more easily than disk gas, perhaps preventing gas from returning to the merger remnant and thereby halting any subsequent star formation. Fourth, we implicitly assume that that the disk surface densities of fainter group members are at least as high as for the $L^*$ galaxy considered here. If the disks of fainter galaxies have substantially lower restoring force pressure, they may be more affected by ram pressure stripping. \\subsubsection{The Effects of Viscous Stripping and Thermal Conduction} Transport processes such as viscous stripping and thermal conduction are also proposed as gas removal mechanisms in clusters (Nulsen 1982; Cayatte \\etal 1994). In contrast to ram pressure stripping, these processes do not depend on the disk's local gravity, except for the most massive and/or slow moving galaxies. Therefore, gas is stripped from the galaxy at all radii, not just from the outermost disk. In addition, gas transport is not sensitive to the orientation of a galaxy with respect to its motion through the intra-group medium, whereas ram pressure depends on the component of the galaxy's velocity perpendicular to its disk. As in the ram pressure calculation above, we assume here that there is no replenishment of the gas once it is swept away and that the galactic gas is uniform. We also assume that ram pressure has not first stripped gas from the disk. The stripping rate due to gas transport is given by ${\\dot M_t} \\approx \\pi {r_d}^2 {\\overline {\\rho_{IGM}}} \\upsilon$, where $r_d \\sim 30$ kpc is the radius of the gas disk (excluding any ram pressure effects) and $\\upsilon \\approx \\sqrt 2 \\sigma_r$ is the approximate velocity of the galaxy through the intra-group medium (Nulsen 1982). For a galaxy within 300\\lith\\inv\\ kpc of the center of NGC 533, the gas mass loss rate is then $\\sim 40 h \\ \\Mdot$ yr$^{-1}$. Therefore, a typical galaxy with $\\sim 8 \\times 10^9 \\Mdot$ of HI (NGC 4192 from Cayatte \\etal 1994) could lose almost $100\\%$ of its atomic gas in a crossing time ($t_c/t_H = 0.02$ for NGC 533). However, the HI detected in many poor group members (Giuricin \\etal 1985) suggests that the effectiveness of this mechanism is overestimated here. As pointed out by Kundi\\'c \\etal (1997), this approximation neglects the formation of a shock in front of the galaxy that prevents some fraction of the intra-group medium from interacting directly with the galactic gas. In addition, the introduction of even a small magnetic field term into the transport equation would limit the effectiveness of the process by reducing the mean free paths of the plasma particles (B. Mathews 1997, priv. comm.). An improved understanding of the significance of gas transport mechanisms in groups awaits detailed simulations. As in the case of ram pressure, gas transport may limit the subsequent star formation of the galaxy, but the requirements for transforming the galaxy into an earlier morphological type, {\\it e.g.}, the total disruption of the stellar disk or an increase in the bulge-to-disk ratio, are not part of this picture. \\subsubsection{Implications} In the preceding discussion, we demonstrated that the effects of ram pressure stripping are weaker in poor groups than in rich clusters. We suggested that, even if gas disruptive processes like ram pressure stripping, viscous stripping, or thermal conduction can induce star formation as well as truncate it, they are unlikely to significantly affect the stellar morphology of a galaxy. Therefore, if one environmental mechanism is responsible for both the high early type fractions of some poor groups and the young stellar populations of certain early type group galaxies, galaxy-galaxy encounters provide a possible explanation. Such interactions can produce bursts of star formation ({\\it e.g.}, Londsdale \\etal 1984; Kennicutt \\etal 1987; Sanders \\etal 1988) in which the gas is consumed or stripped away. Although our observations of poor groups suggest that the effects of cluster environment are not required to produce the early type fractions and star formation episodes of nearby clusters, we suspect that the star formation histories of group and subcluster galaxies will begin to deviate after the subclump and cluster mix. Proposed gas removal processes including ram pressure stripping, the tidal limitation of galaxy halos, and galaxy harassment, which are more efficient in clusters than in groups, may eventually suppress star formation in cluster galaxies. Comparative studies of the HI content of field, group, and cluster galaxies will help to resolve this issue." + }, + "9708/astro-ph9708242_arXiv.txt": { + "abstract": "We have developed a method for the linear reconstruction of an image from undersampled, dithered data, which has been used to create the distributed, combined Hubble Deep Field\\cite{hdf+96} images -- the deepest optical images yet taken of the universe. The algorithm, known as Variable-Pixel Linear Reconstruction (or informally as ``drizzling\"), preserves photometry and resolution, can weight input images according to the statistical significance of each pixel, and removes the effects of geometric distortion both on image shape and photometry. In this paper, the algorithm and its implementation are described, and measurements of the photometric accuracy and image fidelity are presented. In addition, we describe the use of drizzling to combine dithered images in the presence of cosmic rays. ", + "introduction": "The Hubble Space Telescope (HST) is now capable of providing the superb images for which it was designed. However, the detectors on HST are only able to take full advantage of the full resolving power of the telescope over a limited field of view. In particular, the primary optical imaging camera on the HST, the Wide Field and Planetary Camera 2\\cite{tbbc+94}, is composed of four separate 800x800 pixel CCD cameras, one of which, the planetary camera (PC) has a scale of $0\\farcs046$ per pixel, while the other three, arranged in a chevron around the PC, have a scale of $0\\farcs097$ per pixel. These latter three cameras, referred to as the wide field cameras (WFs), are the primary workhorse for deep imaging surveys on HST. However, these cameras greatly undersample the HST image. The width of a WF pixel equals the full-width a half maximum of the optics in the the near-infrared, and greatly exceeds it in the blue. The effect of undersampling on WF images is illustrated by the \"Great Eye Chart in the Sky\" in Figure 1. \\begin{figure} \\centerline{\\psfig{figure=fourbyfour.ps,height=6.0in}} \\caption{In the upper left corner of this figure, we present the ``true image\", {\\it i.e.} the image one would see with an infinitely large telescope. The upper right shows the image after convolution with the optics of the Hubble Space Telescope and the WFPC2 camera -- the primary wide-field imaging instrument presently installed on the HST. The lower left shows the image after sampling by the WFPC2 CCD, and the lower right shows a linear reconstruction of dithered CCD images. \\label{fig:bigeye}} \\end{figure} Fortunately, much of the information lost to undersampling can be restored. In the lower right of Figure 1 we display a restoration made using one of the family of techniques we refer to as ``linear reconstruction.\" The most commonly used of these techniques are shift-and-add and interlacing. The image in the lower right corner has been restored by interlacing dithered images. However, due to the occasional small positioning errors of telescope and the non-uniform shifts in pixel space caused by the geometric distortion of the optics, true interlacing of HST images is often infeasible. The other standard linear reconstruction technique, shift-and-add can easily handle arbitrary dither postions, but it convolves the image yet again with the orginal pixel, adding to the blurring of the image and the correlation of the noise. The importance of avoiding unnecessarily convolving the image with the pixel is emphasized by comparing the upper and lower right hand images in Figure 1. The deterioration in image quality is due entirely to convolution of the image by the WF pixel. Here we present a new method which has the versatility of shift-and-add yet largely maintains the resolution and independent noise statistics of interlacing. ", + "conclusions": "" + }, + "9708/astro-ph9708074_arXiv.txt": { + "abstract": "{ Using CCD photometry obtained by the EROS collaboration in 1991-1993, we have discovered an LMC variable star with a light curve that is oscillating with a mean period of $\\sim 14$ days and an amplitude of $\\sim$ 0.3 mag. The oscillations appear with irregular amplitude variations. The Fourier spectrum shows that the pulsation of this star is phase locked between two modes of frequencies $f_0$ and 1.5$\\times f_0$. Moreover, this object has strong $H \\alpha$ and $H \\beta$ emission lines and neutral lines of Helium that suggest a spectral type between late O and early B. In a preliminary analysis, we derive a luminosity of $ L=3.4-3.8L_\\odot$ and an effective temperature in the range $\\log(T_{eff}) =3.85-4.2$. } ", + "introduction": " ", + "conclusions": "" + }, + "9708/astro-ph9708238_arXiv.txt": { + "abstract": "It is usual, in gamma-ray burst (GRB) studies, to compare the average properties of bright and faint GRBs, with the assumption that brightness classes reflect distance classes. When brightness is intented to reflect the distance to the sources, it is nevertheless important to use a quantity with a small intrinsic dispersion. We propose here a method to compare the intrinsic dispersion of various measures of GRB brightness. This method assumes that nearby bursters are homogeneously distributed in an Euclidean space with no density or luminosity evolution. We then use it to compare 5 measures of GRB brightness in the BATSE Catalog. Our analysis reveals that better (i.e. less dispersed) measures of brightness are obtained at low energy and that GRBs are much closer to standard candles below 100 keV than above. We suggest that a beaming of the emission above 100~keV could explain this behaviour. ", + "introduction": "The GRB intensity distribution has been extensively used to improve our understanding of these sources. The \\vvmm\\ test for instance has been unvaluable to demonstrate the burster spatial inhomogeneity (Meegan et al. 1992). The LogN-LogP distribution has been shown to be in good agreement with the intensity distribution expected for cosmological sources (e.g. Piran 1992, Wickramasinghe et al. 1993, Fenimore et al. 1993). Finally, it is common to define brightness classes for the purpose of searching cosmological signatures in gamma-ray bursts. In all these studies, the burst brightness appears as a key parameter. While it is widely acknowledged that C$_{max}$/C$_{min}$ is the measure of brightness which is appropriate to check the spatial homogeneity of gamma-ray bursters (Schmidt et al. 1988), there is no such agreement on the parameter which should be used for studies where brightness is taken as a distance indicator. These studies require a quantity with small, or no, intrinsic dispersion. We propose below a new way of comparing the intrinsic dispersions of different measures of brightness. When several definitions of brightness are available (e.g. peak flux, fluence...), this method answers the question of which one of these quantities is closer to a standard candle. Section 2 describes the method. In section 3, we apply it to a sample of 1471 GRBs detected by BATSE. The implications for gamma-ray bursts are discussed in section 4. ", + "conclusions": "We have proposed a new method to compare the intrinsic dispersion of various measures of GRB brightness. An essential assumption of this method is that the slope -3/2 at the bright end of the size-frequency curve is due to the spatial homogeneity of nearby bursters. If on the other hand, source evolution dominates the GRB size-frequency curve, the slope -3/2 might have a completely different interpretation, invalidating the conclusions of our analysis (for instance density evolution could cancel the effects of a non-Euclidean space, simulating a homogeneous distribution of bursters in an Euclidean space). In the rest of the discussion we assume that source evolution does not dominate the observed GRB size-frequency distribution. When applied to GRBs in the Current Catalog of BATSE our method shows that the measure of luminosity with the smallest intrinsic dispersion is the time integrated luminosity below 100~keV. This result is however not complete since we restricted our analysis to 5 measures of brightness given in the Catalog. This study calls for a careful definition of the GRB brightness when this quantity is used as a distance indicator (e.g. to compare the properties of nearby and distance bursters); measures at low energies are then clearly preferred. We finally note that the combination of (1) a broad luminosity function and (2) a spatial density which varies with the distance has interesting consequences for the comparison of faint and bright GRBs. Intrinsically bright bursts are detected to large distances (typically larger than the size of the homogeneous region) where the burster spatial density decreases rapidly. Intrinsically faint bursts on the other hand are only visible to much smaller distances where the burster density is constant (if we remain in the homogeneous region) or slowly decreasing. As a consequence, going to lower intensities increases the number of bright GRBs much less (in percentage) than the number of intrinsically faint bursts. In other words, burst classes based on the {\\it observed} brightness do not contain the same proportion of {\\it intrinsically} bright bursters. This changing proportion may produce brightness-dependent average burst properties (spectral and/or temporal) which could strengthen or counteract cosmological effects. Because our study suggests that the GRB luminosity function is more extended above 100~keV, we expect GRB properties to be more brightness-dependent when brightness is measured above 100~keV. For instance the well known hardness-intensity correlation (e.g. Dezalay et al. 1997) could be explained in this way if it appears that it is stronger when the intensity is measured at higher energies. While we do not address here the reasons which make the luminosity at low energies a better standard candle, we note that this behaviour could well be explained by an anisotropy of the emission above $\\sim$100~keV. Such an anisotropy would make the brightness at high energies dependent on the aspect of the source. From the point of view of the size-frequency distribution, a beaming of the emission is equivalent to the existence of a luminosity function. Hence a beaming factor which changes with the energy may just appear as an energy dependent luminosity function, which is precisely what we observe." + }, + "9708/astro-ph9708148_arXiv.txt": { + "abstract": "We present numerical simulations of an isothermal turbulent gas undergoing gravitational collapse, aimed at testing for ``logatropic'' behavior of the form $P_t \\sim \\log \\rho$, where $P_t$ is the ``turbulent pressure'' and $\\rho$ is the density. To this end, we monitor the evolution of the turbulent velocity dispersion $\\sigma$ as the density increases during the collapse. A logatropic behavior would require that $\\sigma \\propto \\rho^{-1/2}$, a result which, however, is not verified in the simulations. Instead, the velocity dispersion {\\it increases} with density, implying a polytropic behavior of $P_t$. This behavior is found both in purely hydrodynamic as well as hydromagnetic runs. For purely hydrodynamic and rapidly-collapsing magnetic cases, the velocity dispersion increases roughly as $\\sigma \\propto \\rho^{1/2}$, implying $P_t\\sim \\rho^2$, where $P_t$ is the turbulent pressure. For slowly-collapsing magnetic cases the behavior is close to $\\sigma \\propto \\rho^{1/4}$, which implies $P_t \\sim \\rho^{3/2}$. We thus suggest that the logatropic ``equation of state'' may represent only the statistically most probable state of an ensemble of clouds in equilibrium between self-gravity and kinetic support, but does not adequately represent the behavior of the ``turbulent pressure'' within a cloud undergoing a dynamic compression due to gravitational collapse. Finally, we discuss the importance of the underlying physical model for the clouds (in equilibrium vs.\\ dynamic) on the results obtained. ", + "introduction": "Molecular clouds and clumps exhibit the well-known velocity dispersion- (or linewidth-)size relation \\begin{equation}\\label{larvel} \\sigma \\sim R^{1/2}, \\end{equation} where $\\sigma$ is the linewidth-determined velocity dispersion, and $R$ the characteristic size. This correlation is observed both in ensembles of clouds (Larson 1981; Leung et al.\\ 1982; Torrelles et al.\\ 1983; Dame et al.\\ 1986; Myers \\& Goodman 1988; Falgarone, et al.\\ 1992; Miesch \\& Bally 1994) or as a function of radius in quiescent cores using various tracers (Fuller \\& Myers 1992; Caselli \\& Myers 1995; Goodman et al.\\ 1997), although the latter studies have suggested that the scaling exponent in relation (\\ref{larvel}) may actually differ between massive and low-mass cores. Furthermore, Goodman et al.\\ have suggested that the exponent may decrease and approach zero as the innermost regions of the cores are considered, in which the turbulent velocity dispersion becomes subsonic. A second scaling relation, between mean density $\\langle\\rho\\rangle$ and size, is also generally reported, reading \\begin{equation}\\label{lardens} \\langle\\rho\\rangle \\sim R^{-1}, \\end{equation} although its authenticity has been questioned on theoretical (Kegel 1989; Scalo 1990) and numerical (V\\'azquez-Semadeni, Ballesteros-Paredes \\& Rodr\\'iguez 1997) grounds, and significantly discrepant scaling exponents have been found (e.g., Carr 1987; Loren 1989), or none at all (e.g., Plume et al.\\ 1997). Relations (\\ref{larvel}) and (\\ref{lardens}) constitute the now famous ``Larson's relations''. In spite of the anomalies at small scales and high-mass regions, Larson's relations are generally accepted as distinctive signatures of turbulence in molecular clouds and clumps (e.g., Larson 1981; Scalo 1987), and together they imply \\begin{equation}\\label{virial} \\sigma \\propto \\rho^{-1/2}. \\end{equation} Relation (\\ref{virial}) is actually a manifestation of virial equilibrium between the turbulent velocity dispersion (possibly magnetohydrodynamic, or MHD) in the clouds (Larson 1981). A turbulent ``pressure'' $P_t$ corresponding to the turbulent velocity dispersion can then be defined by (Lizano \\& Shu 1989; hereafter LS) \\begin{equation}\\label{turbpres} \\sigma^2 \\equiv (dP_t/d\\rho), \\end{equation} in analogy with the relation between thermal pressure and the sound speed. Choosing \\begin{equation}\\label{logatrope} P_t\\propto \\log \\rho \\end{equation} recovers the virial relation (\\ref{virial}) (LS). Relation (\\ref{logatrope}) is commonly refereed to as a ``logatropic equation of state'' \\footnote{Strictly speaking, this is not an equation of state, since it does not invlove all three thermodynamic variables. However, we will allow ourselves the terminology for consistency with common nomenclature.}, or simply, a ``logatrope''. It is important to emphasize that the concept of a turbulent ``pressure'' may not be a very realistic representation of the effects of turbulence, since it implicitly assumes a microscopic and isotropic process. Instead, turbulence is a phenomenon involving a wide range of spatial scales, from the scale size of the system under consideration to the smallest dissipative scales. In particular, the existence of large-scale modes implies coherent motions which are more akin to ram pressure (locally anisotropic, with a well-defined direction) than to an isotropic, thermodynamic pressure. For these reasons, in the present paper we will focus primarily on the turbulent velocity dispersion. References to the turbulent ``pressure'' will be made assuming that it can be defined according to eq.\\ (\\ref{turbpres}), for compatibility with published work, but the above caveat should always be kept in mind. The logatropic equation of state has been used in a number of studies of cloud support and stability, such as quasi-static contraction (LS), nonlinear wave propagation (Adams \\& Fatuzzo 1993; Adams, Fatuzzo \\& Watkins 1994; Gehman et al.\\ 1996), and gravitational stability (McLaughlin \\& Pudritz 1996), among others. Nevertheless, the logatropic equation remains a completely empirical assumption, and there is no direct evidence that turbulent pressure indeed behaves in this manner in fully dynamic situations. In fact, there are some lines of reasoning which suggest that it may not: \\noindent 1. Larson's relations are observed in either {\\it ensembles} of relaxed clouds (e.g., Torrelles et al.\\ 1983; Dame et al.\\ 1986; Myers \\& Goodman 1988; Falgarone, et al.\\ 1992; Miesch \\& Bally 1994), or as a function of radius in {\\it quiescent} cores (Fuller \\& Myers 1992; Caselli \\& Myers 1995; Goodman et al.\\ 1997), but there is no evidence that they hold in fully dynamical processes, such as gravitational collapse. Interestingly, clouds which are strongly perturbed also seem to not follow the Larson scaling relations (e.g., Loren 1989; Plume et al.\\ 1997). In general, there is little sampling of fully out-of-equilibrium, dynamical processes and, in particular, dynamical collapse has not yet been directly detected because dynamical velocities occur only at very small scales. \\noindent 2. Although it might be argued that an ensemble of molecular clouds provides a complete sample of various dynamical stages, in actuality most observations refer to clouds and clumps close to {\\it equilibrium}. Thus, the clouds included in surveys such as Larson's (1989) constitute an ensemble of equilibrium states for clouds of different masses rather than an ensemble of evolutionary steps for a single cloud (of constant mass). That is, instead of representing a number of different states for the same cloud, they represent the same state for different clouds. \\noindent 3. V\\'azquez-Semadeni et al.\\ (1997) have suggested that there may exist large numbers of low-column density clouds which do not satisfy the density-size relation, and that possibly only the highest-column density clouds follow such a scaling relation. Thus, the logatropic equation of state is not expected to apply to such low-column density clouds, which are probably not in self-gravitating equilibrium, but rather pressure- or ram pressure-confined. As a first attempt to decide on this matter, in this paper we present two-dimensional numerical simulations of a turbulent, self-gravitating, magnetized, isothermal gas, aimed at testing the variation of the velocity dispersion as a cloud is compressed by self-gravity. A related calculation has been performed by Bonazzola et al.\\ (1987), who used low-resolution simulations to estimate the correlation between the nonlinear advection term (related to the turbulent pressure) and the density gradient in a compressible turbulent flow. We emphasize that the simulations discussed in this paper are not presented as models of cloud cores and their observed linewidths, but only as numerical ``experiments'' designed to test the applicability of the logatropic equation of state. Furthermore, throughout this paper we will refer exclusively to the {\\it non-thermal} part of the velocity dispersion. In contrast with the observational situation, where the separation between the thermal and non-thermal components is an issue (e.g., Fuller \\& Myers 1992), in the simulations this is a trivial task, since there is no confusion between the fluid velocity and the thermal velocity dispersion, the latter being directly represented by the temperature field. The outline of the present paper is as follows. In \\S \\ref{model} we describe the numerical model; in \\S \\ref{results} we present the results, for both purely hydrodynamic and fully MHD cases, and in \\S \\ref{conclusions} we summarize and discuss our results. ", + "conclusions": "\\subsection{Summary and discussion}\\label{sumandisc} We have argued that Larson's (1981) relations and the resulting logatropic ``equation of state'' (relation [\\ref{logatrope}]) and virial condition (relation [\\ref{virial}]) may describe an ensemble of clouds {\\it in (near) equilibrium} between self-gravity and the turbulent velocity dispersion, but not out-of-equilibrium, dynamical processes occurring on a single cloud. We have tested this assertion by means of numerical simulations of collapsing clouds with an initial turbulent velocity field, in both magnetic and non-magnetic regimes. The simulations exhibit in all cases a turbulent velocity dispersion which {\\it increases} with mean density as the collapse proceeds, in contradiction with the expected behavior for a logatrope, relation (\\ref{virial}). Non-magnetic and strongly self-gravitating runs seem to approach a power-law behavior of the form $\\sigma \\sim \\langle \\rho \\rangle^{1/2}$, while weakly self-gravitating magnetic runs in general tend to have shallower dependences, although always with positive exponents. In particular, the fact that magnetic runs exhibit the same qualitative behavior suggests that weak magnetic fields cannot induce a logatropic (or a $\\sigma \\sim \\rho^{-1/2)}$ behavior either. Interestingly, run M512 exhibits a behavior very close to $\\sigma \\sim \\rho^{1/4}$. Assuming that this run has converged to the true slope, it is noteworthy that the implied turbulent pressure satisfies $P_t \\sim \\rho^{3/2}$. This result is consistent with that of McKee \\& Zweibel (1995) for the polytropic index of Alfv\\'en waves under slow compression. However, runs MM256 and MM512 appear to be closer to the $\\sigma \\sim \\rho^{1/2}$ ($P_t \\propto \\rho^2$) behavior observed in the non-magnetic runs. This distinction is likely to be due to the larger Jeans length used in the M runs, implying a slower collapse (final time $t_{\\rm fin}=2.4$) than for the MM runs ($t_{\\rm fin}=1.0$), so that the M runs are closer to the slow compression assumption of McKee and Zweibel. We emphasize that although convergence may not have been fully achieved yet at the highest resolution we used ($512 \\times 512$ grid points), the trend is towards {\\it faster} increase of the velocity dispersion with density at higher resolution, away from the behavior predicted by the logatropic equation. Thus, the result that the velocity dispersion increases with mean density appears quite robust. Moreover, the trend of increasing $\\sigma$ with $\\langle \\rho \\rangle$ is preserved in run 3D96 (albeit at a slower rate due to the lower resolution of this run), thus ruling out the possibility that our results are a purely 2D effect. The main consequence of our results is that the logatropic ``equation of state'' appears to be inadequate for the description of dynamical processes occurring in a cloud. This implies that the use of the logatropic equation in studies of gravitational collapse and dynamical stability is questionable. Its use in studies of quasi-static processes (e.g., LS) may still be justified, although in general the question remains open as to whether the logatropic equation, which can be thought of as representing the {\\it final} states of the virialization process, also represents the behavior of the turbulent pressure {\\it during} the relaxation processes which lead to virialization. For this reason, it would also be interesting to test its applicability in problems of nonlinear wave propagation (e.g.,Adams \\& Fatuzzo 1993; Adams, Fatuzzo \\& Watkins 1994; Gehman et al.\\ 1996). Finally, we remark that the ensemble consisting of the evolutionary states of our simulated gravitationally-collapsing clouds with a fixed mass is completely different from the ensemble constructed from the observations of many clouds of different masses in near equilibrium. The present work shows that for the former ensemble, the logatropic equation of state is not applicable. \\subsection{Comparison with previous work}\\label{comparison} In our simulations we obtain a polytropic form ($P_t \\propto \\rho^{\\gamef}$) for the effective ``equation of state'' of the turbulence, with polytropic exponents $\\gamef=2$ for the non-magnetic and strongly self-gravitating cases, and $\\gamef=3/2$ for the weakly self-gravitating cases. This result appears to be in contradiction with that of McLaughlin \\& Pudritz (1996, hereafter MP), who conclude that the total pressure (thermal plus turbulent) is {\\it not} expected to behave as a polytrope. MP reach this conclusion on the basis of a stability analysis, noting that truncated polytropic solutions with $0 < \\gamef < 1$ (consistent with the observed lower temperatures of denser structures) have never unstable, or even critically stable, equilibrium solutions. That is, absolutely stable configurations are discarded by MP so that a cloud is eventually able to collapse and form a star. It can then be seen that the difference between our results and those of MP arises from the consideration of different physical models for the clouds. While MP's clouds are in hydrostatic equilibrium, our clouds are always out of equilibrium and are already unstable from the start. These may originate from clumps rendered unstable by external turbulent compressions (V\\'azquez-Semadeni et al.\\ 1996) if the effective equation of state (i.e., the heating and cooling) permits it. In such cases, the clumps never need to pass through a static equilibrium state. Another possibility is the well-known onset of gravitational instability due to the loss of magnetic support caused by ambipolar diffusion (e.g, Nakano 1979; LS). Finally, we note that the polytropic exponents implied by our simulations are larger than the critical value $\\gamma_c=4/3$ below which gravitational collapse can proceed to a singularity (e.g., Chandrasekhar 1961). Thus, if turbulent pressure continued with this behavior unrestrictedly, it would eventually halt the collapse. However, we do not expect this occur since, at late stages of the collapse, dissipation becomes important again due to the large velocity gradients that develop. In fact, in fig.\\ 3 an end to the steady increase of $\\sigma$ is seen at large values of $\\langle \\rho \\rangle$ for several of the runs. Thus, we speculate that turbulent pressure cannot by itself halt the collapse." + }, + "9708/astro-ph9708195_arXiv.txt": { + "abstract": "Continuous CCD photometry of Nova Aquilae 1995 was performed through the standard $B,V,R$ and $I$ filters during three nights in 1995 and with the $I$ filter during 18 nights in 1996. The power spectrum of the 1996 data reveals three periodicities in the light curve: 0.2558 d, 0.06005 d and 0.079 d, with peak-to-peak amplitudes of about 0.012, 0.014 and 0.007 mag. respectively. The two shorter periods are absent from the power spectrum of the 1995 light curve, while the long one is probably already present in the light curve of that year. We propose that V1425 Aql should be classified as an Intermediate - Polar CV. Accordingly the three periods are interpreted as the orbital period of the underlying binary system, the spin period of the magnetic white dwarf and the beat period between them. Our results suggest that no later than 15 months after the outburst of the nova, accretion processes are taking place in this stellar system. Matter is being transferred from the cool component, most likely through an accretion disc and via accretion columns on to the magnetic poles of the hot component. ", + "introduction": "\\subsection{V1425 Aquilae} Nova Aql 1995 was discovered on 1995 February 7 by Takamizawa (1995). Mason et al. (1996) used its similarity to Nova V1668 Cygni 1978 to extrapolate a maximum brightness of $M_{V}\\approx 6.2$ and $t_{2V}\\approx 11$ d. This classifies Nova Aquilae as a fast one. They also deduces from IR considerations that the dust shell of the nova was optically thin shortly after its outburst. The probable precursor star of about 20 mag. was detected by Skiff (1995) and this indicates an outburst amplitude of about 14 mag. Retter, Leibowitz \\& Kovo-Kariti (1996) and Leibowitz, Retter \\& Kovo-Kariti (1997) reported the discovery of two photometric periods in the light curve of Nova Aquilae 1995, and pointed out the resemblance of this object to Intermediate Polars systems. In this work we present further arguments for our interpretation of the periods of V1425 Aql as resulting from the Intermediate Polar nature of the nova. \\subsection{Intermediate Polars and nova systems} Intermediate polars (for reviews see Patterson 1994, Hellier 1995, 1996, 1997) are binary systems, which are a sub-class of AM Her stars (magnetic Cataclysmic Variables). Unlike in other members of the AM Her class, in Intermediate Polars the rotation of the primary white dwarf is not synchronized with the orbital motion of the binary system. The spin periods found in Intermediate Polars are usually much shorter than their orbital periods (Patterson 1994, Hellier 1996). One of the main observational characteristics of Intermediate Polars is the presence of multiple periodicities in the power spectra of their light curves, emanating from the non-synchronous rotation of the primary with the orbital revolution. In fact this characteristic has become, together with modulations of the X-ray radiation of the system, a major criterion for membership in the Intermediate Polar group. The classification of a classical nova as an Intermediate Polar is not new. DQ Her is a prototype of a nova which is an Intermediate Polar (Patterson 1994), and in the last few years a few other novae have been so classified as well. Examples are V533 Her and GK Per, with white dwarf spin periods of 64 and 351 sec, respectively (Patterson 1994). The detected variations of V603 Aql in the optical, X-ray and ultraviolet regimes make it also a likely Intermediate Polar system (Udalski \\& Schwarzenberg-Czeny 1989, Schwarzenberg-Czeny, Udalski \\& Monier 1992). The presence of multiple periods in the light curve of HZ Pup supports the membership of this nova in the Intermediate Polar group, too (Abbott \\& Shafter 1997). In V1974 Cygni, evidence was found to the presence of an accretion disc in the system, (Retter, Ofek \\& Leibowitz 1995, Retter, Leibowitz \\& Ofek 1996, 1997a, 1997b, 1997c, Skillman et al. 1997), together with some indications for an intense magnetic field on the surface of the white dwarf (Chochol et al. 1997). The combination of these two properties makes V1974 Cyg another potential candidate for the Intermediate Polar group. It would therefore appear that the Intermediate Polar phenomenon is not uncommon among classical nova CV systems. ", + "conclusions": "\\subsection{Identification of the two periods} Three periodicities have been identified in the light curve of V1425 Aql 15 months after its outburst; one of them, P$_{1}$, was present in the light curve already three months after outburst. The three periods reflect three genuine modulations of the light emanating from this stellar system. They are not independent of each other as one of them is the beat between the other two. Thus it appears that in 1996 two independent clocks are operating in the nova system, each one modulates the light radiation directly and also indirectly through some combination with the other clock. We suggest that the longest periodicity, P$_{1}\\sim$ 6.14 h, is the orbital period of the nova underlying binary system. This is based mainly on the fact that this period was present in the light curve already in 1995, with no apparent change in its value. The observations by Mason et al. (1996), and their conclusion, that the dust shell of the nova was not optically thick at the time of their and our observations, make it indeed likely that the binary system could have been seen with optical light at that time. Radial velocity measurements should confirm or refute our suggestion. The amplitude of the 6.14-h variation, if indeed present in the 1995 light curve, is comparable, in magnitude units, to the amplitude of this variation in 1996. In 1996, however, the optical brightness of the system was some four mag. fainter than in 1995. This seems to us to suggest that the major source of the binary modulation in the light curve is the reflection effect. The amplitude of modulations by this effect depends on the inclination angle of the system and on the fraction of the radiation from the hot component, that is being intercepted by the companion, and reflected in the direction of the observer. The first parameter is clearly an invariant of the system. The second one may also be constant to first order, provided that the dimensions of the hot component did not change appreciably, relative to the radius of the companion, between the two years. This could be the case either because the size of the hot source, indeed did not change by a large factor from 1995 to 1996, or because in 1995, the size of the hot source, e.g. the white dwarf or the pseudo-photosphere of the white dwarf, was already small relative to the radius of the secondary star. We further propose that Nova Aql 95 is an Intermediate Polar, in which the shorter period, P$_{2}$, is the spin period of the magnetic white dwarf in the system. Spin period that is shorter than the orbital is the rule in almost all known Intermediate Polars. There is only one exception, RX J19402-1025, but this is a nearly synchronous system, in which the spin period has only a marginal excess: $(P_{spin} - P_{orbital}) / P_{orbital}\\sim 0.3$\\% (Patterson et al. 1995). The double structure shape of the 86.5-m period (Fig. 4 middle panel) is typical to the variations emanating from one pole of a rotating magnetic white dwarf (Warner 1995). This evidence supports the interpretation of this period as the spin period of the white dwarf. Finally, the third period is interpreted as the beat period between the two other periods. Its independent presence in the light curve of Intermediate Polar system is well established, observationally as well as in theory (Patterson 1994, Hellier 1995, 1996, 1997). \\subsection{The accretion form -- an Accretion Disc (?)} The fate of the accretion disc in a classical nova binary system during the outburst event and immediately following it is unclear. It is sometimes assumed that if an accretion disc is present in the pre-outburst system, it is being destroyed by this cataclysmic event. However, no theoretical effort was done in the direction of answering the question when is the accretion disc rebuilt in the shattered, slowly decaying system. Up to recent times there were little observational data concerning the existence of the accretion disc in young novae. In the last few years, however, observational data are being accumulated, indicating an early presence of an accretion disc in a few classical novae, already a few weeks or months after the outburst. Leibowitz et al. (1992) discovered an eclipse three weeks after maximum light in Nova V838 Herculis 1991. They interpreted it as the occultation of the accretion disc by the secondary star. Some 30 months after the outburst of the classical nova V1974 Cyg, Retter, Leibowitz \\& Ofek (1997a, 1997b, 1997c) and Skillman et al. (1997) detected permanent superhumps in the light curve of this system. Superhumps characterize the SU UMa class of CVs that are known to have an accretion disc in their underlying stellar system. Thus the observations in V1974 Cyg indicate the early presence of an accretion disc also in that system. It is now believed that in nearly all known Intermediate Polars, a major part of the accretion proceeds through an accretion disc most of the time (Patterson 1994, Hellier 1996). In a few systems the accretion is partly maintained by a different mode - disc overflow, in which the accretion stream from the companion bypasses the accretion disc and interacts directly with the magnetic field of the white dwarf. (Hellier 1993, 1995, 1996, 1997, Hellier \\& Livio 1994). However, only one object (RX J1712-24) out of the 13 Intermediate Polars listed by Hellier 1996 (see his Fig. 2) is believed to be a disc-less system. Based on this statistics, and ignoring the possibility that it is biased by a selection effect due to the excessive brightness of the disc, we may regard it as very likely that no later than 1996 May, Nova Aql 95 already contained an accretion disc within its binary system. Unless the asymmetries in the system are large, the relative amplitude of the spin period to the beat period is a first order measure of the rate of accretion through an accretion disc relative to the rate of accretion via the disc overflow mode (Hellier 1997). In N. Aql 95, if the 86.5-m period is indeed the spin period of the white dwarf, the dominant accretion is via the accretion disc, while a smaller part of it is maintained through the disc overflow mode. This is implied by the (peak-to-peak) amplitude 0.007 mag. of the beat period 0.079 d, that is a half of the amplitude 0.014 of the 0.06005-d spin period. The 86.5-m spin periodicity in V1425 Aql is one of the longest among Intermediate Polars (Patterson 1994, Hellier 1996), especially if the three nearly synchronous systems (Nova V1500 Cygni 1975, BY Cam and RX J19402-1025 - Patterson et al. 1995) are not counted in this class (Warner 1995 groups them with the AM Her systems). Hellier (1996) and Allan et al. (1996) speculated that slow rotators accrete through accretion curtains. A 86.5-m spin cycle makes Nova Aquilae 95 a slow rotator and therefore a system with that type of accretion mode. In the accretion curtains model, the pulse structure is expected to consist of a single hump, because the two poles in this mode, act in phase. In the middle panel of Fig. 4 we see, however, that the pulse in N. Aql 95 has a structure of mid-way between one and two humps. According to Hellier, this would indicate that in this system polecaps are modulating the optical radiation to a large extend, in spite of the slow rotation of the white dwarf. If these ideas are correct, it is another evidence for the presence of an accretion disc in the system. According to the spin-amplitude relation of Patterson (1994) the amplitude of the shorter, spin variation should increase as the nova continues to fade. Further confirmation for the Intermediate Polar nature of this system may come also from future X-ray and polarization measurements. \\subsection{Magnetic Novae and Speed Class} It was proposed by Diaz \\& Steiner (1991) and by Orio, Trussoni \\& Ogelman (1992) that classical novae with strong magnetic fields tend to be of higher speed class. Our findings that Nova Aql 95 is an Intermediate Polar seems to be in line with this suggestion as this nova is a fast one with $t_{2V}\\approx 11$ d (Mason et al. 1996). We checked this claim against observational data concerning some 60 classical novae (Warner 1995), of which 12 are Intermediate Polars or suspected Intermediate Polars. We found that the histograms of both t$_{2}$ and t$_{3}$ parameters of the population of Intermediate Polars novae are not significantly different from the corresponding histograms for the non-magnetic systems. Our test, however, cannot be considered as evidence against the above contention, since the statistic is rather poor due to the still small number of recognized magnetic classical novae." + }, + "9708/astro-ph9708088_arXiv.txt": { + "abstract": "We compute the number counts of clusters of galaxies, the log$N$--log$S$ relation, in several X-ray and submm bands on the basis of the Press--Schechter theory. We pay particular attention to a set of theoretical models which well reproduce the {\\it ROSAT} 0.5-2 keV band log$N$--log$S$, and explore possibilities to further constrain the models from future observations with {\\it ASCA} and/or at submm bands. The latter is closely related to the European {\\it PLANCK} mission and the Japanese Large Millimeter and Submillimeter Array ({\\it LMSA}) project. We exhibit that one can break the degeneracy in an acceptable parameter region on the $\\Omega_0 - \\sigma_8$ plane by combining the {\\it ROSAT} log$N$--log$S$ and the submm number counts. Models which reproduce the {\\it ROSAT} band log$N$--log$S$ will have $N(>S) \\sim (150-300) (S/10^{-12}\\mbox{erg~cm$^{-2}$~s$^{-1}$})^{-1.3}$ str$^{-1}$ at $S \\simgt 10^{-12}$erg~cm$^{-2}$~s$^{-1}$ in the {\\it ASCA} 2-10 keV band, and $N(>S_\\nu) \\sim (10^2-10^4) (S_\\nu/100\\mbox{mJy})^{-1.5} \\mbox { str}^{-1}$ at $S_\\nu \\simgt 100\\mbox{mJy}$ in the submm (0.85mm) band. The amplitude of the log$N$--log$S$ is very sensitive to the model parameters in the submm band. We also compute the redshift evolution of the cluster number counts and compare with that of the X-ray brightest Abell-type clusters. The results, although still preliminary, point to low density ($\\Omega_0\\sim 0.3$) universes. The contribution of clusters to the X-ray and submm background radiations is shown to be insignificant in any model compatible with the {\\it ROSAT} log$N$--log$S$. ", + "introduction": "Clusters of galaxies are among the largest virialized structures in the universe and their importance as a cosmological probe is well-recognized. Thus their observations have been actively carried out in a variety of wavelengths including X-ray, optical, infrared, and radio bands. In the present paper, we focus on the theoretical predictions for the number counts of clusters of galaxies, the \\ns relation, rather than more conventional statistics such as the X-ray temperature and luminosity functions (hereafter XTF and XLF) for several reasons; 1) temperature of X-ray clusters can be reliably determined only for luminous ones, and thus the statistics is inevitably limited, 2) such obtained XTF is to some extent weighted towards relatively rich clusters, and may be biased for the luminous species, 3) the XTF and XLF at high redshifts ($z \\simgt 0.1$) are in fact model-dependent statistics, because the translation of the observed X-ray flux to the absolute luminosity, and of the observed number to the comoving number density can be done only by assuming specific values of the cosmological parameters (the density parameter, $\\Omega_0$, the dimensionless cosmological constant, $\\lambda_0$, and the Hubble constant $H_0$ in units of 100 km/s/Mpc, $h$). On the other hand, the \\ns relation is almost free from the above problems as long as the cluster identification (or separation from point-like sources) and the conversion of count rates to fluxes are reliable. Recent analysis of the {\\it ROSAT} Deep Cluster Survey (RDCS, Rosati et al. 1995, 1997) and the {\\it ROSAT} Brightest Cluster Sample (BCS, Ebeling et al. 1997a,b) has determined the \\ns of clusters over almost four orders of magnitude in flux, i.e. $S(\\mbox{0.5-2.0 keV}) \\sim 10^{-14}-10^{-10}$ \\unit. The number of identified clusters in the \\ns is over 200, an order of magnitude larger than that for the commonly used XTF based on Henry \\& Arnaud (1991), and therefore the \\ns data are statistically more reliable. Kitayama \\& Suto (1997, hereafter KS97) found that a set of cold dark matter (CDM) models reproduce the above {\\it ROSAT} \\ns data remarkably well over whole observed flux range, and simultaneously agree with the observed XTF and the {\\it COBE} 4 year data. Nevertheless, there still exist some degeneracy of acceptable cosmological parameters. In the present paper, we explore possibilities to further constrain the models so as to break such degeneracy and discuss their implications in the following manner. First, we combine the \\ns relations at different wavelengths, in X-ray and submm bands. The latter is of particular significance in relation to the future projects including the European {\\it PLANCK} mission and the Japanese Large Millimeter and Submillimeter Array ({\\it LMSA}) project. The emissions from intracluster gas in X-ray and submm bands are originated from completely different physical mechanisms; the former is mainly due to thermal bremsstrahlung, and the latter is due to the inverse-Compton scattering of the cosmic microwave background (CMB) photons, i.e. the Sunyaev \\& Zel'dovich (1972, hereafter SZ) effect. As a result, the \\ns relations in these bands show very different parameter dependence. We note that the submm \\ns was computed earlier by several authors (e.g., Barbosa et al. 1996; Colafrancesco et al. 1997). Our analysis below differs from theirs in considering several CDM models consistent with the {\\it ROSAT} \\ns data, in including the relativistic correction to the SZ effect, and in making quantitative and extensive predictions for the number counts on the $\\Omega_0 - \\sigma_8$ plane. Second, we consider the cluster number counts incorporating redshift and/or temperature information in addition to the flux. In this way, we are able to discuss the evolution of cluster abundances on the basis of a cosmological model-independent and bias-free observable at high redshifts, which is in contrast to the approaches based on the XTF or XLF. We demonstrate a tentative comparison of our predictions with an observed sample from the {\\it ROSAT} All Sky Survey, the X-ray brightest Abell-type clusters (Ebeling et al. 1996). Finally, we discuss the implications of our results for the X-ray background (XRB) and the submm background radiation (SBR). Since the \\ns relation is closely related to the background radiation in the corresponding energy band, we may rigorously constrain the contribution of clusters of galaxies to the XRB and SBR. ", + "conclusions": "We have presented several cosmological implications of the number counts of clusters of galaxies. We have paid particular attention to the theoretical models which are in good agreement with the {\\it ROSAT} \\ns in the soft X-ray (0.5-2 keV) band, and explored possibilities to further constrain the models from future observations in the {\\it ASCA} hard X-ray and submm bands. In the submm band ($0.85$mm), models which reproduce the {\\it ROSAT} \\ns predict $N(>S_\\nu) \\sim (10^2-10^4) (S_\\nu/100\\mbox{mJy})^{-1.5} \\mbox { str}^{-1}$ at $S_\\nu\\simgt 100\\mbox{mJy}$. We have shown that the amplitude of the above relation depend sensitively on $\\Omega_0$ and $\\sigma_8$, and in a substantially different manner from the {\\it ROSAT} \\ns. Thus, combining the two can break the degeneracy in the acceptable parameter region on the $\\Omega_0 - \\sigma_8$ plane. This indicates that the future observations by the European {\\it PLANCK} mission and the Japanese {\\it LMSA} project would provide powerful probes of these parameters. In the 2-10 keV band, the number counts show similar parameter dependence to those in the {\\it ROSAT} 0.5-2 keV band, and we predict $N(>S) \\sim 200 (S/10^{-12}\\mbox{\\unit})^{-1.3}$ str$^{-1}$ at $S\\simgt 10^{-12}$\\unit in the {\\it ASCA} 2-10 keV band. The {\\it ASCA} \\ns would therefore provide an important cross-check for our interpretation of the {\\it ROSAT} \\ns data. The evolutionary behavior of the number counts is also important to put additional cosmological constraints. We have exhibited that, given a complete flux limited cluster sample with redshift and/or temperature information, one can further constrain the cosmological models. We have performed a tentative comparison between our theoretical predictions and the recent compilation of the XBACs by Ebeling et al. (1996), which is the largest sample of galaxy clusters available to date. While the incompleteness of the sample and uncertainties in the temperature data still make it difficult to draw any definite conclusions from this comparison, it is interesting to note that our predictions reproduce well the evolutionary features of the XBACs and that the results, although preliminary, seem to favor low density ($\\Omega_0 \\sim 0.3$) universes. The cluster \\ns also provides a tight constraint on their contribution to the background radiation in the corresponding energy band. Based on the \\ns relation observed by {\\it ROSAT} in the 0.5-2 keV band, we conclude that clusters of galaxies contribute at most $\\sim 20$\\% of the total XRB and less than $\\sim 5$\\% of the SBR. \\bigskip \\bigskip \\vspace{1pc} \\par We are grateful to H. Ebeling and P. Rosati for generously providing us their X-ray data and helpful comments, and to K. Masai for making his X-ray code available to us. We also thank A. Blanchard for valuable discussions, Y. Rephaeli for useful correspondences on the SZ effect, and the referee G. Zamorani for constructive comments. T.K. acknowledges support from a JSPS (Japan Society for the Promotion of Science) fellowship (09-7408). This research was supported in part by the Grants-in-Aid for the Center-of-Excellence (COE) Research of the Ministry of Education, Science, Sports and Culture of Japan (07CE2002) to RESCEU (Research Center for the Early Universe). \\bigskip \\bigskip \\baselineskip12pt \\parskip2pt" + }, + "9708/astro-ph9708041_arXiv.txt": { + "abstract": "Anomalous transport processes in which the variance of the distance travelled does not necessarily increase linearly with time: $\\left<\\Delta x^2\\right>\\propto t^{\\alpha}$ with $0<\\alpha<2$ are modelled using the formalism of continuous time random walks. We compute particle propagators which have the required dependence on space and time and use these to find the spatial dependence of the synchrotron radiation emitted by a population of continuously injected electrons. As the electrons are transported away from the source they cool, and the synchrotron spectrum softens. Sub-diffusive ($\\alpha<1$) transport -- corresponding to stochastic trapping, or restriction of the transport across the average direction of a stochastic magnetic field -- produces a much slower rate of change of spectral index than does supra-diffusion ($1<\\alpha<2$) -- which occurs when particles move almost without scattering, in a field containing large ordered regions. Application to the diffuse emission of the outer parts of the Coma cluster favours an interpretation in terms of supra-diffusion. ", + "introduction": "In a regular magnetic field topology, the transport of charged particles across the field is due to the collisions of the particles and their finite Larmor radii. However, a perturbation of the magnetic field results in a wandering of the field lines and a potentially much faster transport of the particles. This transport is widely attributed to the effect of the large-scale random field component, which would produce, following the quasilinear theory, a diffusion of the field lines across the direction of the the average field (Jokipii \\& Parker~\\cite{jokipiiparker69}, Kadomtsev \\& Pogutse~\\cite{kadomtsevpogutse79}). In fusion applications, it is generally assumed that particle collisions lead to diffusive transport along each individual field line. In astrophysical applications, where collisions can be neglected, it is usual to assume that small scale fluctuations in the magnetic field play this role, so that here too, particles diffuse along field lines (Chuvilgin \\& Ptuskin~\\cite{chuvilginptuskin93}). As long as the particles remain correlated to a given patch of field lines, the combination of these two diffusions results in sub-diffusion of the particles i.e., $\\langle \\Delta x^2(t) \\rangle \\propto t^{\\alpha}$, with $\\alpha = 1/2$, as described by Getmantsev~(\\cite{getmantsev63}). (In the general case, sub-diffusion refers to all such processes when $0 < \\alpha < 1$.) However, the exponential divergence of neighbouring field lines and the resulting stretching of a field line patch lead to decorrelation of a particle from its field line, and thus to the large-scale diffusion of the particles, as pointed out by Rechester and Rosenbluth~(\\cite{rechesterrosenbluth78}), with a diffusion coefficient given by the expression: \\eqb D_{RR} &=& D_{st} 2 \\kpar/l_{c\\delta} \\eqe where $l_{c\\delta} = l_c \\log(1/k_0 \\delta)$, $\\delta = l_c \\sqrt{\\kperp/\\kpar}$, where $l_c$ is the exponentiation length of the field lines. Here $D_{st}$ is the quasilinear diffusion coefficient of the field lines, which has the dimensions of a length, $k_0$ is a characteristic wave number of the magnetic turbulence, $\\kpar$ and $\\kperp$ are the quasilinear diffusion coefficients of the particles along and across the direction of the local magnetic field, respectively. The transport of particles across a turbulent magnetic field is thus sub-diffusive for short timescales, but crosses over to normal diffusion in the limit of the long timescales. In the following, we investigate the anomalous transport regime, which applies if the natural timescale of a particular problem -- such as escape from the galaxy, acceleration at a shock, or loss of energy by synchrotron radiation, is shorter than the decorrelation time. In a statistically homogeneous medium, the density of particles, $n(\\vec{x},t)$, is related to the source $Q(\\vec{x}',t')$ by the propagator $P$: \\eqb n(\\vec{x},t) &=& \\int \\diff\\vec{x}' \\int \\diff t' \\, P(\\vec{x}-\\vec{x}',t-t') Q(\\vec{x}',t') \\; . \\label{density}\\eqe For diffusive transport, $P$ typically contains a factor $\\exp(-|\\vec{x}|^2/Dt)$, and is the Green's function of the diffusion equation. For sub-diffusive transport with $\\alpha=1/2$, Rax \\& White (\\cite{raxwhite92}) have determined the propagator by combining two diffusive propagators using a Wiener integration method. They considered the transport in cylindrical symmetry across the $z$ direction (i.e., a 2-dimensional problem, with $|z|$ playing the role of time), and obtained: \\eqb P(r,t) &=& \\frac{H(t)}{2\\pi\\left( 4 r^{2} \\bdiff^{2} \\kpar t\\right)^{1/3}} \\exp\\left[ \\frac{-3 r^{4/3}} {4 \\left(4 \\bdiff^{2} \\kpar t\\right)^{1/3}} \\right]\\, \\label{prop2} \\eqe where $r$ is the radius in cylindrical coordinates, and $H(t)$ the Heaviside function. Duffy et al.~(\\cite{duffyetal95}) considered the transport perpendicular to a plane shock front. In a cartesian coordinate system with the shock in the $y$--$z$ plane, and assuming no gradients in the $y$ direction are present, they obtained the propagator \\eqb P(x,t) &=& \\int \\diff z \\, \\frac{\\exp\\left[-x^2 / \\left(4 \\bdiff |z|\\right) -z^2 / \\left(4 \\kpar t\\right)\\right]} {\\sqrt{4 \\pi \\bdiff |z|}\\sqrt{4 \\pi \\kpar t}} \\label{doublediff} \\eqe Approximating the integral using the method of steepest descents, they found \\eqb P(x,t)&\\approx& \\eta (\\bdiff \\kpar^{1/2} |x| t^{1/2})^{-1/3} \\exp \\left( \\frac{-\\beta |x|^{4/3}}{\\bdiff^{2/3}\\kpar^{1/3} t^{1/3}} \\right) \\label{prop1} \\eqe with $\\eta = 2^{1/3} (3\\pi)^{1/2}$ and $\\beta = 3/2^{8/3}$. For this 1-dimensional problem of $\\alpha=1/2$ sub-diffusion, Chuvilgin \\& Ptuskin~(\\cite{chuvilginptuskin93}) derived an equation describing the evolution of the particle density (their Eq.~B.12) and solved it to find the above propagator. Essentially the same equation was also found by Balescu~(\\cite{balescu95}). It can be written: \\eqb \\partial_t n(x,t) &=& D_0 \\Bigg[\\nabla_x^2 n(x,t) \\nonumber\\\\ && -{1\\over2\\sqrt{\\pi}\\tau_D} \\int_{1/\\pi}^t d\\tau \\left(\\frac{\\tau_D}{\\tau}\\right)^{3/2} \\nabla_x^2 n(x,t-\\tau) \\Bigg] \\; , \\label{nmde} \\eqe with $ D_0 = \\sigma^2/2 \\tau_D$. This is a non-Markovian diffusion equation: the integral term is characteristic of the long-time memory of the dynamics. In a realistic situation, the topology of the magnetic field may be more complicated than just a pure stochastic sea with its associated diffusion of the field lines. In fusion plasmas, for example, there can exist ordered structures, \\lq stability islands\\rq, in which particles can be trapped for long periods of time, (referred to as \\lq sticking\\rq), leading to sub-diffusive large-scale transport of the field lines themselves (e.g., White et al.~\\cite{whiteetal93}). On the other hand, the field lines may also wander faster than implied by diffusion. In this case, the field lines are said to perform \\lq flights\\rq, during which they maintain almost the same direction for a relatively long distance. Ultimately, the large-scale transport of field lines is the result of competition between \\lq sticking\\rq\\ and \\lq flights\\rq, and can yield transport regimes of many different kinds, ranging from slow sub-diffusion (almost perfect sticking) to fast supra-diffusion (domination of flights). In terms of $\\alpha$, the physically relevant range is $0<\\alpha <2$, with $\\alpha=2$ corresponding to transport completely dominated by a single straight-line flight. To find the particle propagator, it is necessary to combine the propagator for field lines with that for particle motion along the field, as described above for the case of $\\alpha={1\\over2}$. Here too, the assumption of diffusive transport can be generalised. Particles may, in fact, undergo no scattering at all, in which case they propagate ballistically along the field lines, as assumed by Achterberg \\& Ball~(\\cite{achterbergball94}). On the other hand, they may be trapped between magnetic mirrors on a segment of a field line. Once again, the transport can be characterised by an index $\\alpha$ which lies between 0 and 2. However, as we show below, the combination of two propagators does not extend the overall range of $\\alpha$ which is permitted. In fact, the effect of superimposing transport with the two indices $\\alpha_1$ and $\\alpha_2$ is described by a single value $\\alpha=\\alpha_1\\alpha_2/2$. ", + "conclusions": "In this paper we have presented a simple one-parameter method of modelling the effects of anomalous transport on energetic electrons. Anomalous transport is likely to occur wherever magnetised electrons (i.e., electrons which are tied to field lines) move in a magnetic field which has a stochastic component. In astrophysics, this is the rule rather than the exception and anomalous transport can be expected, for example, in the interstellar medium, as well as in the intra-cluster medium of clusters of galaxies. The primary observational diagnostic is the intensity and spectrum of the synchrotron radiation emitted by the transported particles. For the simplest case in which the magnetic field is of constant magnitude and of random orientation, we have presented general expressions for the surface brightness as a function from position of injection, and also for the spatial variation of the spectral index. Applying these to the diffuse emission observed from the outer parts of the Coma cluster, we note that it is not possible to distinguish between the various forms of transport merely from the profile of the surface brightness. However, the expected spatial dependence of the spectral index is very sensitive to the type of transport. Assuming that the effects of particle acceleration are negligible in the outer parts of the cluster, and that the electron distribution has achieved a steady state, we find that standard diffusive transport cannot produce the observed rapid softening of the spectrum with radius. Under these assumptions, the type of transport indicated is supra-diffusion, in which particles move almost ballistically in a field configuration which has an ordered radial component. The computations we have presented contain several major simplifications. In addition to the assumption of constant, randomly orientated magnetic field, and the simple planar or spherical geometry, we have assumed that the parameters governing the transport are independent of the particle's energy. In reality, the type of transport itself (i.e., the value of $\\alpha$) will change depending on the energy range and timescales considered. Thus, at very large times (which may exceed the synchrotron lifetime), a particle can be expected to decorrelate from the magnetic field and perform diffusion (e.g., Duffy et al.,~\\cite{duffyetal95}). We do not model the situation in which a significant change in $\\alpha$ occurs within the synchrotron lifetime of an electron. Finally, in order to apply such models to well-observed objects such as spiral galaxies, it will be necessary to include additional effects such as that of a galactic wind, as well as bremsstrahlung and ionisation loss processes. However, these processes will not change our basic conclusion that it is the spatial dependence of the synchrotron spectral index which provides the most sensitive measure of the transport properties of the emitting electrons. \\noindent{\\bf Acknowledgements:} BRR thanks the Max-Planck-Institut f\\\"ur Kernphysik for the grant of a visitor's stipend, during which this work was performed. We are grateful to R.O.~Dendy, P.~Duffy and Y.A.~Gallant for stimulating discussions. \\appendix" + }, + "9708/astro-ph9708107_arXiv.txt": { + "abstract": "The efficiency and uniqueness of the diffusive shock acceleration is studied on the basis of the novel kinetic solutions. These solutions obtained earlier (paper I) selfconsistently describe a strong coupling of cosmic rays with the gas flow. They show that the dependence of the acceleration efficiency upon physical parameters is critical in nature. In this paper we investigate a steady acceleration in the parameter space formed by the injection rate \\( \\nu \\), the upper cut-off momentum \\( p_1 \\) and the Mach number \\( M \\) while the flow compression $R $ serves as an order parameter. We determine a manifold of all possible solutions in this parameter space. To elucidate the differences between the present kinetic results and the well known two-fluid predictions we particularly focus on the \\(\\nu \\to 0, \\, M \\to \\infty \\) limit where the two-fluid model suffers from especially serious closure problems and displays an `unphysical' behavior. We show that in contrast to the two-fluid model three different solutions occurs also for arbitrarily large \\( M \\) provided that \\( p_1 \\) is sufficiently high. The three solutions appear together only if the injection rate $\\nu $ lies between two critical values, \\( \\nu_1 < \\nu < \\nu_2\\). For \\( \\nu < \\nu_1(M,p_1) \\) only the inefficient solution is possible. For \\( \\nu > \\nu_2(M,p_1) \\) only the efficient solution with a very high cosmic ray production rate occurs. On the basis of the obtained bifurcation surface \\( R(\\nu, p_1) \\) we consider the limit \\( p_1 \\to \\infty , \\, \\nu \\to 0 \\) which completely uncovers the long debating anomalies of the two-fluid model. The constructed steady state manifold that, at least partially is an attractor of a time dependent system, allows us to speculate on the nonstationary acceleration. ", + "introduction": "The question, how efficient the diffusive shock acceleration may be, arose naturally when the first test particle calculations of this process became available (Krimsky 1977; Axford, Leer \\& Skadron 1977; Bell 1978; Blandford \\& Ostriker 1978; see \\cite{dru}; \\cite{bla:eich} and \\cite{jel91} for a review). This is because the backreaction of accelerated particles (cosmic rays (CRs) in the astrophysical context) on the shock structure is very strong and leads normally to a significant increase of the compression ratio. Such an accelerating shock should therefore be thought of as a strongly nonlinear dynamical system with a pronounced selforganization. Neither the particle spectrum nor the hydrodynamic flow structure can be calculated independently. Furthermore, since the diffusion length of particles increases with momentum, particles with higher momenta sample longer parts of the shock transition. This makes any kind of moment description very difficult. However, the first essentially nonlinear calculations of this acceleration process were performed within the hydrodynamic approach. \\subsection{Success and limitations of the two-fluid model} The above arguments suggest that the problem is kinetic in nature, which makes the usage of any fluid theory for describing the acceleration process questionable. At the same time quite a deep insight can be gained from simple moment equations. The two-fluid model (TFM) introduced by Axford Leer \\& Skadron (1977) and Drury \\& V\\\"olk (1981) (DV, hereafter) treats the thermal and CR populations as separate fluids coupled only through the hydrodynamic equations. The main effect of this coupling is a deceleration of the inflowing gas in front of the shock by the pressure gradient of counterstreaming CRs accelerated at the shock and, as a result, an enhancement of the total shock compression, the multiplicity of solutions and a much higher acceleration efficiency. Unfortunately, the underlying particle distribution implies the pressure divergence and is underdetermined in some other ways (see \\eg \\cite{dru}; \\cite{abp}; \\cite{kj90}, and below). \\subsection{Renormalization of the two-fluid model} A renormalization procedure to overcome the above ultraviolet divergence has been suggested recently by Malkov \\& V\\\"olk (1996) (MV96, hereafter). This theory produces basically the same two-fluid hydrodynamics except the renormalized CR specific heat ratio $\\Gamma$ instead of the usual \\( \\gamma_{\\rm c}\\) which results from the losses. Under the assumption $\\gamma_{\\rm c} =5/3$ upstream (which automatically implies that far upstream $\\Gamma=5/3$ as well, for the case of reacceleration considered in MV96), $\\gamma_{\\rm c} =4/3$ downstream, and in the limit $p_1 \\to \\infty$, the renormalized two-fluid model (RTFM) produces a solution qualitatively similar to that of the unrenormalized theory. In the case of $\\gamma_{\\rm c} =4/3$ upstream, when also \\(\\Gamma \\) decreases $\\Gamma < \\gamma_{\\rm c}$, the results change dramatically. Namely, the shock compression $r$ becomes much larger than the usual unrenormalized result $r=7$ (for strong shocks). What happens is a very fast increase of the losses with decrease of $\\gamma_{\\rm c}$ which rises the compression that even tends to infinity (for infinitely large Mach numbers when, in addition, one takes the maximum possible spectral slope $q =3r/(r-1)$ at the upper cut-off). This regime was not (and could not be) explored in MV96 since such a high compression shock requires a detailed information about $\\gamma_{\\rm c},\\, \\Gamma $, and $\\bar\\kappa$, the spectrum averaged CR-diffusivity across the shock transition. This would practically be equivalent to the full kinetic solution. To obtain such a solution was one of the main objectives of a companion paper (Malkov 1997a, paper I). Further motivations of the present paper will be outlined in the next subsection. To conclude this subsection we note that the assumption \\( \\gamma_{\\rm c} \\approx 4/3 \\) upstream is precisely what the kinetic solution obtained in paper I suggests for the case of injection in contrast to the case of reacceleration, \\( \\gamma_{\\rm c} = 5/3 \\) considered in MV96. Moreover, $\\gamma_{\\rm c} $ is smaller upstream than downstream, again opposite to the reacceleration case. Finally, the \\rtfm results in the injection case are unacceptably sensitive to the values of $\\gamma_{\\rm c} $ and \\( \\Gamma \\) that are not known to the required extent when the kinetic solution is not available. This makes the moment approach especially restrictive for describing namely the injection triggered acceleration process. As in paper I we concentrate here on this, certainly more interesting and at the same time more difficult case. \\subsection{`Pathological' limits of the \\tfm} There exists another difficulty of the \\tfm that has already been criticized in the literature (\\eg \\cite{jel91}). Namely, the shock modification in a steady state occurs while particles are constantly injected at some rate, are then accelerated, and disappear eventually through the upper cut-off or downstream. Once the injection is somehow eliminated, the CR-dominated (or efficient) steady-state solution cannot be justified physically, and the ordinary gas shock remains the only solution possible. The TFM, however, still produces a CR-dominated shock that even becomes completely smooth beyond a certain Mach number, \\( M > M_1 \\). In fact, it represents a fast-mode shock in a two-fluid hydrodynamics associated with the tenuous high pressure CR-fluid (\\cite{ptus}). Since the number density of the CRs \\( n_{\\rm c} \\) is irrelevant in the TFM, solutions that have a finite CR pressure (\\( P_{\\rm c} > 0\\)) are formally permitted without any injection (\\( n_{\\rm c} =0\\)). Moreover, for an arbitrarily small nonzero injection rate there exists a critical Mach number \\( M_2 < \\infty \\) above which this efficient solution is the only one the \\tfm can offer. Clearly, it is difficult to judge the acceleration efficiency because the injection rate is, as a rule, very small and the two-fluid system, on the other hand, does not behave adequately when injection vanishes. The question, however, is whether the consideration of this limit within the \\tfm is admissible. The answer is definitive not. Indeed, a fully kinetic steady state solution obtained in paper I revealed the following nonlinear response of the system to a weak injection of thermal particles. First of all, this response depends not so much on the injection itself but rather on the parameter \\( \\Lambda_1=\\nu/\\delta \\equiv \\nu p_1/p_0 \\), where \\( \\nu \\) and \\( p_0 \\) are the injection rate and injection momentum, respectively, and \\( \\delta \\ll 1 \\). Furthermore, the effect of shock modification completely disappears as \\( \\Lambda_1 \\to 0 \\) (it becomes practically insignificant in an abrupt manner already at \\( \\nu /\\sqrt{\\delta} \\la 1 \\)). Only in the other extreme, \\( \\Lambda_1 \\to \\infty \\) (more precisely, when \\( \\Lambda_1 \\gg M^{3/4} \\), see also subsection \\ref{crit:inj}) a {\\em unique} solution that is indeed injection insensitive appears which is quite in the spirit of the TFM. Quantitatively, this solution can be very different from the respective \\tfm solution but for a different reason which is related to the subshock smoothing. The TFM is of course incapable of describing the dependence of its solution on \\( \\Lambda_1 \\) since it implies \\( p_1=\\infty \\), i.e. \\( \\Lambda_1 =\\infty \\) already on the derivation level; even if \\( \\nu \\) is set to zero afterwards the physically correct behavior of the solution with vanishing injection cannot be recovered since \\( \\Lambda_1 \\) remains infinite. Therefore it is useless to expect from the \\tfm a correct behavior at \\( \\nu \\to 0 \\) and to criticize this model for the lack of it. This is beyond its validity range. The most dramatic consequence of this (\\( p_1 \\to \\infty \\)) degeneracy is the subshock smoothing (\\( r_{\\rm s} \\equiv 1 \\)). The solution becomes enormously different from the kinetic solution that does not pass through the point \\( r_{\\rm s} =1 \\) just because of this fact. Why this is so, has been explained in paper I and we shall look at this problem from a different perspective in subsection \\ref{crit:inj}. Perhaps the most direct explanation why the kinetic solution differs so strongly from its hydrodynamic counterpart is the singular character of the underlying perturbation problem in the small parameter \\(\\delta \\equiv p_0/p_1 \\ll 1 \\). No matter how small it is, the efficient kinetic solution with \\( \\delta = 0 \\) (and, hence, the \\tfm solution) is fundamentally different from that with \\( \\delta > 0 \\) (see paper I). The parameters that the \\tfm usually operates on are the Mach number \\( M \\) and the injection rate \\( \\nu \\) (or the seed particle pressure in the case of reacceleration). As we argued, this is not enough to describe consistently what is going on in the steady nonlinear shock acceleration. The \\rtfm introduces an additional parameter needed, \\(p_1 \\), and accounts of the losses at \\( p = p_1 \\) but then it lets \\( p_1 \\to \\infty \\), \\( \\nu \\) fixed. That means \\( \\Lambda_1 \\equiv \\nu /\\delta \\to \\infty \\), and therefore the results are again insensitive to \\( \\nu \\) when \\( \\nu \\to 0 \\) since the critical information about \\( \\Lambda_1 \\) is lost, exactly as in the TFM. We emphasize that it is the parameter \\( \\Lambda_1 \\) that regulates primarily the budget of energetic particles at a shock, not \\( \\nu \\) alone. That is why, for example, the subshock completely vanishes in the \\tfm as well as in the \\rtfm beyond a certain \\( M =M_1\\) even for \\( \\nu \\to 0 \\); the more important parameter \\( \\Lambda_1 \\) remains infinite which effectively corresponds to the situation with a very strong injection. If we allow for time dependence on the kinetic level of description and assume a slower than in the Bohm limit momentum dependence of the CR diffusion coefficient, completely smooth \\tfm stationary solutions will exist even with zero injection and with no seed particles upstream. This phenomenon has been explained by Drury (1983)-- the high energy particles may be considered as those injected in the past and being then continuously accelerated at the shock. Unfortunately the smooth \\tfm solutions can tell us very little about a steady acceleration of CRs out of the thermal upstream plasma -- if there are no preexisting CRs, there are no such solutions. Since these \\tfm solutions are essentially time dependent on the kinetic level (see \\eg \\cite{fg87}, \\cite{kj90} and \\cite{dvb} for relevant discussions), we may infer that once a natural cut-off \\( p_1 < \\infty\\) exists and is reached, this acceleration regime will be disrupted. Indeed, these solutions correspond to the acceleration of CRs injected in the past whose number density virtually decreases (although being irrelevant in the \\tfm which in fact admits these pseudo-steady solutions) while the CR pressure remains approximately constant, being determined simply by the ram pressure of the inflowing gas. When particles start to leak through the upper cut-off the CR pressure decreases as well and the system relaxes to the ordinary gas shock which is the only steady state solution without injection. In general, the above \\tfm quasi-stationary acceleration scenario implies a rather low production of CRs (in terms of their number, not the energy density) because it operates only on initially injected particles and suppress further injection as soon as this solution is set up. Much more productive would be solutions that allow for permanent losses. This would mean, in fact, the propagation of high-energy CRs into the shock surroundings which decouples them from the gas flow with their replenishment due to the permanent injection at the subshock. But these solutions are essentially kinetic and, as we emphasized, fairly different from the \\tfm solutions. As it was demonstrated in paper I, given the injection rate three different solutions are possible for sufficiently high \\( M \\) and \\( p_1 \\), Figure 1. However, only the most efficient solution with the highest compression ratio has been considered in detail. Accordingly, only a first critical injection, i.e. the injection rate \\( \\nu = \\nu_1 \\) above which this solution exists along with the two other solutions has been calculated. The calculation of a second critical injection that requires an inspection of the two remaining solutions and above which the efficient solution is the only possible, is one of the subjects of the present paper. We consider the entire manifold of stationary solutions and in this context the three above-mentioned solutions are simply its subsets. In the next section we briefly review the physical formulation of the problem and discuss our strategy of a unified description of all the three solutions. In Sec.3 we obtain both the efficient and inefficient solutions from the integral equation derived in paper I and consider their matching in an intermediate range. In Sec.4 we describe the solution space as a whole and calculate the critical injections. We conclude this section with implications of its results for the \\tfm. Further, in Sec.5 we speculate upon possible scenarios of time dependent acceleration on the basis of the emerged bifurcation picture. Sec. 6 discusses the results and some of their most evident consequences for calculations of the acceleration efficiency in real astrophysical objects. ", + "conclusions": "\\label{disc} The early works on the nonlinearly modified CR shocks inspired the hope that, because of a very high acceleration efficiency in the nonlinear regime, the overall CR production should not be very sensitive to injection and just a qualitative understanding of this complicated process suffices for quantitative calculations of acceleration efficiency. The CR dominated (efficient) solutions of the two-fluid model (DV) strongly supported this idea. In fact, such an optimism rests on the limited amount of energy available in the gas flow to be converted into CRs. Indeed, if we consider a strong shock ($ M \\gg 1 $) we may calculate the acceleration efficiency, or the coefficient of the flow energy conversion, as $\\varepsilon_{\\rm conv} = P_{\\rm c}(0)/\\rho_1 u_1^2=1-1/R $, eq.(\\ref{ber}). Whenever the acceleration process is known to be in the nonlinear regime ($R \\gg 1$), almost all the flow energy goes into CRs, practically independent of anything at all. The main issue now is, under which circumstances the system may indeed be in a highly nonlinear acceleration regime. As we have seen, the answer to {\\em this} question depends {\\em critically } on the injection rate: infinitesimal variations of $\\nu $ in the vicinity of $\\nu_1 $ or $\\nu_2 $ can result in finite (and typically very large) variations of \\( R \\) due to transitions between different branches. Also solutions belonging to the same branch are typically very sensitive to the injection rate \\( \\nu \\). This may be easily understood from the inspection of \\eg Figure 2 and Figure 3a where only on the inefficient branch \\( R \\) varies relatively slowly with \\( \\nu \\), while in the cases of intermediate and especially efficient solution \\( R \\) changes very rapidly with injection. There is of course an injection insensitive region belonging to efficient solution where the compression $R(\\nu ) $ approaches its upper bound \\( R \\sim M^{3/4} \\) at given $M$ (see Figure 3a, where this region may be identified as a very sharp growth of the function \\( \\nu(R,M) \\) in the farthermost corner of the plot). Such a behavior is caused by the requirement of a finite subshock. It should be noted that this scaling has indeed been observed in some numerical works with fixed injection rates (\\eg \\cite{bky}). However, this situation occurs at sufficiently large values of \\(\\nu \\) and diminished \\( r_{\\rm s} \\) and it is doubtful that such a high injection rate is possible at a weak subshock. For smaller \\( \\nu \\), when \\( \\nu \\ll \\delta \\cdot M^{3/4} \\) (the system parameter \\( \\Lambda \\ll 1 \\)) an important signature of the stationary acceleration is that the compression ratio is practically independent of \\( M \\) (see paper I for further details). Physically, the injection rate should be calculated selfconsistently using the solution of injection problem given subshock parameters. The solution of this problem provides a function \\( \\nu = \\nu_{\\rm s }(R,M) \\). Then isolated solutions for \\( R = R(M) \\) might be obtained as intersection points of the curves \\( \\nu_{\\rm s} \\) and \\( \\nu \\) (as shown in Figure 1). This solution may or may not be multiple depending on the character of the function \\( \\nu_{\\rm s} \\). In any case, the bifurcation diagrams alone do not suffice for determining the actual acceleration efficiency. The calculation of the injection $\\nu_{\\rm s} $ as a function of shock characteristics is equally important for this purpose. It should also be born in mind that the model considered here and in paper I gives an upper bounds to the actual acceleration efficiency. A number of not included factors may significantly decrease compression ratio \\( R \\), acceleration efficiency and the spectrum hardness (see paper I for a relevant discussion). Turning to the time dependent acceleration we note that unless the scenario suggested in the previous section is totally unrealistic, a critical quantity that would determine the CR production is the cut-off momentum \\( p_1^{(2)} \\) beyond which the system jumps to the efficient acceleration regime. According to eq.({\\ref{nu_2}) \\(p_1^{(2)} \\propto \\exp (1/\\nu) \\). Since in the Bohm limit we may write \\( p_1(t) \\propto t \\), the corresponding critical time \\( t_{\\rm crit} \\propto \\exp (1/\\nu) \\). In an accelerating object of a finite life time \\( \\tau \\) (\\eg supernova remnant, SNR) the main question, of course, is whether the condition \\( \\tau > t_{\\rm crit} \\) is fulfilled. This is again extremely sensitive to the injection rate. Theoretically, the injection rate \\( \\nu_{\\rm s} \\) must not necessarily be as high as \\( \\nu_{2} \\) for the acceleration process to become efficient. The condition \\( \\nu_{\\rm s } > \\nu_1 \\) could suffice provided that the lower branch \\( R_l \\) (Figure 6) looses stability for \\( p_1 \\la p_1^{(1)} \\). Since for large values of \\( p_1 \\), \\( \\nu_2 \\ga 10 \\nu_1 \\) (see Figure 5), this may determine the outcome of the acceleration process completely. Another possibility to overcome the high \\( \\nu_2 \\) threshold is an essentially stronger time dependence of the acceleration process than that discussed in the preceding section. Basically, the bifurcation picture of this system is quite rich and promises an interesting dynamics. This is the more so as governing parameters are themselves subject for a temporal evolution. They may change significantly during the acceleration process in a variety of astrophysical environments. These variations may be of a quasi-external type like \\eg decrease of the Mach number when the shock slows down. Equally important may be an intrinsic variability associated with the growth of the maximum energy or with the heating of the upstream plasma by the CR driven turbulence (\\cite{VDMcK}). Therefore, to comprehend the acceleration dynamics we must face the injection problem together with the physics of subshock dissipation and treat these problems selfconsistently with the above bifurcation analysis. The fact that this system displays very much hysteresis emphasizes the necessity of this approach. It is important to recognize that there is a serious drawback in the way to a full calculation of the acceleration efficiency in concrete astrophysical shocks, like \\eg SNR shocks. It originates from the threshold nature of the acceleration process. Indeed, the flow structure changes quasi-abruptly when the critical injection $\\nu_2 $ drops below $\\nu_{\\rm s}(R_2) $ as \\( p_1 \\) grows (or $\\nu_{\\rm s}(R_1) $ becomes smaller than $\\nu_1 $). Since the function $\\nu_{\\rm s} $ is very sensitive to local subshock conditions (a local orientation of the magnetic field is perhaps the most obvious and very important factor here), this transition occurs first at those parts of the shock surface where $\\nu_{\\rm s} $ reaches its maximum. This must result in `hot spots' or `discharge' zones in the shock front where the acceleration becomes efficient. Then, the flow structure will be essentially 3-dimensional (or at least quasi 2-dimensional), quite complicated and probably unsteady. The inhomogeneity of the ambient medium (see \\eg \\cite{mcKee}) may very well result in a similar effect. Clearly, the one-dimensional calculations, even with a properly determined injection rate \\( \\nu_{\\rm s}(R,M) \\) may give at best an upper bound to the acceleration efficiency. Even if the flow remains quasi-laminar the overall efficiency will be reduced according to the surface density of the hot spots on the shock front. Besides that the losses from the hot spots into the neighboring regions of inefficient acceleration may significantly reduce the maximum energy. A very important consequence of this would be the corresponding increase of the critical injection, eq. (\\ref{nu_1}) which may drive the system below the threshold of the efficient acceleration." + }, + "9708/astro-ph9708113_arXiv.txt": { + "abstract": "We have analyzed the spectral variations of the superluminal black-hole X-ray binary GRS~1915+105 by using data obtained with the PCA on the Rossi XTE. We find that, despite the marked differences in the structure and the time scale of variability, all spectral changes can be attributed to the rapid disappearing of the inner region of an accretion disk, followed by a slower re-filling of the emptied region. The time scale for each event is determined by the extent of the missing part of the disk. The observed relation between the duration of an event and the radius of the disappearing region matches remarkably well the expected radius dependence of the viscous time scale for the radiation-pressure dominated region of an accretion disk. ", + "introduction": "The X--ray source GRS~1915+105 is the best of the two only examples of galactic objects that show superluminal expansion in radio observations (\\cite{mr94}). It is located at a distance of 12.5 kpc and its relativistic jets are directed at 70$^\\circ$ from the line of sight (\\cite{mr94}). After its discovery as an X-ray transient in 1992 (\\cite{cas92}), a number of outbursts has been reported, although it is possible that the source has never completely switched off in between. The Rossi X-ray Timing Explorer (RXTE) started observing GRS~1915+105 in a bright and variable state in April 1996 and continued to observe it regularly at least once per week since then. During this period, the source has displayed a remarkable richness in variability, ranging from quasi-periodic burst-like events, deep regular dips and strong quasi-periodic oscillations, alternated with quiescent periods (\\cite{gmr96},\\cite{mrg96},\\cite{cst97},\\cite{bmk97},\\cite{tcs97}). Because of its similarities with the other galactic superluminal source GRO~1655-40 (\\cite{zh94}), whose dynamical mass estimate implies the presence of a black hole (\\cite{bai95}), and because of its high luminosity exceeding the Eddington limit for a neutron star, the source is suspected to host a black hole. In a previous paper (\\cite{bmk97}, hereafter Paper I) we showed that the complicated light curve of GRS~1915+105 can be described by the rapid ($\\sim$1 s) appearance and disappearance of emission from an optically thick inner accretion disk. In this paper we report on the results of the analysis of a later observation which allowed a more detailed investigation of the spectral variability of the source. We find that the variation time scale $t_{\\rm rec}$ of the disk source is completely specified by the maximum size $R_{\\rm max}$ of it inner radius. The relation between these quantities is precisely as expected for the viscous evolution of a radiation--pressure dominated accretion disk. ", + "conclusions": "The results described in the previous section can be interpreted within the model presented in Paper I, providing a unified picture of the variability observed in GRS~1915. In Paper I, we modeled the large amplitude changes as emptying and replenishing of the inner accretion disk caused by a viscous-thermal instability. The small radius observed during the quiescent period was identified with the innermost stable orbit around the black hole, while the large radius during the burst phase was the radius of the emptied section of the disk. The smaller oscillations were interpreted as failed attempts to empty the inner disk. As it can be seen from Figure 1, from this observation we find that all variations, from major events like the ones described in Paper I to small oscillations observed at the end of a large event, can be modeled in exactly the same fashion (see Figure 1). In this scenario, the ``flare state'' presented in Paper I is simply a sequence of small events following a big one, similar to the small oscillations in Figure 1. Both the spectral evolution and the duration of the event are determined by one parameter only, namely the radius of the missing inner section of the accretion disk. For a large radius, the drop in flux will be large and the time needed for re-fill will be long. Following Paper I, it is natural to associate the length of the quiescent part of an event $t_q$ to the viscous time scale of the radiation-pressure dominated part of the accretion disk. This can be expressed as $t_{\\rm visc}=R^2/\\nu$, where $\\nu = \\alpha c_S H$. From Frank, King \\& Raine (1992) the scale-height $H$ and sound speed $c_S$ can be found, leading to \\begin{equation} t_{\\rm visc} = 30 \\alpha_2^{-1} M_1^{-1/2} R_7^{7/2}\\dot M_{18}^{-2}\\ {\\rm s} \\end{equation} where $\\alpha_2 = \\alpha/0.01$, $R_7$ is the radius in units of $10^7$ cm, $M_1$ the central object mass in solar masses, and $\\dot M_{18}$ is the accretion rate in units of $10^{18}$ g/s. Notice that even the largest radii derived here are well within the radiation-pressure dominated part of the disk (see Equation 2 in Paper I). The line in Figure 4 represents the best fit to the data with a relation of the form $t_{\\rm q} \\propto R^{7/2}$. The fit is excellent, with the exception of the point corresponding to the longest event. The qualitative agreement with the theoretical expectation is striking, although by substituting the appropriate values for the mass and accretion rate we find that our best fit predicts rather small values of $\\alpha_2$ (0.004 and 0.05 for the Schwarzschild and extreme Kerr cases respectively). This indicates a small viscosity in the disk, although we stress that $t_{visc}$ is only a time scale, so that additional corrections might be necessary in order to allow a precise quantitative comparison. An event can therefore be pictured in the following way. At the start of a quiescent period, the disk has a central hole, whose radius is R$_{max}$. The hole is either empty or filled with gas whose radiation is too soft to be detected. Slowly the disk is re-filled by a steady accretion rate $\\dot M_0$ from outside. Each annulus of the disk will move along the lower branch of its S-curve in the $\\dot M-\\Sigma$ plane trying to stabilize at $\\dot M_0$ (see Paper I). The surface gravity increases as the annulus moves towards the unstable point at a speed determined by the local viscous time scale. During this period, no changes are observed in the radius of the hole, since all the matter inside does not radiate in the PCA band. The observed accretion rate is $\\dot M_0$. At some point, one of the annuli will reach the unstable point and switch to the high-$\\dot M$ state, where the accretion rate is larger than $\\dot M_0$, causing a chain-reaction that will ``switch on'' the inner disk. The observed accretion rate is now higher than the external value $\\dot M_0$. A smaller, hot radius is now observed. At the end of the outburst, the inner disk runs out of fuel and switches off, either jumping back to the $\\dot M < \\dot M_0$ state or emptying completely. A new hole is formed and a new cycle starts. Notice that in this scenario the more ``normal'' state for the source is the one at high count rates, where the disk extends all the way to the innermost stable orbit: in this state the energy spectrum is similar to that of conventional black-hole candidates (see \\cite{tl95}). Not only the start and end of a major burst, but also all the small amplitude oscillations within a burst show the same timing signature of decaying faster than they rise. This is in agreement with what was already noticed in Paper I: the rise time is determined by the speed at which a heating wave moves through the central disk, while the faster decay time is due to the rapid fall of matter into the black hole (or into a relativistic jet). Chen, Swank \\& Taam (1997) found that when HR2 exceeds 0.1, the power density spectrum is similar to the one observed in black-hole transient systems during the Very High State (see \\cite{vdk95}). In these occasions, a strong 1-6 Hz QPO peak was found, positively correlated with the count rate. Notice that the limit HR2$>$0.1 is an indication that the source was in a quiescent state. Our spectral results show that the fast timing features (both QPO and band-limited noise, see \\cite{mrg96}) cannot originate from the innermost regions of the optically thick accretion disk, since those are missing during the quiescent phases. The fact that the QPO frequency increases with count rate is in qualitative agreement with the model, since a higher count rate indicates a smaller inner disk radius, and therefore shorter time scales. The radii for the disappearing region of the disk found here are considerably smaller than that reported in Paper I. This is entirely due to the improved knowledge of the spectral response of the PCA. During the burst phase the accretion rate through the optically thick accretion disk is found to be higher than in the quiescent phase, contrary to what we reported in Paper I. Although we can model the spectral evolution and the time scale of the events, most of the variability is yet to be explained. The main question to be answered is: what determines the length of the following outburst? Or, in more physical terms, what determines how large the next missing section of the disk will be? In some observations the events are very regular, in some others they are extremely random, and in some others no events are observed at all. In the observation reported here a striking one to one relation between quiescent and burst time is observed, a relation which applies to the ``outburst'' and ``quiescent'' states of the observation presented in Paper I but is obviously not satisfied in other observations (see e.g. \\cite{tcs97}) nor during the ``flare'' state of Paper I. Moreover, as already mentioned, one observation among the ones we analyzed does not fit this pattern and requires a different interpretation (1996 June 16th). Nevertheless, the model sketched here provides a satisfactory interpretation of the cause of the changes in the X-ray emission. It fits not only the energy spectral distribution, but also the complex variability shown by GRS~1915+105." + }, + "9708/astro-ph9708263_arXiv.txt": { + "abstract": "Starbursts are episodes of intense star-formation that occur in the central regions of galaxies, and dominate the integrated emission from the galaxy. They are a significant component of the present- day universe, being the site of $\\sim$ 25\\% of the high-mass star- formation. They offer unique `laboratories' for testing our ideas about star-formation, the evolution of high-mass stars, and the physics of the interstellar medium. They serve as local analogs of the processes that were important in the origin and early evolution of galaxies and in the heating and chemical enrichment of the inter-galactic medium. In this contribution I review starbursts from this broad cosmogonical perspective, stressing several key lessons we have learned from starbursts: 1) Violent, transient events play a significant role in the origin and evolution of galaxies. 2) Galaxies do not evolve as `Island Universes': starbursts are triggered by galaxy interactions and produce outflows of hot chemically-enriched gas that `pollute' the inter- galactic medium. 3) Dust dramatically affects of view of high-mass star-formation in starbursts and (probably) in high-redshift galaxies. Throughout this review I emphasize the importance of space-based observations in understanding starbursts. ", + "introduction": " ", + "conclusions": "" + }, + "9708/gr-qc9708068_arXiv.txt": { + "abstract": "Diffeomorphism freedom induces a gauge dependence in the theory of spacetime perturbations. We derive a compact formula for gauge transformations of perturbations of arbitrary order. To this end, we develop the theory of Taylor expansions for one-parameter families (not necessarily groups) of diffeomorphisms. First, we introduce the notion of knight diffeomorphism, that generalises the usual concept of flow, and prove a Taylor's formula for the action of a knight on a general tensor field. Then, we show that any one-parameter family of diffeomorphisms can be approximated by a family of suitable knights. Since in perturbation theory the gauge freedom is given by a one-parameter family of diffeomorphisms, the expansion of knights is used to derive our transformation formula. The problem of gauge dependence is a purely kinematical one, therefore our treatment is valid not only in general relativity, but in any spacetime theory. ", + "introduction": "\\setcounter{equation}{0} In the theory of spacetime perturbations \\cite{bi:SW,bi:waldbook,bmms}, one usually deals with a family of spacetime models $M_\\lambda:=({\\cal M},\\{T_\\lambda\\})$, where $\\cal M$ is a manifold that accounts for the topological and differential properties of spacetime, and $\\{T_\\lambda\\}$ is a set of fields on $\\cal M$, representing its geometrical and physical content. The numerical parameter $\\lambda$ that labels the various members of the family gives an indication of the `size' of the perturbations, regarded as deviations of $M_\\lambda$ from a background model $M_0$. Perturbations are described as additional fields in the background, defined as $\\Delta T^\\varphi_\\lambda :=\\varphi^*_\\lambda T_\\lambda -T_0$, where $\\varphi_\\lambda : {\\cal M}\\to {\\cal M}$ is a diffeomorphism that provides a pairwise identification between points of the perturbed spacetime and of the background, and $\\varphi_\\lambda^*$ denotes the pull-back. Of course, such an identification is arbitrary, and this leads to a gauge freedom in the definition of perturbations. Under a change $\\varphi_\\lambda\\to \\psi_\\lambda$ of the point identification mapping, a perturbation transforms as $\\Delta T^\\varphi_\\lambda \\to \\Delta T^\\psi_\\lambda$, with \\beq \\Delta T^\\psi_\\lambda=\\Phi_\\lambda^* \\Delta T^\\varphi_\\lambda + \\left(\\Phi^*_\\lambda T_0-T_0\\right)\\;, \\lab{gauge}\\eeq where $\\Phi_\\lambda:=\\varphi_\\lambda^{-1}\\circ\\psi_\\lambda$ is a diffeomorphism on $\\cal M$. In the perturbative approach, one tries to approximate $T_\\lambda$ expressing $\\Delta T^\\varphi_\\lambda$ as a series, \\beq \\Delta T^\\varphi_\\lambda=\\sum_{k=1}^{n-1}{\\lambda^k\\over k!}\\,\\delta^k T^\\varphi+O(\\lambda^n)\\;, \\lab{exp}\\eeq where $n$ is the order of differentiability with respect to $\\lambda$ of $\\Delta T^\\varphi_\\lambda$, and then solving iteratively the field equations for the various terms $\\delta^k T^\\varphi$. It is then important to know how the latter transform under a change of gauge. Until very recently, only the first order terms, $\\delta^1 T^\\varphi$, have been considered; in this case, it is well-known that the representations of a perturbation in two different gauges differ just by a Lie derivative of the background quantity $T_0$ \\cite{bi:SW}. However, non-linear perturbations are now becoming a valuable tool of investigation in black hole and gravitational wave physics \\cite{bi:gleiseretal}, as well as in cosmology \\cite{bi:cosmo}. Their behaviour under gauge transformations can be derived by Taylor-expanding (\\ref{gauge}) with respect to $\\lambda$. This apparently straightforward procedure presents a difficulty, though. Even if one chooses, as usual, point identification maps that are one-parameter groups with respect to $\\lambda$, the family of diffeomorphisms $\\Phi_\\lambda$ is {\\em not\\/} a one-parameter group \\cite{bmms}, i.e., it does not correspond to a flow on $\\cal M$. While flows on manifolds are well understood and widely discussed in the literature, more general one-parameter families of diffeomorphisms are not. Only some fragmentary statements about them can be found in a few papers \\cite{taub,schutz1,gl,fw}. Therefore, in order to extract from (\\ref{gauge}) the relationship between $\\delta^k T^\\varphi$ and $\\delta^k T^\\psi$, one must first develop the theory of Taylor expansions for general one-parameter {\\em families\\/} of diffeomorphisms, not necessarily forming a local group. The purpose of the present article is to provide the mathematical framework needed for this purpose. Roughly, the discussion generalises section 2 of reference \\cite{bmms} from the analytic to the $C^n$ case, but we also derive here a compact formula that gives directly the gauge transformation to an arbitrary order $k$. The paper is organised as follows. In the next section we define particular combinations of flows that we dub knight diffeomorphisms, and present our main result (Theorem \\ref{theorem1}). This establishes that arbitrary one-parameter families of diffeomorphisms can be approximated by families of knights, so that all one needs is a suitable expression for the Taylor expansion of knights, which is derived in section 3. Then, Theorem \\ref{theorem1} is proved in section 4. Section 5 contains the application to (\\ref{gauge}), i.e., our formula (\\ref{gtransf}) and some concluding remarks. In the following, we shall work on a finite-dimensional manifold $\\cal M$, smooth enough for all the statements below to make sense. In order to avoid cumbersome talking about neighbourhoods, we shall often suppose that maps are globally defined. This assumption simplifies the discussion, without altering the results significantly. Also, we specify the class of differentiability of an object only when it is really needed. Finally, let us recall that a one-parameter family of diffeomorphisms of $\\cal M$ is a differentiable mapping $\\Phi:{\\cal D}\\to{\\cal M}$, with $\\cal D$ an open subset of ${\\rm I\\!R}\\times{\\cal M}$ containing $\\{0\\}\\times{\\cal M}$, and $\\Phi(0,p)=p$, $\\forall p\\in{\\cal M}$. As we have already been doing, we shall write, following the common usage, $\\Phi_\\lambda(p):= \\Phi(\\lambda,p)$, for any $(\\lambda,p)\\in{\\cal D}$. ", + "conclusions": "\\setcounter{equation}{0} In the previous sections we have presented the theory of Taylor's expansions for one-parameter families of diffeomorphisms on a manifold $\\cal M$. Taking the simple case of a flow as our basic element, we have first defined the notion of knights, and then shown that an arbitrary one-parameter family of diffeomorphisms can always be approximated by a family of knights of a suitable rank. We can now return to the problem stated in the introduction, of finding the relationship between the $k$th order perturbations of a tensor $T_\\lambda$ in two gauges $\\varphi_\\lambda$ and $\\psi_\\lambda$. Let $n$ be the lowest order of differentiability of the objects contained in (\\ref{gauge}). It follows from Theorem \\ref{theorem1} that the action of $\\Phi_\\lambda$ is equivalent, up to the order $\\lambda^n$, with the one of a knight $\\Psi_\\lambda$, constructed as in (\\ref{theorem}). Therefore, we can expand (\\ref{gauge}) using (\\ref{lemma2}), and find, $\\forall k 0.3$) X-ray selected cluster sample was that of the EINSTEIN Extended Medium Sensitivity Survey (EMSS; Henry et al. \\markcite{h92}1992; Gioia \\& Luppino \\markcite{gl-94}1994). Henry et al. \\markcite{h92}(1992) used the EMSS to show X-ray clusters evolved `negatively' --- the space density of high luminosity clusters being lower in the redshift range $z = 0.30 - 0.60$ compared to $z = 0.14 - 0.20$. This result was in conflict with popular models of cluster formation (e.g. Kaiser \\markcite{k86}1986) and prompted further significant theoretical work (e.g. Kaiser \\markcite{k91}1991; Evrard \\& Henry \\markcite{eh91}1991). The maturing of the ROSAT database has sparked much recent interest in testing this result. A large cluster sample with a median depth of $z \\sim 0.1$, created from the ROSAT All Sky Survey (RASS), shows no sign of evolution out to $z = 0.3$ (Ebeling et al. \\markcite{e97}1997). Castander et al. \\markcite{RIXOS}(1995) presented the first look at the high redshift cluster population with ROSAT, claiming that the evolution seen by Henry et al. \\markcite{h92}(1992) extends to luminosities $\\sim 10^{44}$ erg s$^{-1}$. We have recently shown (Collins et al. \\markcite{c97}1997), using the Southern SHARC\\footnote{Serendipitous High-redshift Archival ROSAT Cluster} sample of serendipitously detected clusters from deep ROSAT PSPC pointings, that the number of high redshift clusters is consistent with a no evolution model, in direct contrast to Castander et al. \\markcite{RIXOS}(1995). Finally, the EMSS sample has been re-analysed in the light of new optical and X-ray data, which indicates that the evidence for evolution seen by Henry et al. \\markcite{h92}(1992) is not statistically significant (Nichol et al. \\markcite{n97}1997). In this letter we present the high redshift X-ray luminosity function (XLF) of the Collins et al. \\markcite{c97}(1997) cluster sample. In section \\ref{xlf} we describe the calculation of the XLF from this sample and in section \\ref{discussion} we discuss the results. Throughout this letter we have assumed H$_0 = 50$ km/s/Mpc and $q_0 = 0.5$ and quoted luminosities in the 0.5 to 2.0 keV pass band, unless explicitly stated otherwise. ", + "conclusions": "\\label{conclusion} We have used the Southern SHARC survey to create the first $z > 0.3$ XLF derived from ROSAT detected clusters of galaxies. Comparison with the low redshift cluster XLF of both ROSAT (Ebeling et al. \\markcite{e97}1997) and EMSS (Henry et al. \\markcite{h92}1992) clusters shows that there is no evolution in the X-ray luminosities of $L \\sim 10^{44}$ erg s$^{-1}$ clusters at a median depth of $z=0.44$. This is consistent with our analysis of the redshift distribution of this cluster sample (Collins et al. \\markcite{c97}1997) and adds further weight to the body of evidence for no evolution in the cluster population." + }, + "9708/gr-qc9708044_arXiv.txt": { + "abstract": "Assuming that spacetime tunnels -wormholes and ringholes- naturally exist in the universe, we investigate the conditions making them embeddible in Friedmann space, and the possible observable effects of these tunnels, including: lensing and frequency-shifting of emitting sources, discontinuous change of background temperature, broadening and intensity enhancement of spectral lines, so as a dramatic increase of the luminosity of any object at the tunnel's throat. ", + "introduction": "Solutions to Einstein equations corresponding to spacetimes with closed timelike curves (CTCs) have stirred the relativistics, following developments by Lanczos [1], van Stockum [2], G\\\"odel [3], Misner [4] and, more recently, Morris and Thorne [5], Gott [6] and Jensen and Soleng [7] The reason for the excitement and subsequent general disbelief resided much in that, being consistent solutions to Einstein equations for very special kinds of matter, the proposed spacetimes allow for the possibility of time travel and, hence for potential violations of causality [8]. Here we look at the idea that, rather than being interpreted as constructs to be eventually built up from future highly developed technology, spacetime tunnels generating CTCs may spontaneously exist in some regions of our universe, and give rise to observable effects which could be detected even with present technology. Thus, we consider CTCs generated along spacetime tunnels whose mouths embed in distant regions with the Friedmann geometry of the overall universe. The necessary condition for such tunnels to occur in a given region is that, in that region there is a certain proportion of matter with negative energy [8,9]. Two tunnel topologies have been considered so far. That of a two-sphere which gives rise to wormholes [5,10], and that of a two-torus which is associated with the so-called ringholes [11]. Both tunneling types are traversable and convertible into timemachines generating CTCs by simply letting one of the hole's mouths to move toward the other [5,10,11], but whereas the energy density is everywhere negative near the throat of a wormhole, it still becomes positive for values of the angle $\\varphi_2$ (defining the position on the surface circles determined by the torus sections) such that $2\\pi -\\varphi_h >\\varphi_2 > \\varphi_h$, with $\\varphi_h=\\arccos\\frac{b}{a}$, where $a$ and $b$ are the radius of the circumference generated by the circular axis of the torus and that of a torus section, respectively, at the ringhole's throat [11]. The purpose of the present work is to investigate under what conditions can a tunnel be embedded in a cosmological spacetime, and explore the effects that the inner properties of spacetime tunnels may have on the observable characteristics of astronomical objects placed beyond, or passing through these tunnels, relative to an observer whose line of sight to the object traverses or does not traverse the given tunnel. Most of the emphasis will be placed on ringholes, but the results will be always compared with those expected from wormholes. ", + "conclusions": "" + }, + "9708/astro-ph9708003_arXiv.txt": { + "abstract": "In order to better utilize the information contained in the shower images generated by imaging Cherenkov telescopes (IACTs) equipped with cameras with small pixels, images are fit to a parametrization of image shapes gained from Monte Carlo simulations, treating the shower direction, impact point, and energy as free parameters. Monte Carlo studies for a system of IACTs predict an improvement of order 1.5 in the angular resolution. The fitting technique can also be applied to single-telescope images; simulations indicate that the shower direction in space can be reconstructed event-by-event with a resolution of $0.16^\\circ$ to $0.20^\\circ$, allowing to generate genuine source maps. Data from Crab observations with a single HEGRA telescope confirm this prediction. ", + "introduction": "Over the last decade, imaging atmospheric Cherenkov telescopes have proven the prime instrument for $\\gamma$-ray astronomy in the TeV domain~\\cite{review}. Both the orientation and the shape of Cherenkov images are exploited to supress cosmic-ray background in the search for point sources of $\\gamma$ rays: $\\gamma$ shower images seen in the camera point back to the source location, and are characterized by narrow, compact images. In contrast, cosmic rays generate hadronic showers with wider and more diffuse images, and random orientation. Recent improvements of the imaging Cherenkov technique include the use ``high-resolution'' cameras with pixel sizes of $0.15^\\circ$ or less, capable of resolving fine details of the image, and the stereoscopic technique, where a shower is observed simultaneously by several telescopes, allowing to geometrically reconstruct the shower axis, and hence the direction of the primary particle. Since Cherenkov shower images in the camera point to the image of the source, the apparent source of air shower can be reconstructed by superimposing the images of several cameras and intersecting the image axes (Fig.~\\ref{fig_geometry}) \\cite{kohnle_paper,crab_stereo,whipple_source}. \\begin{figure}[hb] \\begin{center} \\mbox{ \\epsfysize7.5cm \\epsffile{geom.eps}} \\end{center} \\caption {Image of a shower in the camera, with the coordinate system $(x,y)$ fixed to the camera, and the system $(x',y')$ defined by the source image and the direction to the shower impact point (after accounting for the reversal of signs occuring for the mirror image). Also illustrated is the reconstruction of the shower direction using images from multiple telescopes.} \\label{fig_geometry} \\end{figure} Opposite to the source image, the shower image points towards the location where the shower axis intersects the plane of the telescope dish. The impact location can be therefore be derived by extrapolating the image axes, starting from the locations of the telescopes, until the lines intersect at one point \\footnote{This simple method applies only if all telescope dishes lie in a plane perpendicular to the telescope axis; the extension to the general case is however straightforward.}. Cherenkov images are traditionally described by parameters related to the first and second moments of the intensity distribution in the camera~\\cite{hillas_param}: the center of gravity, the length of the major and minor axis of the image ``tensor of inertia'', and its orientation. While rather powerful, this method has two shortcomings: it does not make use of the full information obtained with todays cameras, where a typical shower lights up twenty and more pixels, and it requires the use of a ``tail cut'' to eliminate pixels which do not belong to the image, yet show some signal due to night-sky background light. Since this tail cut is usually set at a level around 5 photoelectrons, tails of the image are excluded, too. To provide tools for an improved image analysis, we have, over the last years, developed a technique~\\cite{ulrich,ulrichphd} which makes better use of the information contained in the image, by fitting the observed intensities to a model of Cherenkov images, with the showers characteristics - direction, core location, and energy - as free parameters. In this paper, we describe first the simulation and parametrization of the images, then the fitting procedure and its predicted performance when used in a system of imaging Cherenkov telescopes. Finally, we apply the fitting technique also to single-telescope data and demonstrate its performance using images obtained with one of the HEGRA Cherenkov telescopes during observations of the Crab Nebula. Emphasis on the present work is on this interpretation of single-telescope images; multi-telescope data has become available, but the analysis is still in the early stages~\\cite{crab_stereo}. A similar reconstruction technique has been studied by the CAT group~\\cite{cat_reco}. ", + "conclusions": "" + }, + "9708/astro-ph9708235_arXiv.txt": { + "abstract": "s{ The observations of microlensing events in the Large Magellanic Cloud suggest that a sizable fraction ($\\sim$ 50\\%) of the galactic halo is in the form of MACHOs (Massive Astrophysical Compact Halo Objects) with an average mass $\\sim 0.27 M_{\\odot}$, assuming a standard spherical halo model. We describe a scenario in which dark clusters of MACHOs and cold molecular clouds (mainly of $H_2$) naturally form in the halo at galactocentric distances larger than 10--20 kpc. } ", + "introduction": "A central problem in astrophysics concerns the nature of the dark matter in galactic halos, whose presence is implied by the flat rotation curves in spiral galaxies. As first proposed by Paczy\\'nski \\cite{pa}, gravitational microlensing can provide a decisive answer to that question \\cite{kn:Derujula1}, and since 1993 this dream has started to become a reality with the detection of several microlensing events towards the Large Magellanic Cloud \\cite{al,au}. Today, although the evidence for MACHOs is firm, the implications of this discovery crucially depend on the assumed galactic model. It has become customary to take the standard spherical halo model as a baseline for comparison. Within this model, the average mass reported by the MACHO team is $0.5^{+0.3}_{-0.2}~M_{\\odot}$, which is based upon their first two years data \\cite{al}. The inferred optical depth is $\\tau = 2.1^{+1.1}_{-0.7} \\times 10^{-7}$ when considering 6 events \\footnote{In fact, the two disregarded events are a binary lensing and one which is rated as marginal.} (or $\\tau = 2.9^{+1.4}_{-0.9} \\times 10^{-7}$ when considering all the 8 detected events). Correspondingly, this implies that about 45\\% (50\\% respectively) of the halo dark matter is in form of MACHOs assuming a standard spherical halo model. Instead, using the mass moment method yields an average MACHO mass \\cite{je} of $0.27~M_{\\odot}$. Unfortunately, because of the presently available limited statistics different data-analysis procedures lead to results which are only marginally consistent. Apart from the low-statistics problem -- which will automatically disappear from future larger data samples -- we feel that the real question is whether the standard spherical halo model correctly describes our galaxy \\cite{kn:Ingrosso}. Besides the observational evidence that spiral galaxies generally have flattened halos, recent determinations of the disk scale length, the magnitude and slope of the rotation at the solar position indicate that our galaxy is best described by the maximal disk model, which implies a minimal halo model. This conclusion is further strengthened by the microlensing results towards the galactic centre, which imply that the bulge is more massive than previously thought. For such halo models the expected average MACHO mass should be smaller than within the standard halo model. Indeed, a value $\\sim 0.1~M_{\\odot}$ looks as the most realistic estimate to date and suggests that MACHOs are brown dwarfs. ", + "conclusions": "" + }, + "9708/astro-ph9708221_arXiv.txt": { + "abstract": "We present the results of the detailed surface photometry of a sample of elliptical galaxies in the Hubble Deep Field. In the $<\\mu_e>$--$r_e$ plane the elliptical galaxies of the HDF turn out to follow a `{\\it rest frame}' Kormendy relation, once the appropriate $K+E$ corrections are applied. This evidence, linked to the dynamical information gathered by \\cite{st:et:al}, indicates that these galaxies, even at $z \\simeq 2-3$, lie in the Fundamental Plane, in a virial equilibrium condition. At the same redshifts a statistically signifcant lack of large galaxies (i.e. with $\\log r_e^{kpc} > 0.2$) is observed. ", + "introduction": "The present sample of {\\it early--type} galaxies has been extracted from the second release of the $WFPC2-HDF$ frames, in the $V_{606}$ band. The selection is based on the photometry carried out by the $ESO-STECF-HDF$ Group (\\cite{cl:cou}) with the automated SExtractor algorithm (\\cite{be:ar}). A preliminary list of candidate ellipticals was defined, including all the galaxies satisfying the following criteria: {\\it a)} Kron {\\it STMAG} magnitude in the $V_{606}$ band $\\le$ 26.5; {\\it b)} number of pixels above the threshold limit of $1.3\\sigma$ of the background noise $\\ge$ 200; {\\it c)} star/galaxy classifier $(s/g)\\ \\le$ 0.6 ($s/g$=1 means 'star'). The $\\sim$400 objects matching the above limits were examined to produce a first screening against late-type objects. The preliminary morphological classification was also compared with the ones by \\cite{vdb:et:al} and \\cite{statler}, finding a general good agreement. Detailed surface photometry (luminosity and geometrical profiles) was carried out on the resulting list of 162 {\\it early--type} candidates. As a consequence, 99 objects were left in the final sample of 'bona-fide' {\\it early--type} galaxies. We derived the near infrared ($J$,$H$,$K$) total magnitudes of these galaxies by applying the SExtractor algorithm to the deep images provided by \\cite{dick}. We have restricted our analysis to the galaxies for which a spectroscopic redshift is available (24), or a reliable redshift estimate is possible (24), on the basis of multi--band (optical {\\it and} infrared) photometry. Many galaxies in the present sample show a luminosity profile close to the instrumental PSF. Therefore, in order to extract useful morphological information, we used the '{\\it Multi--Gaussian Expansion}' deconvolution technique (\\cite{bend}) to restore the profiles. Actually, \\cite{fas:et:al} have shown that this technique gives good results in recovering the '{\\it true}' equivalent half-light radius $r_e$ and the corresponding average surface brightness $<\\mu_e>$ down to $r_e^{true} \\simeq FWHM$. ", + "conclusions": "" + }, + "9708/astro-ph9708017_arXiv.txt": { + "abstract": "{\\baselineskip 0.4cm The MACHO Project monitors millions of stars in the Large Magellanic Cloud, the Small Magellanic Cloud and the bulge of the Milky Way searching for the gravitational microlensing signature of baryonic dark matter. This Project has yielded surprising results. An analysis of two years of data monitoring the Large Magellanic Cloud points to {$\\sim 50\\%$} of the mass of the Milky Way's halo in compact objects of {$\\sim 0.5 M_{\\odot}$}. An analysis of one year of monitoring the bulge has yielded more microlensing than predicted without the invocation of a massive bar or significant disk dark matter. The huge database of light curves created by this search is yielding information on extremely rare types of astrophysical variability as well as providing temporal detail for the study of well known variable astrophysical phenomena. The variable star catalog created from this database is previewed and example light curves are presented. } ", + "introduction": "The MACHO Project is monitoring millions of stars every night searching for the gravitational microlensing signature of massive compact halo objects (Machos). This project has the dedicated use the Great Melbourne Telescope at Mount Stromlo. This endeavor was stimulated by Paczy\\'nski's suggestion \\cite{pac86} that gravitational microlensing was a possible way to detect baryonic dark matter in the halo of the Milky Way. The principle of microlensing is simple; if a Macho lies near the line of sight to a background star (the source), it will deflect light from the source and produce two images. For galactic scales, these images cannot be resolved even by HST ($ \\sim 0.001$ arcsec and thus microlensing), but the two unresolved images combine to give an apparent increase in the source brightness. Due to the relative motions of the observer, lens and source, this magnification is transient, so the effect appears as a symmetrical and unique brightening in an otherwise constant star. The duration of the event is a function of the mass of the lens, the relative distances of the lens and source and the motion of the lens with respect to the line of sight. The magnification is just a function of the distance of the lens from the line of sight. The probability that a source will be microlensed is termed the optical depth to microlensing, $\\tau$. The Large Magellanic Cloud (LMC) was chosen as the primary target because its line of sight passes through much of the halo yet it is relatively nearby, and contains millions of stars resolvable in modest seeing with a small telescope. Our secondary target was chosen as the Galactic bulge for two reasons. Although the LMC is circumpolar from Mount Stromlo, it is at too high an airmass to be usefully observed when the bulge is overhead. Calculations before we began taking data \\cite{kim91} suggested that there should be an observable microlensing signal from known stellar populations in the disk and the bulge so observing the bulge would serve to prove whether microlensing could be detected with our system. We also observe the Small Magellanic Cloud (SMC) at a lower priority. The SMC provides a different line of sight through the halo, but is sufficiently farther away and smaller than the LMC that we can only monitor a few million stars with our system. ", + "conclusions": "The unevenly spaced sampling, the wide field of view, the simultaneous color information and the long, densely sampled light curves make the MACHO database a treasure trove for the study of time variability of astrophysical sources. While we have only begun to examine this database outside the context of the search for Machos, the early returns promise much for the future. The MACHO Project is also moving rapidly toward achieving its goal of determining the baryonic content of the halo. The MACHO Project plans to continue operation at least through 1999. At the end of this time, we expect to have a sample of about 50 microlensing events across the face of the LMC, about 1000 events toward the bulge, and a database of light curves for more than 50 million stars in the LMC, the SMC and the bulge, spanning 8 years." + }, + "9708/astro-ph9708198_arXiv.txt": { + "abstract": "The possibility of significant systematic errors due to the use of 1D homogeneous atmospheres in lithium-abundance determinations of cool stars motivates a study of non-local-thermodynamic-equilibrium (\\nlte) effects on \\lii\\ line formation in a 3D solar-granulation simulation snapshot. The \\nlte\\ effect on the equivalent width of the 671\\,nm resonance line is small in 1D models or in integrated light from the granulation model. The line-strength variations over the granulation pattern are however markedly different in \\nlte\\ compared to \\lte\\ -- observations of this may provide diagnostics to \\nlte\\ effects. The effects of horizontal photon exchange found in the granulation model are moderate and due entirely to bound-bound processes, ultraviolet overionization is unimportant. ", + "introduction": "Stellar lithium abundances are potentially very useful for testing astrophysical theories. The abundances derived from observed spectra have spawned a number of scientific debates, for a recent review see Thorburn (1996), also the conference proceedings of Crane (1995) and Spite \\& Pallavicini (1995), and the introduction to Carlsson et al. (1994). It is important in this context that we can be confident in the derived abundances -- which so far mostly have been derived using the questionable assumptions of line formation in local thermodynamic equilibrium (\\lte) and plane-parallel homogeneous \\lte\\ photospheres. Efficient computer codes and extensive atomic data sets have made possible realistic \\nlte\\ spectral-line modeling for light atoms in plane-parallel cool-star photospheres that can provide \\nlte\\ abundance corrections for the convenience of the stellar-abundance community (e.g., Carlsson et al. 1994; Kiselman \\& Carlsson 1996). A notable departure from plane-parallel homogeneity in the quiet solar photosphere is the granulation, now understood as a visible manifestation of convection below the photosphere (e.g., Spruit, Nordlund, \\& Title 1990). We have also some knowledge of granulation on stars adjacent to the Sun in the HR diagram (Gray \\& Nagel 1989; Nordlund \\& Dravins 1990; Dravins \\& Nordlund 1990a, 1990b). How does granulation influence line strengths and abundance determinations? Holweger, Heise, \\& Kock (1990) argued that solar abundance ratios should not be seriously in error since they are based on line-strength ratios which are approximately constant over the solar granulation pattern in the simulations of Steffen (1989). This notion was confirmed observationally by Kiselman (1994a) who found that lines of both neutral and singly ionized species of several elements behaved similarly by being stronger in bright granular regions and weaker in dark intergranular lanes. For other work on lines in realistic granulation simulations that is relevant for abundance analysis, see Nordlund (1984), Bruls \\& Rutten (1992), Atroshchenko \\& Gadun (1994), and Kiselman \\& Nordlund (1995). Of special interest is the question of \\nlte\\ effects on line formation in granulation. Nordlund (1984) found that 3D \\nlte\\ effects could be of some importance for \\ion{Fe}{1} lines in solar granulation, while Kiselman \\& Nordlund (1995) found rather small such effects for \\ion{O}{1}. Note also that Mihalas, Auer, \\& Mihalas (1978) did not find any important 2D effects in their investigation of artificial but solar-like atmospheric structures. Kurucz (1995) claimed that \\nlte\\ effects in extremely metal-poor solar-type stars can cause standard analyses to underestimate lithium abundances with a factor of ten. In this scenario, \\lii\\ lines will be weak in hot photospheric regions where Li is largely ionized. The cooler regions would show strong \\lii\\ lines if \\lte\\ was valid, but the ultraviolet radiation from adjacent hot regions will keep lithium largely ionized also there. The \\lii\\ resonance line in integrated light would thus be very much weaker than the result from a 1D model representing a spatial and temporal average of the photospheric structure. Thus the abundance would be underestimated when an analysis is performed using standard plane-parallel models. So far, there is no quantitative model to verify this effect, which would certainly have important impact on the debates related to stellar lithium abundances. This Letter reports on \\nlte\\ aspects of \\lii\\ line formation as found with experiments on a 3D solar-granulation model. The results are thus directly applicable only for Sun-like stars, but they should give some indication on what we can expect for other stars as well. ", + "conclusions": "Line formation in granulation is not well approximated with simple ``cold'' and ``hot'' regions or streams. The strongly varying vertical photospheric temperature gradient, the typical inversion of the temperature pattern above a certain height, and the presence of inclined thermal inhomogeneities complicate the situation to the extent that conclusions regarding spectral lines must be based on realistic granulation simulations. A case has been made that such simulations together with spatially resolved solar spectroscopy can be used as diagnostics of \\nlte\\ effects. Another paper (Kiselman 1997) will follow up on this and compare simulations with solar observations of \\lii. The current calculations do not show evidence for any large \\nlte\\ effects in quiet granulation that would seriously affect lithium abundance determinations for solar-like stars. This is not to say that effects of granulation are altogether unimportant -- more comprehensive studies are needed to quantify ``granulation abundance corrections'' corresponding to the commonly used \\nlte\\ abundance corrections. One may also speculate that 3D \\nlte\\ effects could be important in regions of enhanced magnetic activity and thus be relevant to problems such as the apparent lithium-abundance spread in the Pleiades (cf. Stuik, Bruls, \\& Rutten 1997). To draw conclusions about very metal-poor solar-type stars and the \\nlte\\ effect proposed by Kurucz (1995) would be to extrapolate -- the nature of these stars' surface inhomogeneities is unknown but probably different from solar (Allende~Prieto et al. 1995). One can note, however, that bound-bound processes may be important for the statistical equilibrium of Li there as well as in the solar case studied here. Wherever line transitions are the dominating drivers of departures from \\lte, any 3D effects are likely to be milder than if bound-free processes govern the \\nlte\\ behavior, because the contrast of granulation or other thermal inhomogeneity is lower at the longer wavelengths of the \\lii\\ spectral lines than in the blue and ultraviolet continua." + }, + "9708/astro-ph9708151_arXiv.txt": { + "abstract": "A multi-wavelength investigation of the candidate supernova remnant G63.7+1.1 and its surrounding interstellar medium is presented. On the basis of radio continuum data we conclude that the object is a filled-center supernova remnant, perhaps in the course of becoming a composite remnant. The morphology of the remnant, along with HI, $^{12}$CO and high resolution IRAS data, suggest that G63.7+1.1 is interacting directly with the ISM, and does not lie in a low density region of the ISM. This in turn strongly suggests that the detected nebula is not surrounded by an invisible halo of supernova ejecta. The association between the SNR and HI and CO features near the tangent point implies a kinematic distance for G63.7+1.1 of $3.8 \\pm 1.5$~kpc. ", + "introduction": "Of the 215 Galactic supernova remnants (SNRs) cataloged by Green (1996), only 9 have been classified as filled-center (FC, also known as ``Crab-like'' or ``plerionic''). An object classified as a FC SNR must have a centrally brightened radio morphology, a flat ($\\alpha>-0.3$), non-thermal, radio spectral index, and a complete lack of an associated limb-brightened shell. It is assumed that the radio emission from these objects is powered by a pulsar interior to the nebula, even if no pulsar has been detected. An object which is centrally brightened but has an associated limb-brightened shell is classified as a composite remnant. The distinction between FC and composite SNRs highlights an important problem associated with FC SNRs, the absence of the limb-brightened shell. If both these types of objects are formed by supernovae which produce pulsars, why should one type have an associated shell but not the other? The most appealing theory to explain the lack of a shell around FC SNRs is that put forward by Chevalier (1977) who proposed that the Crab Nebula (and by extension other FC SNRs) consists of two components, a central pulsar-powered core, which makes up the detected component of the SNR, and an invisible halo of ejecta, formed at the same time as the central pulsar, moving at $\\sim 10^4$~km/s. He surmised that the fast-moving ejecta around these objects are invisible because the FC SNRs lie in low density regions of the ISM and the surrounding halos of ejecta have interacted with only minimal amounts of material; presumably composite SNRs then lie in more ``normal'' environments and the shell is a direct result. This hypothesis can be tested by directly imaging the interstellar medium (ISM) around FC SNRs to determine whether they lie in low density environments. Romani et al. (1990) used archival IRAS and HI data and concluded that the Crab Nebula lies in a large-scale, low-density void in the ISM. Similarly, Wallace et al. (1994) observed the Crab, 3C58, G74.9+1.2 and G21.5-0.9, and found evidence that all but G21.5-0.9 lie in voids (the data for G21.5-0.9 were inconclusive). These studies were limited by their poor resolution ($\\sim 36'$) however, and higher resolution studies are required to image the ISM distribution immediately around these, and similar, objects. This paper is the second of a series in which the ISM around FC SNRs is imaged at high resolution. The first paper (\\cite{wall97}) came to the surprising conclusion that G74.9+1.2 does {\\it not} lie in a low density region of the ISM (despite a low density region being suggested by the observations presented in \\cite{wall94}). In this paper we present the results of a multi-frequency investigation into the FC SNR candidate G63.7+1.1 and its surroundings. G63.7+1.1 was identified as a FC SNR candidate by Taylor et al. (1992) based upon its non-thermal radio spectrum, low ratio of infrared to radio flux density, and morphology (the entire survey on which Taylor et al. (1992) is based can be found in \\cite{wsrtsurv}). The data of Taylor et al. (1992) are limited by their relatively poor resolution however, and further observations are required to confirm the nature of G63.7+1.1. Presented here are radio continuum and line, and high resolution far infrared data. The observations are discussed in Sec.~2 and the data presented in Sec.~3. In Sec.~4 the nature of G63.7+1.1 and its surroundings is discussed, and a model for the interaction between the two is presented. The paper is summarized in Sec.~5. ", + "conclusions": "High resolution radio continuum observations of the SNR candidate G63.7+1.1 have been presented, along with far-infrared, HI, and CO observations of its surroundings. These observations show that G63.7+1.1 has a centrally-brightened radio morphology with no detectable limb brightening; no limb brightened shell is seen down to a level of $\\Sigma_{1\\ GHz} \\sim 2 \\times 10^{-22}$W~m$^{-2}$~Hz$^{-1}$~sr$^{-1}$. The spectral index of the emission, $\\alpha=-0.28 \\pm 0.02$, along with the detection of polarization, imply a non-thermal emission process. G63.7+1.1 is thus most likely a filled-center SNR. The morphology of G63.7+1.1 can be broken down into two components, a bright resolved core and a fainter, more uniform, plateau. The plateau is shown to have a sharp, unresolved, intensity gradient marking its outer edge around most of the remnant, but exhibits a much shallower gradient to the south-east. Modeling of the emissivity profile within the remnant suggests that an increase in emissivity must occur around the portions of the remnant having a sharp intensity gradient. This rise in emissivity implies an increase in magnetic field strength, the injection of additional emitting particles, or both, at the edges of the SNR. This may suggest that G63.7+1.1 is in the process of forming a limb-brightened component at its edges and may thus be in transition from FC to composite SNR. The observations of the ISM around G63.7+1.1 suggest that the remnant lies adjacent to molecular material with which it is interacting. Some of this molecular material may have been injected into the SNR interior, perhaps creating a ridge of brighter radio emission with associated features seen at 100$\\mu m$ and in CO~Feature~D. The radial velocity of the molecular material with which the SNR is interacting indicates that the SNR lies at or near the tangent point of the Galactic rotation curve in this direction, implying a distance to the remnant of $3.8 \\pm 1.5$~kpc. The interaction between G63.7+1.1 and the ISM implies both that the SNR is {\\it not} lying in a low density region of the ISM, and that the SNR is {\\it not} surrounded by a fast moving halo of ejecta. This result demands a rethinking of models used to describe FC SNRs and their evolution." + }, + "9708/astro-ph9708036_arXiv.txt": { + "abstract": "We present a new approach to calculating the statistical distributions for magnification, shear, and rotation of images of cosmological sources due to gravitational lensing. In this approach one specifies an underlying Robertson-Walker cosmological model together with relevant information on the clumping of matter on scales much smaller than the Hubble radius. The geodesic deviation equation is then integrated backwards in time until the desired redshift is reached, using a Monte Carlo procedure wherein each photon beam in effect ``creates its own universe'' as it propagates. The approach is somewhat similar to that used in ``Swiss cheese'' models, but the ``cheese'' has been completely eliminated, the matter distribution in the ``voids'' need not be spherically symmetric, the total mass in each void need equal the corresponding Robertson-Walker mass only on average, and we do not impose an ``opaque radius'' cutoff. The case where the matter in the universe consists of point masses is studied in detail, and it is shown that the statistical distributions of the lensing images are essentially independent of both the mass spectrum and the clustering properties of the point masses, provided that the clustering is spherical. Detailed results for the distribution of the magnification of images are presented for the point mass case, as well as a number of other matter distributions. We apply our results (i) to argue that the positive correlation recently found between quasar luminosity and the number of absorption line systems is not likely to be due to lensing, and (ii) to determine the amount of ``noise'' and possible bias produced by lensing in measurements of $q_0$ using distant supernovae. ", + "introduction": "\\label{intro} In recent years there has been a great deal of interest in studying the effects on cosmologically distant sources produced by gravitational lensing due to intervening matter. In many cases of interest, the lensing effects can be assumed to be produced by a single galaxy or cluster of galaxies, and one can use the detailed structure of the images produced by lensing to extract a great deal of information about the mass distribution of the galaxy or cluster. However, in other circumstances of interest one may be interested in the cumulative lensing effects produced by many different objects (or voids), and one may be primarily interested in statistical distributions of the image brightenings and/or distortions, rather than the detailed modeling of any individual lens system. Two examples of the latter circumstances are the following: (1) Vanden Berk {\\em et al.}~\\cite{vandenberk} have presented evidence for a positive correlation between quasar luminosity and the number of intervening Carbon IV absorption clouds. Could this correlation be the result of the cumulative gravitational lensing effects produced by the mass distributions associated with these clouds? (2) Efforts are currently underway to use supernovae occurring at cosmological distances as standard candles for tests of $q_0$~\\cite{perlmutter}. How much ``noise'' in the apparent luminosity distribution of the supernovae would be expected from gravitational lensing effects? Could any useful information about the distribution of matter in the universe be extractable from this ``noise''? The main purpose of this paper is to present a new approach for determining cumulative gravitational lensing effects on cosmological scales due to inhomogeneities in the matter distribution of the universe. As explained further below, in this approach one specifies an underlying Robertson-Walker cosmological model together with one's assumptions concerning the detailed clumping and clustering of matter in the universe. Both the Robertson-Walker model and the clumping/clustering of matter may be specified arbitrarily, provided that the clustering of matter occurs only on scales much smaller than the Hubble radius and that the average density of the matter distribution corresponds to that of the underlying Robertson-Walker model. Our approach then enables one to accurately obtain statistical distributions for the luminosity, shear, and rotation of images of ``standard candle'' (nearly) point sources at any cosmological redshift. When multiple images occur, however, even statistical information about the number of images and the relationships between the images cannot be easily extracted using our approach, since that would require us to keep track of the relationship between finitely (as opposed to infinitesimally) separated null geodesics. Nevertheless, statistical information about the luminosity, shear, and rotation of the individual images occurring in multiple images is included in our distributions. The rest of this section will be devoted to an overview of our approach for determining statistical lensing effects in inhomogeneous universes. Subsection~\\ref{Cosmological Model} introduces our cosmological model, presenting and justifying the metric which provides the framework for our results. Subsection~\\ref{Propagation of Photon Beams} discusses lensing effects on the propogation of photon beams within the cosmology, while Subsection~\\ref{The Local Nature} discusses the local nature of these effects. Subsection~\\ref{Our Method} gives a general overview of our method, and Subsection~\\ref{ss:relevant_scales} discusses the relevant scales of the model. In Section~\\ref{method} we present our procedure for calculating statistical lensing effects in more explicit detail. In Section~\\ref{points_and_clustering} we analyze the case where all of the matter in the universe can be treated as being comprised of point masses (satisfying Eq.~(\\ref{mlim})). Other distributions of mass are considered in Subsection~\\ref{ss:cases}, and then in Subsection~\\ref{ss:checks} we perform some consistency checks on our results. Applications of our work to the analysis of lensing effects by quasar absorption systems are given in Section~\\ref{york}, and applications to the effects of lensing on supernovae luminosity are given in Section~\\ref{supernovae}. \\subsection{Cosmological Model} \\label{Cosmological Model} To explain our approach, we first need to state our cosmological assumptions with more precision. We assume that the spacetime metric of the universe is globally well approximated (on {\\em all} scales) by a ``Newtonianly perturbed Robertson-Walker metric'' of the form \\begin{equation} ds^2 = -(1 + 2\\phi)\\,d \\tau^2 + (1 - 2\\phi) a^2(\\tau) \\left[\\frac{dr^2}{1-kr^2} + r^2 (d \\theta^2 + sin^2 \\theta\\,d \\varphi^2)\\right], \\label{metric} \\end{equation} where $k = 0, \\pm 1$. We shall refer to the metric obtained by setting $\\phi = 0$ in Eq.~(\\ref{metric}) as the {\\em underlying Robertson-Walker model}. The spatial metric of this underlying Robertson-Walker model is $a^2 h_{ab}$, where \\begin{equation} h_{ab} \\equiv \\frac{1}{1-kr^2}\\,dr_a dr_b + r^2 (d \\theta_a d \\theta_b + sin^2 \\theta\\,d \\varphi_a d \\varphi_b) \\label{smetric} \\end{equation} is either the metric of a unit 3-sphere ($k = 1$), a unit 3-hyperboloid ($k=-1$), or flat 3-space ($k = 0$). Without loss of generality, we may assume that the spatial average of $\\phi$ vanishes, since a spatially constant part of $\\phi$ could be absorbed into the definitions of $\\tau$ and $a$. We also assume that throughout spacetime---or at least out to distance scales of order $R_H$, where $R_H \\equiv H^{-1} = a/\\dot{a}$ denotes the Hubble radius of the underlying Robertson-Walker model---we have \\begin{equation} |\\phi| \\ll 1 . \\label{d1phi} \\end{equation} We further assume that time derivatives of $\\phi$ are much smaller than spatial derivatives, i.e., \\begin{equation} |\\partial \\phi/ \\partial \\tau|^2 \\ll a^{-2} h^{ab} D_a \\phi D_b \\phi, \\label{d2phi} \\end{equation} with similar relations holding for the higher time derivatives. Here $D_a$ denotes the spatial derivative operator associated with $h_{ab}$, and $h^{ab}$ denotes the inverse of $h_{ab}$ (so $a^{-2} h^{ab}$ is the inverse spatial metric of the underlying Robertson-Walker model). It is important to note that spatial derivatives of $\\phi$ may locally be very large compared with scales set by the underlying Robertson-Walker model. However, we assume that products of first spatial derivatives of $\\phi$ are small compared with second derivatives, i.e., \\begin{equation} (h^{ab} D_a \\phi D_b \\phi)^2 \\ll h^{ac} h^{bd} D_a D_b \\phi D_c D_d \\phi. \\label{d3phi} \\end{equation} Finally, we assume that the matter stress-energy tensor, $T_{ab}$ ({\\em not} including the cosmological constant term), is everywhere such that, in the rest frame of the underlying Robertson-Walker model, the energy density of matter greatly dominates the other components of $T_{ab}$. In this case $T_{ab}$ is approximately of the ``matter dominated'' form \\begin{equation} T_{ab} \\approx \\rho u_a u_b, \\label{Tab} \\end{equation} where $u^a$ is the unit (in the metric of Eq.~(\\ref{metric})) timelike vector field orthogonal to the surfaces of constant $\\tau$. Eqs.~(\\ref{d1phi})--(\\ref{Tab}) are the only assumptions we shall need to obtain eqs.~(\\ref{ee1'})--(\\ref{poisson}) below.\\footnote{ E. Linder (private communication) has claimed that the approximation $\\epsilon^2/\\kappa \\ll 1$ of references \\cite{linder}, \\cite{futamase}, and \\cite{futamase2} is also needed for the validity of our equations below. We do not agree with this claim.} However, in Subsection~\\ref{The Local Nature} we shall also assume that there is a (co-moving) scale ${\\cal R} \\ll R_H$ such that no strong correlations in the density of matter occur on scales greater than ${\\cal R}$. We now substitute the metric form of Eq.~(\\ref{metric}) and the matter stress-energy of Eq.~(\\ref{Tab}) into Einstein's equation, possibly with a nonvanishing cosmological constant, $\\Lambda$. We make the approximations of {Eqs.~(\\ref{d1phi})--(\\ref{Tab})}, and also drop all terms (like $\\Lambda \\phi$ and $\\rho \\phi$) which are small compared with the curvature of the underlying Robertson-Walker metric. The nonvanishing components of Einstein's equation then yield\\footnote{In addition to the two equations given here---which correspond to the time-time and diagonal space-space components of Einstein's equation---there are also contributions to the time-space components of Einstein's equation of the form $\\rho v_a$ (where $v_a$ denotes the velocity of the matter relative to the Hubble flow), $(\\dot{a}/a) D_a \\phi$, and mixed time-space derivatives of $\\phi$. These terms need not everywhere be small compared with the curvature of the underlying Robertson-Walker metric. If only these terms were considered, the time-space components of Einstein's equation would yield additional equations for $\\phi$ which would be inconsistent with Eq.~(\\ref{poisson}) below. This difficulty is resolved by allowing for the presence of nonvanishing time-space components of the metric, $g_{0 \\mu}$ (with $\\mu = 1,2,3$), satisfying $|g_{0 \\mu}| \\ll |\\phi|$. The time-space components of Einstein's equation then become, in essence, equations which determine $g_{0 \\mu}$ (see Sec.~4.4a of~\\cite{wald} for further details in the ordinary Newtonian case). However, since $g_{0 \\mu}$ makes a negligible correction to the effects calculated in this paper, we shall ignore its presence below and, correspondingly, will not consider the time-space components of Einstein's equation.} \\begin{eqnarray} &3 \\ddot{a}/a = \\Lambda - 4 \\pi \\rho + a^{-2} h^{ab} D_a D_b \\phi& \\label{ee1} \\\\ &3(\\dot{a}/a)^2 = \\Lambda + 8 \\pi \\rho -2a^{-2} h^{ab} D_a D_b \\phi - 3k/a^2,& \\label{ee2} \\end{eqnarray} where the dots denote derivatives with respect to $\\tau$. The spatial average of these equations yields the usual form of the matter dominated Einstein equations for the underlying Robertson-Walker metric, namely \\begin{eqnarray} &3 \\ddot{a}/a = \\Lambda - 4 \\pi \\bar{\\rho}& \\label{ee1'} \\\\ &3(\\dot{a}/a)^2 = \\Lambda + 8 \\pi \\bar{\\rho} - 3k/a^2,& \\label{ee2'} \\end{eqnarray} where $\\bar{\\rho}$ denotes the spatial average of $\\rho$. Subtracting Eqs.~(\\ref{ee1'}) and~(\\ref{ee2'}) from Eqs.~(\\ref{ee1}) and~(\\ref{ee2}), we find the remaining content of Einstein's equation is that $\\phi$ satisfies the Poisson equation\\footnote{ Nonlinear terms in $\\phi$, such as $a^{-2}h^{ab} D_a\\phi\\,D_b\\phi=D_a\\phi\\,D^a\\phi$, are neglected in eq.~(\\ref{poisson}) because they are small compared with the term linear in $\\phi$ (see eq.~(\\ref{d3phi})). On the other hand, since the spatial average of $D_aD^a\\phi$ vanishes, the neglect of the spatial average of nonlinear terms like $D_a\\phi\\,D^a\\phi$ in eqs.~(\\ref{ee1'}) and~(\\ref{ee2'}) is justified as follows. We have \\begin{eqnarray*} \\int_{\\rm V} \\!D_a\\phi\\,D^a\\phi\\,\\,d{\\rm V} &=&-\\int_{\\rm V} \\!\\phi\\,D_aD^a\\phi\\,\\,d{\\rm V}\\\\ &=&-4\\pi\\int_{\\rm V}\\!\\phi\\,\\delta\\rho\\,\\,d{\\rm V}\\\\ &=&-4\\pi\\int_{\\rm V}\\!\\phi\\,(\\rho-\\bar\\rho)\\,\\,d{\\rm V}. \\end{eqnarray*} The integral of $\\phi\\,\\rho$ is much less than the integral of $\\rho$, as $\\phi\\ll1$ and $\\rho$ is nonegative. The same argument holds for the $\\phi\\,\\bar\\rho$ term. Thus, under our assumptions, the spatial average of $D_a\\phi\\,D^a\\phi$ is much less than $\\bar\\rho$, which justifies dropping the former in eqs.~\\ref{ee1'} and~\\ref{ee2'}. } \\begin{equation} a^{-2} h^{ab} D_a D_b \\phi = 4 \\pi \\delta \\rho, \\label{poisson} \\end{equation} where \\begin{equation} \\delta \\rho \\equiv \\rho - \\bar{\\rho}. \\label{drho} \\end{equation} We emphasize that it is completely consistent with our assumptions to have, locally, $\\delta \\rho \\gg \\bar{\\rho}$. It is essential that this be allowed if Eq.~(\\ref{metric}), together with Eqs.~(\\ref{d1phi})--(\\ref{d3phi}), are intended as an accurate description of our universe, since we commonly find $\\delta \\rho \\sim 10^{30} \\bar{\\rho}$ in our vicinity. Thus, in our model the matter is assumed to have an energy density much greater than its stresses, and is assumed to move non-relativistically with respect to the Hubble flow defined by the underlying Robertson-Walker model. However, unlike a Robertson-Walker model, this matter may be distributed in a very inhomogeneous manner; in particular, as already noted, the fluctuations in the mass density may be very large compared with the spatial average of the mass density. Consequently, the local curvature of spacetime may differ drastically from that of a Robertson-Walker model. Nevertheless, in our cosmological model, the Hubble flow of the matter and the causal structure of spacetime correspond very closely to the underlying matter dominated Robertson-Walker model, whose mass density is equal to the average density of matter in the universe. It is useful to examine the form taken by the metric of Eq.~(\\ref{metric}) in a locally Minkowskian frame associated with an observer moving with the Hubble flow, which, for convenience, we take to be located at $r = 0$. To do so we define a new radial coordinate, $R$, by \\begin{equation} R = a r, \\label{R} \\end{equation} and a new time coordinate, $T$, by \\begin{equation} T = \\tau + {1\\over2}\\frac{\\dot{a}}{a} R^2. \\label{T} \\end{equation} In these new coordinates the metric of Eq.~(\\ref{metric}) takes the form \\begin{eqnarray} ds^2 &=& - (1 + 2 \\phi - R^2 \\ddot{a}/a)\\,dT^2 + \\left(1 - 2 \\phi + R^2\\left[\\left({\\dot{a}}/{a}\\right)^2+ k/a^2\\right]\\right)\\,dR^2 \\nonumber \\\\ & & \\mbox{} + (1 - 2 \\phi) R^2\\, d\\Omega^2, \\label{n1} \\end{eqnarray} where we have dropped all terms of order $R^3$ and higher in distance from the origin. Transforming to an isotropic radial coordinate, then further transforming to the corresponding Cartesian coordinates $X,Y,Z$, and, finally, substituting from Einstein's equations (Eqs.~(\\ref{ee1'}) and~(\\ref{ee2'})) for the underlying Robertson-Walker model, we obtain \\begin{equation} ds^2 = - (1 + 2 \\Phi - \\Lambda R^2/3)\\,dT^2 + (1 - 2 \\Phi - \\Lambda R^2/6)[dX^2 + dY^2 + dZ^2], \\label{n2} \\end{equation} where \\begin{equation} \\Phi \\equiv \\phi + 2 \\pi R^2 \\bar{\\rho}/3, \\label{Phi} \\end{equation} and where, to the approximation in which we are working (i.e., dropping terms of order $R^3$ and higher), we have $R^2 = X^2 + Y^2 + Z^2$. Thus, $\\Phi$ satisfies the ordinary Poisson equation \\begin{eqnarray} \\nabla^2 \\Phi & = & \\nabla^2 \\phi + 4 \\pi \\bar{\\rho} \\nonumber \\\\ & = & 4 \\pi (\\delta \\rho + \\bar{\\rho}) \\nonumber \\\\ & = & 4 \\pi \\rho. \\label{Poisson} \\end{eqnarray} When $\\Lambda = 0$, Eq.~(\\ref{n2}) is precisely the usual form of Newtonianly perturbed Minkowski spacetime (see, e.g., Sec.~4.4a of \\cite{wald}). Thus, in the spacetime of Eq.~(\\ref{metric}), when $\\Lambda = 0$, Newtonian gravity holds to a very good approximation in the vicinity of any observer following the Hubble flow, where ``in the vicinity'' here means on scales much smaller than the Hubble radius. Even when $\\Lambda \\neq 0$, if $|\\delta \\rho| \\gg \\bar{\\rho}$ in the neighborhood of the observer, realistic values of $\\Lambda$ have $\\Lambda R^2 \\ll \\Phi$ out to distances much smaller than the Hubble radius. Thus, Newtonian gravity holds to an excellent approximation in the vicinity of such observers as well. In summary, we may characterize our cosmological model of Eq.~(\\ref{metric}), together with Eqs.~(\\ref{d1phi})--(\\ref{d3phi}), as one which corresponds closely to a Robertson-Walker model as far as the Hubble flow of the matter and the causal structure of the spacetime are concerned, but in which the local distribution of matter may be highly inhomogeneous. In addition, as we have just noted, on scales small compared with those set by the underlying Robertson-Walker model, Newtonian gravity holds to a very good approximation. Apart from negligibly small regions of spacetime which contain black holes or other strong field objects, we believe that our universe is accurately described by this model. In any case, our model is a relatively precise, mathematically consistent cosmological model which describes the spacetime structure and distribution of matter on all scales, and is not in obvious conflict with any observed properties of our universe. \\subsection{Propagation of Photon Beams} \\label{Propagation of Photon Beams} Let us now consider this cosmological model from the perspective of photons ($\\equiv$ null geodesics) propagating in it, and compare this to what photons would encounter in a Robertson-Walker model. All gravitational focusing and shearing effects on an infinitesimal beam of light rays in the vicinity of a null geodesic $\\gamma$ are described by the geodesic deviation equation (see, e.g. \\cite{wald}) \\begin{equation} \\frac{d^2 \\eta^a}{d \\lambda^2} = - {R_{bcd}}^a k^b k^d \\eta^c, \\label{gd} \\end{equation} where $k^a$ is the tangent to $\\gamma$ corresponding to affine parameter $\\lambda$, and $\\eta^a$ is the deviation vector to an infinitesimally nearby null geodesic in the beam. The Riemann curvature tensor appearing in Eq.~(\\ref{gd}) can be decomposed into its Ricci and Weyl pieces in the usual way (see, e.g. \\cite{wald}) \\begin{equation} R_{abcd} = C_{abcd} +\\left(g_{a[c}R_{d]b}-g_{b[c}R_{d]a}\\right) -{1\\over3}Rg_{a[c}g_{d]b}. \\label{wr} \\end{equation} The Ricci curvature directly produces a rate of change of convergence of the beam of geodesics, while the Weyl curvature directly produces a rate of change of shearing. In a Robertson-Walker model the Weyl tensor vanishes and, by Einstein's equation, the Ricci tensor is of the form $R_{ab} = 8 \\pi (T_{ab} - 1/2\\,T g_{ab})$, with $T_{ab}$ given by Eq.~(\\ref{Tab}). The geodesic deviation equation then takes the form \\begin{equation} \\frac{d^2 \\eta^a}{d \\lambda^2} = - 4 \\pi \\omega^2 \\rho \\eta^a, \\label{gd2} \\end{equation} where $\\omega$ is the frequency of the photon as measured in the Robertson-Walker rest frame. This corresponds to a steady increase in the convergence of the beam of geodesics, with no shear. Contrast this behavior with the propagation of photons in the cosmological model of Eq.~(\\ref{metric}) in the case where the matter is highly clumped on various scales, but with no (or negligible) matter distributed between the clumps. In this case, the Ricci tensor vanishes along the geodesic, except for rare instances when the photon propagates through a clump of matter. On these rare occasions, the Ricci curvature briefly becomes extremely large compared with that of the underlying Robertson-Walker model. The Weyl curvature also will be small except in similarly rare instances of propagation through (or very near) a sufficiently dense clump of matter. Thus, when the matter distribution is highly clumped, at almost all times the propagation of a beam of photons in the spacetime of Eq.~(\\ref{metric}) would be indistinguishable from propagation in flat spacetime. Occasionally, however, the beam may receive a strong ``kick'' of Weyl and/or Ricci curvature. Thus, the local history of a photon propagating in the spacetime of Eq.~(\\ref{metric}) could hardly be more different from the local history of a photon propagating in a Robertson-Walker model! Nevertheless, there are some global correspondences. In particular, since the causal structure of the spacetime of Eq.~(\\ref{metric}) corresponds closely to that of the underlying Robertson-Walker metric, at each redshift\\footnote{Since $|\\phi|\\ll1$ and the velocity of matter relative to the Hubble flow is small, we neglect the difference between redshifts in the metric of Eq.~(\\ref{metric}) and in the underlying Robertson-Walker model.} the area of the boundary of the past of an event in the spacetime of Eq.~(\\ref{metric}) must be very nearly equal to the area of the past light cone of the corresponding event in the underlying Robertson-Walker metric. We will return to this point in Subsection~\\ref{ss:checks}. In order to calculate magnification and shear effects on a (nearly) point source due to gravitational lensing, we need to integrate the geodesic deviation equation (Eq.~(\\ref{gd})) along a null geodesic connecting the source to the observer. To do this, we need to know the curvature along the geodesic. The curvature is determined directly by a knowledge of the underlying Robertson-Walker model together with $\\phi$. We will assume that, in the underlying Robertson-Walker model, the distance scales set by the spatial curvature and $\\Lambda$ are at least as large as the Hubble radius, $R_H$. The spacetime curvature of the Robertson-Walker model is then of order $1/R_H{}^2$. Contributions of $\\phi$ to the spacetime curvature which are smaller than $1/R_H{}^2$ will therefore be neglected. From Eq.~(\\ref{poisson}), together with the assumption that $\\phi$ is bounded and has vanishing spatial average, it follows that $\\phi$ is uniquely determined by specifying the matter distribution $\\delta \\rho$. However, Eq.~(\\ref{poisson}) is a nonlocal equation, so in principle the locally encountered curvature could depend upon the distribution of matter in arbitrarily distant parts of the universe.\\footnote{Note that since, for an open universe, $\\delta \\rho$ does not fall off to zero at infinity, we cannot assume, a priori, that $\\phi$ is given in terms of $\\delta \\rho$ by the usual Poisson integral expression that would hold for a localized mass distribution.} Nevertheless, we shall now argue that, under our cosmological assumptions, only the distribution of matter within $R_H$ is relevant. \\subsection{The Local Nature of the Influence of Matter on Photon Beams} \\label{The Local Nature} Let $S$ be a sphere of proper radial distance $R_H$ centered about the point $x$ at which we wish to evaluate $\\phi$. Let $G_D(x,x')$ denote the Dirichlet Green's function for the equation $a^{-2} h^{ab} D_a D_b G(x,x') = - 4 \\pi \\delta(x,x')$ for the region enclosed by $S$. (A simple, explicit formula for $G_D$ in the case of flat geometry can be found, e.g., in Sec.~2.6 of~\\cite{jack}.) Then, by Green's identity, we have \\begin{equation} \\phi(x) = - \\int_V G_D (x,x') \\delta \\rho (x') dV' - \\frac{1}{4 \\pi} \\int_S \\phi (x') \\hat{r}'{}^a D'_a G_D (x,x')\\,dS', \\label{gi} \\end{equation} where the volume integral extends only over the region enclosed by $S$. Under our above assumptions, the contribution of $\\phi$ to the curvature is given directly in terms of the second spatial derivatives of $\\phi$, since the contributions from the time derivatives of $\\phi$, products of first derivatives of $\\phi$, etc., have been assumed to be negligible compared with the linear contributions from the second spatial derivatives of $\\phi$. Differentiating Eq.~(\\ref{gi}), we obtain \\begin{eqnarray} D_a D_b \\phi (x) &=& - \\int_V D_a D_b G_D (x,x') \\delta \\rho (x')\\,dV' \\nonumber \\\\ & & \\mbox{} - \\frac{1}{4 \\pi} \\int_S \\phi (x') \\hat{r}'{}^a D'_a D_a D_b G_D (x,x')\\,dS'. \\label{gi2} \\end{eqnarray} However, the surface term in Eq.~(\\ref{gi2}) is of order $|\\phi|/R_H{}^2$, and thus, in view of Eq.~(\\ref{d1phi}), it can be neglected. Therefore the curvature at $x$ is determined by the matter distribution only within a Hubble radius of $x$, as we desired to show. It should be emphasized that this conclusion is {\\em not} a consequence of any causality arguments but, rather, follows directly from our above {\\em assumption} that $\\phi$ is small at distances of order $R_H$, as is necessary for the underlying Robertson-Walker metric to be a good description of spacetime structure on cosmological scales. We now make the additional assumption that there is a (co-moving) scale ${\\cal R} \\ll R_H$ such that no strong correlations in the distribution of matter occur on scales greater than ${\\cal R}$. Under these circumstances it seems clear that the curvature at a given point can be accurately calculated---at least for the purposes of determining geodesic deviation---by taking into account only the matter distribution within a distance ${\\cal R}$ of that point. We have not attempted to give a precise formulation or proof of this claim, but a justification for it can be given as follows. First we note that, by Einstein's equation, the Ricci curvature is determined by the matter distribution in a completely local manner. Therefore, matter can have a nonlocal influence on a photon beam only via Weyl curvature. To calculate the Weyl curvature associated with a distribution of matter we need to evaluate the trace-free part of the second derivatives of $\\phi$, as given by Eq.~(\\ref{gi2}) with the surface term omitted. We break up the volume $V$ in Eq.~(\\ref{gi2}) into a union of regions of size ${\\cal R}$, excluding the ball of radius ${\\cal R}$ centered at $x$. In the case of flat spatial geometry, each of these regions will make a contribution of order $m/D^3$ to the Weyl tensor at $x$, where $D$ is the distance of the region from $x$, and $m$ is of the order of the expected mass, $\\bar{\\rho}\\,{\\cal R}^3$, contained in that region. However, by our assumption, there will be no correlations between the contributions from the different regions. Hence, by a simple ``random walk'' estimate, we find that the total contribution to the Weyl tensor at $x$ from all of $V$ except for the ball of radius ${\\cal R}$ centered at $x$ should be no greater than $\\sim m/{\\cal R}^3 \\sim \\bar{\\rho}$. Similar estimates hold if the geometry is curved or a cosmological constant is present, since $G_D$ will differ significantly from the flat case only at distances comparable to $R_H$, and the contributions from these regions should be negligible. We note that $\\bar{\\rho}$ is the same order of magnitude as the curvature of the underlying Robertson-Walker metric. A Ricci curvature of this magnitude and having a consistent sign (as occurs in the Robertson-Walker model) could have a significant effect on the convergence of a beam of photons propagating over cosmological distances. However, a randomly fluctuating Weyl curvature of this magnitude should have a completely negligible effect upon the shear (merely adding a tiny bit of ``noise'' to the Weyl curvature resulting from nearby matter), and an even smaller effect upon the convergence. Thus, no significant error should be made by considering only the curvature resulting from the presence of matter within ${\\cal R}$ of the photon path, as we desired to show. Since we have assumed that ${\\cal R} \\ll R_H$ and that the distance scales set by the spatial curvature and/or $\\Lambda$ are at least as large as $R_H$, the Dirichlet Green's function within ${\\cal R}$ of $x$ will be well approximated by $1/r$, where $r$ denotes the proper distance between $x$ and $x'$. Thus, Eq.~(\\ref{gi})---with the surface term omitted and the volume integral restricted to a ball of radius ${\\cal R}$ around $x$---reduces to the usual Poisson integral formula, and the curvature can be obtained from formulas arising from ordinary Newtonian gravity (see Sec.~\\ref{method} below). It is somewhat more convenient to work with the potential $\\Phi$ of Eq.~(\\ref{Phi}) rather than $\\phi$. It follows that $\\Phi$ is given by the usual Poisson integral formula of $\\rho$ (rather than $\\delta \\rho$) over the region enclosed by ${\\cal R}$. \\subsection{Our Method} \\label{Our Method} The basic idea of our procedure in its most general context can now be explained. We choose an underlying Robertson-Walker model and (co-moving) scale, ${\\cal R}$, with ${\\cal R} \\ll R_H$ in the present universe.\\footnote{More generally, we could specify a probability distribution for ${\\cal R}$, although we shall not do so in this paper.} We then specify a probability distribution for how the matter is distributed within ${\\cal R}$. This probability distribution may vary with cosmological time; it is constrained only by the requirement that the average amount of mass contained within ${\\cal R}$ agree with that occurring in the underlying Robertson-Walker model. We then perform a ``Monte Carlo'' propagation of a beam of photons backward in time, starting from the present, in the following manner: We prescribe a matter distribution (chosen from our probability distribution) in a ball of radius ${\\cal R}$. We calculate the Newtonian potential for this matter distribution, and the corresponding curvature. Then we choose a random impact parameter for the entry of a photon into this ball, and we integrate Eq.~(\\ref{gd}) through the ball. (In this step, we take the photon trajectory to be a ``straight line'', i.e., we do not attempt to include the (completely negligible) corrections due to the tiny bending angle.) When the photon exits from this ball, we use the underlying Robertson-Walker model to update the frequency of the photon relative to the local rest frame of the matter, and to update the proper radius corresponding to the comoving scale ${\\cal R}$. Then we choose a matter distribution in a new ball of comoving radius ${\\cal R}$, choose another random impact parameter for entry of the photon into this ball,\\footnote{Note that, in general, this would require the balls to overlap slightly. We neglect this overlap in our analysis.} and repeat the above calculations. We continue until the photon has reached the desired redshift. By repeating this sequence of calculations a large number of times---for most of our models we performed about 2,000 such ``runs''---we build up good statistics on what happens to beams of photons on our past light cone. From this we obtain, for any given model, good statistical information on the magnification, shear, and rotation of images of (nearly) point sources at any redshift. We will spell out the details of our procedure more explicitly in the next section. In comparison with other approaches, ours most closely resembles the ``Swiss cheese'' models, wherein one takes a matter dominated Robertson-Walker model, removes the dust from spherical balls, and redistributes the mass within these balls in some other (arbitrarily chosen) spherically symmetric manner. However, it differs from the Swiss cheese models in the following significant ways: (i) The ``cheese'' has been completely eliminated. (ii) The mass within a given ball need not be equal to the corresponding Robertson-Walker mass, though equality must still hold on average. (iii) The matter distribution within the balls need not be spherically symmetric. (iv) We do not consider the propagation of photons in a single, fixed cosmological model. Rather, each photon in effect ``creates its own cosmological model'' via our Monte Carlo procedure during the course of its propagation. (v) Although it is not a necessary facet of the Swiss cheese models, most analyses of the Swiss cheese models~\\cite{kantowski,dyerroeder1} have attempted to calculate only averages of certain lensing quantities, and, in the course of doing so, have imposed an ``opaque radius'' cutoff---within which photons are absorbed---which biases the results towards defocusing relative to Robertson-Walker models. Our analysis determines the probability distributions for magnification, shear, and rotation of sources by doing an exact, Monte Carlo calculation, imposing no opaque radius cutoff. As we shall see, our results show no bias towards defocusing relative to the underlying Robertson-Walker model, provided that all of the high luminosity images are included (see Sec.~\\ref{supernovae} for further discussion). Our approach also bears some similarity to analyses which start with a model of the matter distribution in the universe---obtained analytically~\\cite{schneider,SEF,rauch} or from N-body codes~\\cite{matzner,wambsganss2,wambsganss,tomita}---and then project the matter into lens planes lying between the source and observer. Ray shooting methods are used to numerically obtain bending angles of a large sample of photons, from which the amplification and shear distribution of images can then be computed. Our approach uses the geodesic deviation equation rather than the lens equation and is considerably simpler and more flexible. It also avoids any artifacts resulting from putting all the matter into lens planes. \\subsection{The Relevant Scales of Clustering and Clumping} \\label{ss:relevant_scales} Two final issues remain to be addressed: (1) What clustering scale ${\\cal R}$ should be chosen to adequately model statistical lensing effects in our universe, i.e., what is the largest scale on which the clustering of matter has an important effect upon lensing? (2) On what scales (below ${\\cal R}$) does one have to model the details of the matter distribution in order to adequately treat statistical lensing effects, i.e., what is the smallest scale on which the clumping of matter has an important effect? In analyzing these questions, it is convenient to view galaxies as the basic ``building blocks'' of the distribution of matter in the universe. (Although we do not exclude the possibility that substantial amounts of matter may be distributed between galaxies, we assume that such matter is distributed in a relatively uniform way.) It is essential to take into account the clumping of matter on the scale of galaxies in order to adequately model lensing effects. In essence, the first question above asks to what extent the clustering of the galaxies themselves must be taken into account, while the second question asks to what extent the clumping substructure of the matter within galaxies must be taken into account. As already noted, it follows from Einstein's equation that the Ricci curvature is determined by the matter distribution in a completely local manner. The effects of Ricci curvature on lensing should therefore depend only upon the density contrasts associated with galaxies, and not upon the ``shape'' of galaxies. This will be verified explicitly in Subsection~\\ref{ss:checks}. Furthermore, these Ricci curvature effects should depend only weakly on the clustering of galaxies, since the clustering should merely produce some correlations in the times of passage of a photon through different galaxies, and these effects should largely ``wash out'' over cosmological distance scales. Thus, we believe that the clustering of galaxies should have a negligible influence on lensing effects produced by Ricci curvature. On the other hand, simple estimates show that the Weyl curvature of a spherical aggregate of matter of mass $m$ and radius $r$ can have a substantial effect on lensing only if the matter ``lies within its own Einstein radius'', $r_E$, i.e., only if \\begin{equation} r^2 \\lap r_E{}^2 \\sim m D, \\label{er} \\end{equation} where $D$ denotes a cosmological distance and we use units where $G = c = 1$. Note that this relationship is marginally satisfied by individual galaxies (or at least by their central cores), so the Weyl curvature of individual galaxies can (at least occasionally) produce significant lensing effects. Clustering of galaxies can produce important Weyl curvature effects only in circumstances when the clusters themselves satisfy Eq.~(\\ref{er}). This {\\em does} occur in the central portions of rich clusters of galaxies, so the effects of clustering cannot always be assumed to be negligible. However, in the limit where galaxies can be treated as ``point masses''---as occurs if Eq.~(\\ref{er}) is satisfied by a wide margin---it follows from the analysis given in Subsection~\\ref{ss:clustering} below that even very strong clustering of the galaxies will have at most a tiny effect on the lensing probability distributions for the magnification, shear, and rotation of (nearly) point sources. (On the other hand, clustering {\\em would} still have an important effect on some lensing quantities, such as bending angles, which we do not calculate here.) Thus, clustering effects can be of importance for the statistical lensing quantities treated here only only when individual galaxies fail to satisfy Eq.~(\\ref{er}), but these galaxies form clusters which satisfy Eq.~(\\ref{er}) (at least in their core regions). In these circumstances the neglect of the clustering of galaxies should underestimate the lensing effects somewhat. However, we do not believe that such circumstances arise frequently enough to have an important influence on the statistical lensing quantities we calculate. Furthermore, as we shall conjecture in Subsection~\\ref{ss:conjecture}, the point mass results should provide a firm upper limit to lensing effects, even when galactic clustering is present. Consequently, in this paper we shall take ${\\cal R}$ to be the scale of the separation between galaxies, thereby neglecting lensing effects resulting from the clustering of galaxies. For the reasons detailed above, we do not expect that this will result in any significant errors in our calculations of the probability distributions for magnification, shear, and rotation of images of cosmologically distant sources. Some evidence in favor of this expectation will be given in Section~\\ref{supernovae}, where we will obtain results in close agreement with~\\cite{wambsganss}, despite our neglect of the effects of clustering.\\footnote{This expectation could be further tested by re-doing our analysis taking ${\\cal R}$ to be the scale of separation of clusters of galaxies and using appropriately chosen probability distributions for the distribution of mass within clusters. We have not yet attempted to carry out such an analysis.} We turn now to the issue of how small a scale of clumping of matter we must consider in order to calculate gravitational lensing effects. In principle the clumping of matter on arbitrarily small scales (including atomic and sub-atomic scales) could have an important effect on lensing---though we would have to use physical, rather than geometric, optics to calculate these effects when the objects are so small that the scale of variation of the gravitational field becomes less than the wavelength of the light. However, the finite size of the source which is being lensed provides an effective cutoff to lensing produced by clumping on small scales. This follows because the lens merely magnifies (as well as shears and rotates) the image of the source, keeping the surface brightness constant~\\cite{SEF}. Thus, if the angular size of the (assumed to be uniform) source is much larger than the angular scale associated with the lens, the lensing effects caused by clumping should have little effect, as only a relatively small part of the source would be magnified by the presence of a clump of matter (and the rest of the source may be correspondingly demagnified by the absence of matter between clumps). In other words, the net angular size of the image of a source of finite size will not be significantly affected by sufficiently small scale lensing, and consequently, the luminosity of the image also will not be greatly affected.\\footnote{If the angular size of the source is much larger than the angular scale of separation between the clumps of matter, then the lensing effects of the matter should wash out completely.} The angular scale of the source is $\\sim r_S/D_S$, where $r_S$ denotes the size of the source and $D_S$ denotes its distance, and the angular scale associated with the lens is $\\sim r_E/D_L$, where $D_L$ denotes the distance of the lens. Taking $D_S$ and $D_L$ to be cosmological in scale and using Eq.~(\\ref{er}) for $r_E$, we find that lensing effects should not be important unless the mass of the lens satisfies \\begin{equation} m \\gap r_S{}^2 /D. \\label{mlim} \\end{equation} The smallest sources of interest here (central regions of quasars and supernova shells at an early stage of expansion) have $r_S \\gap 10^{-3}$ light years, so taking $D \\sim 10^{10}$ light years, we find \\begin{equation} m \\gap 10^{-3} M_\\odot. \\label{mlim2} \\end{equation} Clumping of matter on mass scales smaller than Eq.~(\\ref{mlim2}) should not be relevant for the sources we consider. However, the clumping of matter down to the scale of Eq.~(\\ref{mlim2}) is potentially of importance. In particular, the clumping of matter in galaxies into stars can have a significant effect upon the probability distribution for the magnification of light emitted from quasars and supernovae. Fortunately, it is not necessary to model a galaxy as $10^{11}$ or so point mass stars in order to calculate its lensing effects. The clumpiness of matter will be relevant only very close to the path of the photon. If, say, we let $d \\sim 100 r_E$, where $r_E$ denotes the Einstein radius of a single star, then the discreteness of the galactic mass distribution due to stars which lie outside of a tube of radius $d$ around the photon path can be ignored, i.e., outside of the tube the galactic matter distribution can be treated as continuous. Consequently, in our analysis we will take account of all ``microlensing'' effects due to small scale clumping of matter (say, into stars) in the following manner: First, we model the galaxy as a continuous mass distribution and compute its Newtonian potential. Then, when a photon passes through the galaxy in our Monte Carlo simulations, we remove the continuous galactic matter lying within cylindrical radius $d$ of the path (or the portion of this matter assumed to be clumped into stars), and subtract the Newtonian potential of this removed matter. Finally, we randomly redistribute this removed mass back into the cylinder in the form of stars, and we add in the Newtonian potential of these ``point masses''. In this manner we take full account of the small scale clumping of matter in a computationally efficient way. ", + "conclusions": "Note also that since lensing simply magnifies or demagnifies images relative to the underlying Robertson-Walker model---but does not affect the surface brightness of the images---the apparent luminosity of an image of a source is proportional to $1/A$. Since, as just noted above, the probability that a beam ``hits'' a given source is proportional to $A$, the expected luminosity (i.e., photon flux) in each beam is exactly the same as in the underlying Robertson-Walker model. In particular, our analysis automatically builds in the fact that the expected total luminosity agrees with that of the underlying Robertson-Walker model. For beams of photons which have not undergone caustics, the largest possible area is the ``flat space'' (or ``empty beam'' \\cite{dyerroeder2}) value, $A_{\\rm Flat}$, corresponding to setting the curvature to zero in the geodesic deviation equation. This value is marked on (most of) the figures. It should be noted that after a photon beam undergoes a caustic, its area typically becomes very large (and negative)---significantly larger in magnitude than the flat space value. (An indication of this fact can already be seen in Fig.~1.) To avoid problems with the scale of our figures, we did not attempt to plot any area values less than $- A_{\\rm Flat}$. This accounts for the ``gap'' at the beginning of our plots. We will refer to an image associated with a photon beam which has not undergone a caustic as a {\\it primary image}.\\footnote{This corresponds to the ``type I'' image of~\\cite{SEF}.} If $p$ denotes the event representing our telescope at the present time, then any event $q$ which lies on the boundary of the past of $p$ must be connected to $p$ by a null geodesic whose corresponding photon beam has not undergone a caustic (see, e.g., \\cite{wald}). Since the world line of any source must intersect the boundary of the past of $p$, it follows that every source must have at least one primary image (see also~\\cite{SEF}). In Subsection~\\ref{ss:checks} below we shall argue that for spherical distributions of matter it is very rare that a source would have more than one primary image, but for very dense cylindrical matter distributions, multiple primary images are common. Every primary image of a source must be at least as bright as it would be if it were placed in flat spacetime at the same affine parameter distance~\\cite{SEF}. On the other hand, secondary images (corresponding to photon beams which have undergone one or more caustics) can be arbitrarily faint. Of course, a secondary image of a source can also be bright and, in particular, can be brighter than a primary image of that source. Each secondary image of a source must have at least one associated primary image of the same source, and (since the total number of images must be odd~\\cite{SEF}) must also have other associated secondary images. If the angular separation of these images is very small, it may not be possible to resolve the separate images. One of the shortcomings of our method is that we do not have a good way of determining (even statistically) which primary and secondary images are associated with each other, since this would require us to analyze photon trajectories which differ by finite (as opposed to infinitesimal) separations. Thus, if the different images of the same source are not resolved---the case of microlensing---we are unable to predict the probability distribution in total luminosity. Figs.~\\ref{F:z_05}--\\ref{F:z_3} show our results for a universe filled with point masses corresponding to an underlying spatially flat Robertson-Walker cosmology with $\\Omega = 1$ and $\\Lambda = 0$. As with all plots shown here, we have taken $H_0 = 70\\ \\mbox{km}\\,\\mbox{s}^{-1}\\,\\mbox{Mpc}^{-1}$. We also took ${\\cal R} = 2\\ \\mbox{Mpc}$, although as argued above, the results should be independent of the choice of ${\\cal R}$. As can be seen from the graphs, the percentage of photon beams which have undergone caustics ranges from about $5\\%$ at redshift $z = 1/2$ to over $35\\%$ at redshift $z = 3$. Note also that by redshift $z=3$ about $20\\%$ of the primary images are less than half as bright ($A > 2$) as they would have been in in the underlying Robertson-Walker cosmology. Taking account of the factor of $|A|$ mentioned above, and assuming that each source has only one primary image (see Subsec.~\\ref{ss:checks} below), we find that the probability that the primary image of a randomly placed source at $z=3$ will be demagnified relative to Robertson-Walker by at least a factor of $2$ is $1/2$. Since these photon trajectories do not pass near any of the point masses, it seems unlikely that such sources will have any (bright) secondary images. Thus, even if multiple images cannot be resolved, it appears that in this cosmology, at redshift $z=3$, $50\\%$ of all sources should be dimmer by at least a factor of $2$ relative to the underlying Robertson-Walker model. \\begin{figure} \\begin{center} \\begin{picture}(0,0)% \\epsfig{file=z_05.eps}% \\end{picture}% \\setlength{\\unitlength}{3947sp}% \\begingroup\\makeatletter\\ifx\\SetFigFont\\undefined% \\gdef\\SetFigFont#1#2#3#4#5{% \\reset@font\\fontsize{#1}{#2pt}% \\fontfamily{#3}\\fontseries{#4}\\fontshape{#5}% \\selectfont}% \\fi\\endgroup% \\begin{picture}(5238,2742)(4360,-3691) \\put(7887,-2753){\\makebox(0,0)[lb]{\\smash{\\SetFigFont{12}{14.4}{\\familydefault}{\\mddefault}{\\updefault}\\special{ps: gsave 0 0 0 setrgbcolor}\\% of photon beams\\special{ps: grestore}}}} \\put(4360,-1538){\\makebox(0,0)[lb]{\\smash{\\SetFigFont{12}{14.4}{\\familydefault}{\\mddefault}{\\updefault}\\special{ps: gsave 0 0 0 setrgbcolor}$A/A_{RW}$\\special{ps: grestore}}}} \\end{picture} \\end{center} \\caption{Area vs. \\% of photon beams at $z=0.5$, for an $\\Omega=1$, $\\Lambda=0$ universe, with matter distributed in the form of point masses. The dashed line represents the flat spacetime (empty beam) area, and the dotted line represents the Robertson-Walker area.} \\label{F:z_05} \\end{figure} \\begin{figure} \\begin{center} \\begin{picture}(0,0)% \\epsfig{file=z_1.eps}% \\end{picture}% \\setlength{\\unitlength}{3947sp}% \\begingroup\\makeatletter\\ifx\\SetFigFont\\undefined% \\gdef\\SetFigFont#1#2#3#4#5{% \\reset@font\\fontsize{#1}{#2pt}% \\fontfamily{#3}\\fontseries{#4}\\fontshape{#5}% \\selectfont}% \\fi\\endgroup% \\begin{picture}(5239,2742)(4359,-3691) \\put(7878,-2760){\\makebox(0,0)[lb]{\\smash{\\SetFigFont{12}{14.4}{\\familydefault}{\\mddefault}{\\updefault}\\special{ps: gsave 0 0 0 setrgbcolor}\\% of photon beams\\special{ps: grestore}}}} \\put(4359,-1674){\\makebox(0,0)[lb]{\\smash{\\SetFigFont{12}{14.4}{\\familydefault}{\\mddefault}{\\updefault}\\special{ps: gsave 0 0 0 setrgbcolor}$A/A_{RW}$\\special{ps: grestore}}}} \\end{picture} \\end{center} \\caption{Area vs. \\% of photon beams at $z=1.0$, for an $\\Omega=1$, $\\Lambda=0$ universe. Matter is distributed in the form of point masses.} \\label{F:z_1} \\end{figure} \\begin{figure} \\begin{center} \\begin{picture}(0,0)% \\epsfig{file=z_2.eps}% \\end{picture}% \\setlength{\\unitlength}{3947sp}% \\begingroup\\makeatletter\\ifx\\SetFigFont\\undefined% \\gdef\\SetFigFont#1#2#3#4#5{% \\reset@font\\fontsize{#1}{#2pt}% \\fontfamily{#3}\\fontseries{#4}\\fontshape{#5}% \\selectfont}% \\fi\\endgroup% \\begin{picture}(5239,2742)(4359,-3691) \\put(4359,-1714){\\makebox(0,0)[lb]{\\smash{\\SetFigFont{12}{14.4}{\\familydefault}{\\mddefault}{\\updefault}\\special{ps: gsave 0 0 0 setrgbcolor}$A/A_{RW}$\\special{ps: grestore}}}} \\put(7870,-2753){\\makebox(0,0)[lb]{\\smash{\\SetFigFont{12}{14.4}{\\familydefault}{\\mddefault}{\\updefault}\\special{ps: gsave 0 0 0 setrgbcolor}\\% of photon beams\\special{ps: grestore}}}} \\end{picture} \\end{center} \\caption{Area vs. \\% of photon beams at $z=2.0$, for an $\\Omega=1$, $\\Lambda=0$ universe. Matter is distributed in the form of point masses.} \\label{F:z_2} \\end{figure} \\begin{figure} \\begin{center} \\begin{picture}(0,0)% \\epsfig{file=z_3.eps}% \\end{picture}% \\setlength{\\unitlength}{3947sp}% \\begingroup\\makeatletter\\ifx\\SetFigFont\\undefined% \\gdef\\SetFigFont#1#2#3#4#5{% \\reset@font\\fontsize{#1}{#2pt}% \\fontfamily{#3}\\fontseries{#4}\\fontshape{#5}% \\selectfont}% \\fi\\endgroup% \\begin{picture}(5442,2876)(4156,-3825) \\put(4156,-1660){\\makebox(0,0)[lb]{\\smash{\\SetFigFont{12}{14.4}{\\familydefault}{\\mddefault}{\\updefault}\\special{ps: gsave 0 0 0 setrgbcolor}$A/A_{RW}$\\special{ps: grestore}}}} \\put(7816,-2793){\\makebox(0,0)[lb]{\\smash{\\SetFigFont{12}{14.4}{\\familydefault}{\\mddefault}{\\updefault}\\special{ps: gsave 0 0 0 setrgbcolor}\\% of photon beams\\special{ps: grestore}}}} \\end{picture} \\end{center} \\caption{Area vs. \\% of photon beams at $z=3.0$, for an $\\Omega=1$, $\\Lambda=0$ universe. Matter is distributed in the form of point masses.} \\label{F:z_3} \\end{figure} The results at $z = 3$ for a universe filled with point masses corresponding to an open Robertson-Walker model with $\\Omega_0=0.1$ and $\\Lambda=0$ is plotted in Fig.~\\ref{F:om01}. It can be seen that the lensing effects here are dramatically weaker than in the $\\Omega=1$ model. In particular, in this cosmology less than $10\\%$ of the photon beams have undergone a caustic by $z=3$, and the maximum de-magnification relative to Robertson-Walker is only $0.85$ (but over half the primary images suffer nearly this de-magnification). \\begin{figure} \\begin{center} \\begin{picture}(0,0)% \\epsfig{file=om01.eps}% \\end{picture}% \\setlength{\\unitlength}{3947sp}% \\begingroup\\makeatletter\\ifx\\SetFigFont\\undefined% \\gdef\\SetFigFont#1#2#3#4#5{% \\reset@font\\fontsize{#1}{#2pt}% \\fontfamily{#3}\\fontseries{#4}\\fontshape{#5}% \\selectfont}% \\fi\\endgroup% \\begin{picture}(5239,2742)(4359,-3691) \\put(4359,-1627){\\makebox(0,0)[lb]{\\smash{\\SetFigFont{12}{14.4}{\\familydefault}{\\mddefault}{\\updefault}\\special{ps: gsave 0 0 0 setrgbcolor}$A/A_{RW}$\\special{ps: grestore}}}} \\put(7878,-2753){\\makebox(0,0)[lb]{\\smash{\\SetFigFont{12}{14.4}{\\familydefault}{\\mddefault}{\\updefault}\\special{ps: gsave 0 0 0 setrgbcolor}\\% of photon beams\\special{ps: grestore}}}} \\end{picture} \\end{center} \\caption{Area vs. \\% of photon beams at $z=3.0$, for an $\\Omega_0=0.1$, $\\Lambda=0$ universe. Matter is distributed in the form of point masses.} \\label{F:om01} \\end{figure} Finally, the results at $z = 3$ for a universe filled with point masses corresponding to a spatially flat Robertson-Walker model with $\\Omega_0=0.1$ and $\\Omega_\\Lambda \\equiv \\Lambda/3{H_0}^2 =0.9$ are plotted in Fig.~\\ref{F:om01la09}. This distribution is intermediate between the cases of $\\Omega=1$, $\\Lambda=0$ and $\\Omega_0=0.1$, $\\Lambda=0$. \\begin{figure} \\begin{center} \\begin{picture}(0,0)% \\epsfig{file=om01la09.eps}% \\end{picture}% \\setlength{\\unitlength}{3947sp}% \\begingroup\\makeatletter\\ifx\\SetFigFont\\undefined% \\gdef\\SetFigFont#1#2#3#4#5{% \\reset@font\\fontsize{#1}{#2pt}% \\fontfamily{#3}\\fontseries{#4}\\fontshape{#5}% \\selectfont}% \\fi\\endgroup% \\begin{picture}(5246,2809)(4352,-3158) \\put(4352,-985){\\makebox(0,0)[lb]{\\smash{\\SetFigFont{12}{14.4}{\\familydefault}{\\mddefault}{\\updefault}\\special{ps: gsave 0 0 0 setrgbcolor}$A/A_{RW}$\\special{ps: grestore}}}} \\put(7884,-2131){\\makebox(0,0)[lb]{\\smash{\\SetFigFont{12}{14.4}{\\familydefault}{\\mddefault}{\\updefault}\\special{ps: gsave 0 0 0 setrgbcolor}\\% of photon beams\\special{ps: grestore}}}} \\end{picture} \\end{center} \\caption{Area vs. \\% of photon beams at $z=3.0$, for an $\\Omega_0=0.1$, $\\Omega_\\Lambda=0.9$ universe. Matter is distributed in the form of point masses.} \\label{F:om01la09} \\end{figure} A sample of our results for shear is given in Fig.~\\ref{F:shear1}. Here we have plotted the magnification, $\\mu$, relative to the empty beam value, versus the axial ratio, $\\epsilon$, of the beam at redshift $z = 2$ for a universe filled with point masses for the case $\\Omega=1$ and $\\Lambda=0$. This figure corresponds to Fig.~11.12 of \\cite{SEF}, except that we also have included the points with $\\mu < 1$, arising from beams which have undergone caustics. The agreement between the figures appears to be excellent. \\begin{figure} \\begin{center} \\begin{picture}(0,0)% \\epsfig{file=shear1.eps}% \\end{picture}% \\setlength{\\unitlength}{3947sp}% \\begingroup\\makeatletter\\ifx\\SetFigFont\\undefined% \\gdef\\SetFigFont#1#2#3#4#5{% \\reset@font\\fontsize{#1}{#2pt}% \\fontfamily{#3}\\fontseries{#4}\\fontshape{#5}% \\selectfont}% \\fi\\endgroup% \\begin{picture}(5029,3219)(3969,-4750) \\put(6751,-4711){\\makebox(0,0)[lb]{\\smash{\\SetFigFont{12}{14.4}{\\familydefault}{\\mddefault}{\\updefault}\\special{ps: gsave 0 0 0 setrgbcolor}$\\epsilon$\\special{ps: grestore}}}} \\put(3969,-3015){\\makebox(0,0)[lb]{\\smash{\\SetFigFont{12}{14.4}{\\familydefault}{\\mddefault}{\\updefault}\\special{ps: gsave 0 0 0 setrgbcolor}$\\mu$\\special{ps: grestore}}}} \\end{picture} \\end{center} \\caption{Magnification vs. axial ratio at $z = 2$, for an $\\Omega = 1$, $\\Lambda=0$ universe filled with point masses. The solid line gives the fit $\\mu=(1+\\epsilon)^2/(4\\epsilon)$, which would hold if the lensing was done by a single point mass, as described in~\\protect\\cite{SEF}. This figure compares well with Fig.~11.12 of that reference.} \\label{F:shear1} \\end{figure} A sample of our results for rotation is given in Fig.~\\ref{F:rotation1}. Here we plot the magnitude of rotation angle, $|\\Theta|$ (in radians), versus photon beam number (ordered by area, as described above) at redshift $z = 3$ for a universe filled with point masses for the case $\\Omega=1$ and $\\Lambda=0$. From the figure it can be seen that the photon beams which have not undergone caustics generally have a very small rotation, but those which have undergone caustics have undergone such a large rotation that their orientation is practically random (see \\cite{ehlers} for a general discussion of the behavior of beams near caustics). As noted above, no rotation would occur for lensing produced by a single point mass. \\begin{figure} \\begin{center} \\begin{picture}(0,0)% \\epsfig{file=rotation1.eps}% \\end{picture}% \\setlength{\\unitlength}{3947sp}% \\begingroup\\makeatletter\\ifx\\SetFigFont\\undefined% \\gdef\\SetFigFont#1#2#3#4#5{% \\reset@font\\fontsize{#1}{#2pt}% \\fontfamily{#3}\\fontseries{#4}\\fontshape{#5}% \\selectfont}% \\fi\\endgroup% \\begin{picture}(5044,3108)(3954,-4282) \\put(3954,-2183){\\makebox(0,0)[lb]{\\smash{\\SetFigFont{12}{14.4}{\\familydefault}{\\mddefault}{\\updefault}\\special{ps: gsave 0 0 0 setrgbcolor}$|\\Theta|$\\special{ps: grestore}}}} \\put(5889,-4246){\\makebox(0,0)[lb]{\\smash{\\SetFigFont{12}{14.4}{\\familydefault}{\\mddefault}{\\updefault}\\special{ps: gsave 0 0 0 setrgbcolor}\\% of photon beams\\special{ps: grestore}}}} \\end{picture} \\end{center} \\caption{Magnitude of rotation angle vs. \\% of photon beams at redshift $z = 3$, for an $\\Omega=1$, $\\Lambda=0$ universe filled with point masses. The demarcation between beams which have undergone caustics and those which have not occurs at 36\\% (see Fig.~\\ref{F:z_3}). The restriction of $|\\Theta|$ to the range $0$ to $\\pi/2$ (rather than $0$ to $\\pi$) for beams which have undergone a single caustic is due to our convention in the definition of $\\Theta$ in that case, as explained below Eq.~(\\ref{shear}). The first 3\\% of the photon beams have undergone two caustics.} \\label{F:rotation1} \\end{figure} Finally Fig.~\\ref{F:effective1} shows how remarkably small the effects of clustering are. The right-most curve shows the magnification versus photon beam number for point mass galaxies in a universe with $\\Omega=1$ and $\\Lambda=0$, at a redshift of 3; it is the same curve as shown in Fig.~\\ref{F:z_3} above. Also shown is the curve for (point mass) stars clustered into uniform density galaxies of radius $200\\ \\mbox{kpc}$. This curve is statistically indistinguishable from the curve for point mass galaxies. The left-most curve is for (point mass) stars clustered into uniform density galaxies of radius $20\\ \\mbox{kpc}$. This clustering distribution was chosen (in a parameter search, varying the galactic radius) so as to {\\em maximize} the deviation from the random distribution. As expected, the maximum deviation occurs for galaxies (composed of point mass stars) whose radii are close to their Einstein radii. It can be seen from the figure that there is a slight (but statistically significant) diminution of the lensing effectiveness due to the clustering. \\begin{figure} \\begin{center} \\begin{picture}(0,0)% \\epsfig{file=effective1.eps}% \\end{picture}% \\setlength{\\unitlength}{3947sp}% \\begingroup\\makeatletter\\ifx\\SetFigFont\\undefined% \\gdef\\SetFigFont#1#2#3#4#5{% \\reset@font\\fontsize{#1}{#2pt}% \\fontfamily{#3}\\fontseries{#4}\\fontshape{#5}% \\selectfont}% \\fi\\endgroup% \\begin{picture}(5445,2876)(3553,-3825) \\put(5611,-1561){\\makebox(0,0)[lb]{\\smash{\\SetFigFont{12}{14.4}{\\familydefault}{\\mddefault}{\\updefault}\\special{ps: gsave 0 0 0 setrgbcolor}$r=20$ kpc\\special{ps: grestore}}}} \\put(7209,-2746){\\makebox(0,0)[lb]{\\smash{\\SetFigFont{12}{14.4}{\\familydefault}{\\mddefault}{\\updefault}\\special{ps: gsave 0 0 0 setrgbcolor}\\% of photon beams\\special{ps: grestore}}}} \\put(3553,-1562){\\makebox(0,0)[lb]{\\smash{\\SetFigFont{12}{14.4}{\\familydefault}{\\mddefault}{\\updefault}\\special{ps: gsave 0 0 0 setrgbcolor}$A/A_{RW}$\\special{ps: grestore}}}} \\put(6181,-3451){\\makebox(0,0)[lb]{\\smash{\\SetFigFont{12}{14.4}{\\familydefault}{\\mddefault}{\\updefault}\\special{ps: gsave 0 0 0 setrgbcolor}$r=200\\ \\mbox{kpc}$ and $r=0$\\special{ps: grestore}}}} \\end{picture} \\end{center} \\caption{Area vs. \\% of photon beams at $z = 3$, for an $\\Omega=1$, $\\Lambda=0$ universe. Graphs are shown for point mass galaxies ($r=0$), and for uniform density spheres of $r=20\\ \\mbox{kpc}$ and $r=200\\ \\mbox{kpc}$, each composed of (point mass) stars.} \\label{F:effective1} \\end{figure} \\subsection{A Conjecture} \\label{ss:conjecture} We conclude this section with a conjecture, based upon the fact that spherical clustering of point masses appears to slightly reduce their lensing effectiveness, together with our expectation (borne out in all of our simulations) that point masses are more effective in lensing than any bodies of finite extent: {\\it Conjecture:} For any underlying Robertson-Walker cosmological model at any redshift $z$, randomly distributed point masses provide the most ``effective'' distribution of matter for lensing in the following sense: Let $A_{\\rm rpm}(x)$ denote the area as a function of the percentage of photon beams for a universe filled with a random distribution of point masses (see Figs.~\\ref{F:z_05}--\\ref{F:om01la09}). Let $x_1$ denote the $x$-value such that $A_{\\rm rpm}$ equals the Robertson-Walker area, i.e., $A_{\\rm rpm}(x_1) = 1$. Then for any other matter distribution, we have $A(x) > A_{\\rm rpm}(x)$ for all $x \\leq x_1$. In particular, the greatest number of caustics is achieved for the case of randomly distributed point masses." + }, + "9708/astro-ph9708200_arXiv.txt": { + "abstract": "We present the results of a wide-field [O~III] $\\lambda 5007$ survey for planetary nebulae (PN) in M87 and its surrounding halo. In all, we identify 338 PN candidates in a $16\\arcmin \\times 16\\arcmin$ field around the galaxy; 187 of these objects are in a statistical sample which extends to $m_{5007} = 27.15$. We show that the planetary nebula luminosity function (PNLF) of M87's halo is unlike any PNLF observed to date, with a shape that differs from that of the empirical law at the 99.9\\% confidence level. In addition, we find that the PNLF of M87's outer halo differs from that of the galaxy's inner regions at a high degree of certainty ($\\sim 92\\%$). We show that both these effects are most likely due to the existence of intracluster PN, many of which are foreground to M87. These intracluster objects explain the ``overluminous'' [O~III] $\\lambda 5007$ sources previously identified by Jacoby, Ciardullo, \\& Ford (1990), and present us with a new tool with which to probe the morphological and dynamical properties of the cluster. By modifying the maximum likelihood procedures of Ciardullo \\etal (1989a) to take into account the presence of ``field objects,'' and using an assumed M31 distance of 770~Kpc (Freedman \\& Madore 1990) with a Burstein \\& Heiles (1984) reddening law, we derive a distance modulus to M87 of $30.79 \\pm 0.16$ ($14.4 \\pm 1.1$~Mpc). This value is in excellent agreement with the previous survey of Jacoby, Ciardullo, \\& Ford (1990) and contradicts the assertion of Bottinelli \\etal (1991) and Tammann (1992) that the PNLF distance to Virgo has been underestimated due to inadequate survey depth. ", + "introduction": "The planetary nebula luminosity function (PNLF) technique is one of the simplest methods for determining extragalactic distances out to $\\sim 20$~Mpc. One takes a deep exposure of a galaxy through a filter which passes light at [O~III] $\\lambda 5007$, and compares the image to a slightly deeper exposure off the emission line. Those stellar objects that appear on the [O~III] $\\lambda 5007$ image but not on the off-band frame are either planetary nebulae (PN) or compact H~II regions. If the target object is a normal elliptical or S0 galaxy with no star formation, then the presence of H~II regions can generally be discounted, and one is left with a list of PN, from which one can form a luminosity function. The power of the PNLF technique comes from the shape of the luminosity function. At faint magnitudes, the PNLF has the power law form predicted from models of uniformly expanding shells surrounding slowly evolving central stars (Henize \\& Westerlund 1963; Jacoby 1980). However, observations have demonstrated that the bright end of the PN luminosity function dramatically breaks from this relation, and falls to zero very quickly, within $\\sim 0.7$~mag (cf.~Jacoby \\etal 1992). It is the constancy of this cutoff magnitude, $M^*$, and its high monochromatic luminosity, that makes the PNLF such a useful standard candle. The shape and absolute magnitude of the PNLF cutoff has been successfully reproduced theoretically by a number of authors, including Jacoby (1989), Dopita, Jacoby, \\& Vassiliadis (1992), M\\'endez \\etal (1993), Han, Podsiadlowski, \\& Eggleton (1994), and Richer, McCall, \\& Arimoto (1997). Nevertheless, Bottinelli \\etal (1991) and Tammann (1992) have argued that the bright-end of the PNLF is actually a power law, and thus observations which do not reach the break in the luminosity function are not useful for distance measurements. In support of this model, Bottinelli \\etal and Tammann point to the ``overluminous'' [O~III] $\\lambda 5007$ sources found by Jacoby, Ciardullo, \\& Ford (1990) in Virgo elliptical galaxies, which can plausibly be argued to be part of a high-luminosity tail to the PNLF\\null. By adopting the power-law model, and ignoring evidence for curvature in the observed PNLF of Virgo, Bottinelli \\etal and Tammann have argued that the PNLF distance to this cluster is biased towards too low a value. Although Kolmogorov-Smirnov and $\\chi^2$ tests show this interpretation is highly unlikely, the most unambiguous way to test the hypothesis is to perform a deep, wide-field [O~III] $\\lambda 5007$ imaging survey of the Virgo ellipticals and better define the shape of the faint-end of the PNLF\\null. In this paper, we report the results of a $16\\arcmin \\times 16\\arcmin$ [O~III] $\\lambda 5007$ survey centered on the central elliptical of Virgo, M87. In \\S 2, we give the details of the survey, present the positions and magnitudes of 338 planetary nebulae found in the galaxy's envelope and outer halo, and estimate the photometric accuracy of our measurements by comparing our derived PN magnitudes with those given in Jacoby, Ciardullo, \\& Ford (1990). In \\S 3, we select two statistically complete subsets of these planetaries, and demonstrate that our planetary nebula luminosity function extends well onto the Henize \\& Westerlund (1963) tail, making a distance determination to the galaxy possible. In \\S 4, we discuss the surprising result that the PNLF of M87's outer halo has a cutoff that is $\\sim 0.2$~mag brighter than that for the inner part of the galaxy. We then show that, in retrospect, this behavior could have been predicted, since the intracluster stars of Virgo should produce a considerable number of planetary nebulae, and some of these objects will be foreground to M87{}. In \\S 7, we include this ``field'' contribution in our maximum likelihood analysis, and derive a distance to M87 of $14.4 \\pm 1.1$~Mpc, in good agreement with the previous PNLF distance determination to the galaxy. This result vitiates the hypothesis of Bottinelli \\etal (1991) and Tammann (1992) that the bright-end of the PNLF is an unbounded power-law. We conclude by discussing the implications intracluster PN have for morphological and dynamical studies of nearby clusters. \\section {Observations and Reductions} On 6 and 7 April 1995 we surveyed a $16\\arcmin \\times 16\\arcmin$ region of sky around M87 with the T2KB CCD on the Kitt Peak 4-m telescope, which afforded a plate scale of $0\\parcsec 47$ per pixel. Our on-band data consisted of seven exposures totaling 6.8~hours through a $\\sim 30$~\\AA\\ full-width-half-maximum (FWHM) interference filter centered at 5030~\\AA\\ in the converging f/2.7 beam of the telescope. (The transmission curve of this filter at the ambient temperature of $11^\\circ$~C is displayed in Figure~1.) Our off-band data was composed of five 540~sec exposures through a 267~\\AA\\ wide filter centered at 5312~\\AA. In addition, an \\Halpha image, consisting of nine 900~sec exposures through a 75~\\AA\\ FWHM interference filter centered at 6606~\\AA, was obtained on 8 April 1995. These latter data were used to discriminate PN from compact H~II regions, supernova remnants, and emission associated with M87's cooling flow. The seeing for our $\\lambda 5007$ on-band survey was $1\\parcsec 2$; our \\Halpha data was taken in better than $1\\parcsec 0$ seeing. Planetary nebula candidates were identified and measured in a manner similar to that described in detail by Jacoby \\etal (1989), Ciardullo, Jacoby, \\& Ford (1989b), and Jacoby, Ciardullo, \\& Ford (1990). After spatially registering all the individual frames, we combined the on-band, off-band, and \\Halpha frames of each field, using the {\\tt imcombine} task in IRAF to reject radiation events. We then ``blinked'' the on-band [O~III] $\\lambda 5007$ sum against the off-band $\\lambda 5312$ and \\Halpha sum. Objects clearly visible on the on-band image, but absent on the off-band and \\Halpha frame were noted as possible planetaries. We confirmed these identifications by examining the candidates on each individual on-band frame, and then looking closely at the appearance of each object on our [O~III] $\\lambda 5007$ ``difference'' picture. Equatorial coordinates for the PN candidates were derived using 86 reference stars from the USNO-A.1.0 Astrometric Catalog (Monet 1996) to define the CCD's coordinate system; the internal error in these coordinates is $\\sim 0\\parcsec 5$. [O~III] $\\lambda 5007$ photometry was accomplished relative to bright field stars with the DAOPHOT point-spread-function fitting routines (Stetson 1987) within IRAF\\null. These measurements were placed on the standard system by comparing large aperture measurements of the field stars with similar measurements of the Stone (1977) and Oke (1974) spectrophotometric standards G191B2B, Feige~34, BD+25~3941, and BD+40~4032. The dispersion in the photometric zero point computed from these stars was $0.03$~mag. Finally, we computed the standard $\\lambda 5007$ magnitudes for the PN by modeling the filter transmission curve (Jacoby \\etal 1989) and using the photometric procedures for emission-line objects described by Jacoby, Quigley, \\& Africano (1987). For this computation, the systemic velocity of M87 was taken from the Third Reference Catalog of Bright Galaxies (de Vaucouleurs \\etal 1991), and the galaxy's envelope velocity dispersion was estimated from Sargent \\etal (1978). Note that since the systemic velocity of M87 is near the peak of the filter transmission curve, a $\\sim 100$~\\kms\\ error in the latter quantity translates into a flux error of only $\\sim 0.03$~mag. Table~1 lists the PN candidates identified in the field of M87, and follows on from the numbering scheme of Jacoby, Ciardullo, \\& Ford (1990). Columns~2, 3, and 4 of the table list the objects' epoch 2000 coordinates and $\\lambda 5007$ magnitudes as defined by Ciardullo \\etal (1989a), \\begin{equation} m_{5007} = -2.5 {\\rm \\,log \\,} F_{5007} - 13.74 \\end{equation} Column~5 gives the semi-major axis of the isophote upon which each PN is superposed. For $r_{\\rm iso} < 5\\parcmin 8$, these values were determined using the surface photometry of Cohen (1986); at larger distances, the isophotal radii were computed from an assumed axis ratio $b/a = 0.77$ and a galactic position angle of ${\\rm p.a} = 158^\\circ$. Table~2 lists an additional 9 PN that are projected very near other galaxies in the field and are presumably bound to them. These objects are included only for completeness and are not used in any of our analyses. Table~3 gives the mean errors in our photometric measurements as reported by the PSF-fitting algorithms of DAOPHOT\\null. However, because portions of M87 have been previously surveyed by Jacoby, Ciardullo, \\& Ford (1990), it is possible to independently assess our errors by comparing the two data sets. Of the 55~PN identified by Jacoby {\\refit et al.,} 45 were recovered in this survey. A comparison of the magnitudes of the four PN with the highest signal-to-noise ratio shows that there is no statistical difference between the magnitude system of the two surveys: the zero point of the new observations is $0.03 \\pm 0.06$~mag brighter than that from 1990. However, as Figure~2 demonstrates, a comparison of the entire dataset indicates that there is an additional source of scatter $\\sim 0.1$~mag above that expected from the combined errors of the two measurements. Part of the scatter is probably due to differences in filter transmission curves, as M87's internal velocity dispersion will shift the emission lines of some objects on the filters' wings. (We correct for this effect in the mean using the techniques outlined in Jacoby \\etal (1989) and Ciardullo, Jacoby, \\& Ford (1989b), but corrections for individual objects cannot be made without velocity information. Most of the additional error probably comes from the 1990 data, which was compromised by variable seeing and a high readout noise RCA CCD\\null. Nevertheless, for the analysis below, we have added an additional 0.07~mag uncertainty in quadrature to the errors listed in Table~3. In practice, the amplitude of the error term makes very little difference to our final results. Nine PN candidates from Jacoby, Ciardullo, \\& Ford (1990) were not detected in this survey: PN \\# 30, 35, 45, 49, 50, 51, 52, 53, and 54. In addition, PN candidate \\# 55 also was not recovered, but it fell at the position of a CCD defect, and thus could not be checked. Eight of these objects were at the limit of the previous survey and below the stated completeness limit; the other two were near the limit of completeness. All of the brighter PN from the previous survey were easily recovered in this new, wide-field survey. ", + "conclusions": "We have detected 338 planetary nebulae in our new, wide-field survey of M87 and its surrounding halo. The analysis of the luminosity function of these PN demonstrates that M87 is at a distance of $14.4 \\pm 1.1$~Mpc; this number is in excellent agreement with the earlier PNLF measurement, as well as recent distance determinations from the surface brightness fluctuation method and the globular cluster luminosity function. The result is, however, in sharp disagreement with the hypothesis of Bottinelli \\etal (1991) and Tammann (1992) that PNLF measurements in Virgo are biased due to a limited sample size. Our observations of M87 also present strong evidence for the presence of a substantial population of intergalactic stars, which extends over $\\sim 2$~Mpc in front of M87. The intergalactic stars are, in all likelihood, the explanation for the ``overluminous'' planetaries encountered by Jacoby, Ciardullo, \\& Ford (1990) in Virgo, but not seen any where else (Jacoby, Ciardullo \\& Harris 1996). The analysis of the intracluster PN luminosity function suggests that the Virgo Cluster core is not virialized, but is instead dynamically young. However, the mismatch between the width of our [O~III] $\\lambda 5007$ filter and the Virgo Cluster velocity dispersion precludes a definitive statement or a more detailed analysis at the present time. Planetary nebula observations offer a new opportunity for morphological and dynamical studies of nearby clusters. To date, the only way to study intergalactic light in clusters has been through deep surface photometry, and as a result, only a few, very rich, Abell Clusters have been measured (cf.~Uson, Boughn, \\& Kuhn 1991; V\\'ilchez-G\\'omez, Pell\\'o, \\& Sanahuja 1994). Planetary nebula observations offer an alternative to these studies, and provide information on both the two and three dimensional structure of the cluster. Moreover, a PN's [O~III] $\\lambda 5007$ emission line is an excellent target for a radial velocity measurement. Since there are never enough sufficiently bright galaxies in a cluster to fully define the cluster's velocity field, planetary nebulae can provide invaluable data for cluster dynamics. Future [O~III] $\\lambda 5007$ surveys may therefore allow us to probe the effects of galactic mergers, cluster accretion, and tidal-stripping within several nearby clusters, and enable new investigations into the distribution of dark matter in clusters and of the initial conditions of cluster formation. We would like to thank Mike Pierce for first suggesting the possible intracluster origin of some of M87's halo PN\\null. This work was supported in part by NASA grant NAGW-3159 and NSF grant AST-9529270. Support for this work was also partially provided by NASA through grant number GO-0612.01-94A from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555." + }, + "9708/astro-ph9708170_arXiv.txt": { + "abstract": "In the gravitational lens system Q2237+0305 the cruciform quasar image geometry is twisted by ten degrees by the lens effect of a bar in the lensing galaxy. This effect can be used to measure the mass of the bar. We construct a new lensing model for this system with a power-law elliptical bulge and a Ferrers bar. The observed ellipticity of the optical isophotes of the galaxy leads to a nearly isothermal elliptical profile for the bulge with a total quasar magnification of $16^{+5}_{-4}$. We measure a bar mass of $7.5\\pm 1.5\\times10^{8}\\,{\\rm h}_{75}^{-1}{\\cal M}_{\\sun}$ in the region inside the quasar images. ", + "introduction": "Gravitational lensing provides a unique way to weigh objects at cosmological distances without any assumption about the connection between light and dark matter. Since the discovery of the first gravitational lens (Walsh, Carswell \\& Weymann 1979) several gravitational lenses have been found and this method has been used many times to explore the mass distribution of galaxies. In this paper, we model the lens system Q2237+0305 in order to weigh the bar in the lensing galaxy. The quasar Q2237+0305 ($z_{q} = 1.695$) was found by Huchra {\\etal} \\shortcite{Huchra1985} at the centre of an SBb spiral galaxy ($z=0.0394$) that is situated in the outskirts of the Pegasus~II cluster. The quasar was later resolved into four images that are situated around the core of the galaxy within a radius of one arcsecond \\cite {Yee1988,Schneider1988}. Two fundamentally different approaches have been used to model the lensing galaxy. One was to fit a parametric mass profile with several free parameters to the observed quasar image configuration~\\cite{Kent1988}; the other to use a model of the light distribution of the galaxy and to fit for the mass-to-light ratio as the single free parameter \\cite {Schneider1988,Rix1992}. The former approach was naturally much more precise in the reproduction of the observed image geometry due to the greater number of free parameters. The length scale over which the lensing galaxy influences a light bundle from the quasar is small compared to the cosmological distances between observer, lens and source. The lens can therefore be treated as a mass sheet at the position of the galaxy. Since the galaxy disk of 2237+0305 is inclined with respect to the sky, elliptical surface mass distributions must be used in the models of this system. Interestingly, the position angle, counted counterclockwise from north, of the major axis of the elliptical lens models found by Kent \\& Falco \\shortcite {Kent1988} was about $67\\degr$. This is almost parallel to the axis through images C and D and just between the angle of the inclination axis of the galaxy ($77\\degr$, Yee 1988) and the angle of the bar ($39\\degr$, also Yee 1988). This situation is shown in figure~\\ref {ImageGeometry}.\\begin{figure} \\plotone{figure1.ps} \\caption{Illustration of the image geometry of Q2237+0305, motivated by figure~1 in Kent \\& Falco \\protect\\shortcite {Kent1988}. The images are labelled using the convention by Yee \\protect\\shortcite {Yee1988} with positions from Crane {\\etal} \\protect\\shortcite {Crane1991}. The relative areas of the circles correspond to the radio flux ratios from Falco {\\etal} \\protect\\shortcite {Falco1996}. The position of the galaxy centre is indicated with a cross. The long arrows indicate the directions of the galaxy inclination axis and the bar, as well as the position angle of $67\\degr$. North is up, East to the left.} \\label{ImageGeometry} \\end{figure} This was also found by Kochanek \\shortcite {Kochanek1991} and Wambsganss \\& Paczy\\'{n}ski \\shortcite {Wambsganss1994} who used simple circular mass distributions with an additional quadrupole perturbation. When fitted to the observed image geometry, the direction of the perturbation turned out to be close to the one Kent \\& Falco \\shortcite {Kent1988} found for their model major axis. More recent investigations by Witt, Mao \\& Schechter \\shortcite {Witt1995}, Kassiola \\& Kovner \\shortcite {Kassiola1995} and Witt \\shortcite {Witt1996} obtained the same position angle for the perturbation or major model axis. On the other hand, the bar shows up prominently in CCD images of the galaxy. It has been noted~\\cite {Tyson1985,Yee1988,Foltz1992} that it might contribute significantly to the lensing in the system -- initially this was actually an aid to explain the lensing effect when only the images A, B and C were known~\\cite {Tyson1985}. The motivation for this paper is the idea that the apparent misalignment of the predicted model major axis and observed galaxy inclination axis as shown in figure~\\ref {ImageGeometry} is due to the lensing influence of the bar. In section~\\ref {TheoreticalModel} we construct and analyse a lensing model that includes the bar component and takes the observed position angle for the inclination axis of the galaxy into account. Section~\\ref {PropertiesOfTheBar} deals with the implications of this model for the bar. In section~\\ref {Discussion} we finally discuss our results. We use a cosmological model with $H_0 = 75\\,{\\rm h_{75}}\\, {\\rm km\\, s^{-1}\\, Mpc^{-1}}$, $\\Omega=1$ and $\\Lambda=0$. ", + "conclusions": "\\label{Discussion} In this paper we have presented a barred galaxy model for the gravitational lens 2237+0305. We used a power--law elliptical mass distribution for the bulge and chose the model for which the observed ellipticity is predicted. It turned out that this model has an exponent close to unity, which is compatible with other determinations of lens mass profiles through gravitational lensing \\cite {Kochanek1995,Grogin1996}. The bar represents a small perturbation of the deflection field of the bulge of the galaxy, amounting to $7.5\\pm 1.5\\times10^{8}\\,{\\rm h}_{75}^{-1}{\\cal M}_{\\sun}$ or about 5\\%~of the bulge mass in the critical region inside the quasar images. The relative magnifications our model predicts for the quasar images can be compared with the 3.6cm radio flux ratios published by Falco {\\etal} \\shortcite {Falco1996}. Their results are also given in table~\\ref {ParametersofModels}. Falco {\\etal} argue that it is unlikely that the radio flux densities are variable from microlensing due to the larger size of the radio emitting region as compared to the optical continuum emitting region, although they cannot completely rule out microlensing as an important effect in the radio. If microlensing is not important, the model magnification ratios should be identical to these measured flux ratios. It can be seen that only the ratio $\\mu_{\\rm DA}$ between images D and A is not compatible with their results although no effort was made to fit the flux ratios. In order to find out more about the discrepancy of the ratio $\\mu_{\\rm DA}$ between observation and model one has to make the relatively large error bars from Falco {\\etal} smaller through longer radio observation of the object. Unfortunately, Q2237+0305 has a radio flux density of only $\\approx 1$~mJy \\cite {Falco1996}, so that radio observations of this object are very time-consuming; Falco {\\etal} observed for 11 hours of which only five could be used eventually due to weather conditions. To get additional, independent arbiters for the model, it would be very helpful to measure the time--delays in this system. Since the time--delays are of the order of several hours, this has to be done in a wavelength domain where the necessary intra-day variability is likely to occur for a radio-quiet quasar, for example in the x-ray regime as proposed by Wambsganss \\& Paczy\\'{n}ski \\shortcite {Wambsganss1994}. Also, monitoring in the radio would show if the quasar image flux densities vary at these frequencies. Unless we learn more about the radio flux densities and the time delays, it is not possible to decide whether or not it is microlensing that causes the low magnification of image D. In the optical, image D has always been the faintest quasar image. In fact, in the first resolved image of the quasar, image D was not visible at all \\cite {Tyson1985}. There is also spectroscopic evidence from optical data that image D is undergoing demagnification \\cite {Lewis1996}. If image D is in fact microlensed in the radio, the consequences are interesting. The scale size of the radio region could be less than the characteristic scale of the caustic network. Alternatively, the radio source could have an asymmetric structure like a jet that would have differing microlensing properties for different paths of the microlenses across the source. Yet another way to significantly change the radio flux density of image D would be a globular cluster or black hole \\cite {Lacey1985} with a mass of about $10^6 {\\cal M_{\\sun}}$ in the halo of the lensing galaxy that is situated close to image D. An object of this mass would magnify or demagnify the radio image of the quasar, depending on its location with respect to the direction of the local shear. This effect, the perturbation of lens models by $10^6 {\\cal M_{\\sun}}$ objects, has recently been treated by Mao \\& Schneider \\shortcite {Mao1997}. With this, we can estimate that a surface mass density of globular clusters or black holes of approximately $0.04\\,\\Sigma_{\\rm crit}$ ($\\Sigma_{\\rm crit}$ is defined in section~\\ref {TheBulge}) or $470\\,{\\rm h}_{75} {\\cal M}_{\\sun}/{\\rm pc}^2$ is needed to observe a demagnification of image D by $40\\%$ or more with a probability of 20\\%. Higher surface mass densities would make it more likely. This is much more than the globular cluster surface mass density of about $1\\,{\\cal M_{\\sun}}/$pc$^2$ seen in our Galaxy within 5~kpc of the Galactic centre \\cite {Mao1997}. It thus seems unlikely that the demagnification is due to a globular cluster. The question of the existence of such a massive object near image D could be solved with a method that was proposed by Wambsganss \\& Paczy\\'{n}ski \\shortcite {Wambsganss1992}. They showed that these objects would bend or even create holes in the radio maps of milliarcsecond jets of gravitationally lensed quasars. If we could observe extended structure of the images of the radio-weak Q2237+0305, features due to globular cluster or black hole lensing could be easily identified because they would not be seen in the other gravitationally lensed images of the quasar. It has recently been suggested (Keeton, Kochanek \\& Seljak 1996; Witt \\& Mao 1997) that a strong lensing perturbation is necessary for a number of lenses in addition to an elliptical mass distribution in order to model the observed image geometry. In our model for Q2237+0305, the additional perturbation from the bar is small, but the system is, nevertheless, unique in that we can see the perturbing agent. For the lens systems mentioned in these papers, the better fit for the models was obtained with an additional external shear. We saw, when we compared our results with Wambsganss \\& Paczy\\'{n}ski \\shortcite {Wambsganss1994}, that the predictions for magnifications or time--delays from shear models differ by up to a factor of two from the predictions from elliptical mass deflectors. If one wants to get lensing models with reliable predictions for magnifications or time--delays, the necessary perturbations will ultimately have to be generated in the models by mass components like dark matter haloes \\cite {Keeton1996} or bars." + }, + "9708/astro-ph9708058_arXiv.txt": { + "abstract": "We simulate the growth of large-scale structure, for 3 different cosmological models, an Einstein-de Sitter model (density parameter $\\Omega_0=1$), an open model ($\\Omega_0=0.2$) and a flat model with nonzero cosmological constant ($\\Omega_0=0.2$, cosmological constant $\\lambda_0=0.8$), using a cosmological N-body code ($\\rm P^3M$) with $64^3$ dark matter particles in a comoving cubic volume of present comoving size 128 Mpc. The calculations start at $z=24$ and end at $z=0$. We use the results of these simulations to generate distributions of galaxies at the present ($z=0$), as follows: Using a Monte-Carlo method based on the present distribution of dark matter, we located $\\sim40000$ galaxies in the computational volume. We then ascribe to each galaxy a morphological type based on the local number density of galaxies in order to reproduce the observed morphology-density relation. The resulting galaxy distributions are similar to the observed ones, with most ellipticals concentrated in the densest regions, and most spirals concentrated in low-density regions. By ``tying'' each galaxy to its nearest dark matter particle, we can trace the trajectory of that galaxy back in time, by simply looking at the location of that dark matter particle at earlier time-slices provided by the N-body code. This enables us to reconstruct the distribution of galaxies at high redshift, and the trajectory of each galaxy from its formation epoch to the present. We use these galaxy distributions to investigate the problem of morphological evolution. Our goal is to determine whether the morphological type of galaxies is primarily determined by the initial conditions in which these galaxies form, or by evolutionary processes (such as mergers or tidal stripping) occurring after the galaxies have formed, and eventually altering their morphology, or a combination of both effects. Our main technique consists of comparing the environments in which galaxies are at the epoch of galaxy formation (taken to be at redshift $z=3$) with the environment in which the same galaxies are at the present. Making the null hypothesis that the morphological types of galaxies do not evolve, we compare the galaxies that form in low density environments but end up later in high density environment to the ones that form also in low density environment but remain in low density environment. The first group contains a larger proportion of elliptical and S0 galaxies than the second group. We assume that the initial galaxy formation process cannot distinguish a low density environment that will always remain low density from one that will eventually become high density. Therefore, these results are absurd and force us to discard the null hypothesis that morphological evolution does not occur. Our study suggests that $\\sim75\\%$ of the elliptical and S0 galaxies observed at present formed as such, while the remaining $\\sim25\\%$ of these galaxies formed as spiral galaxies, and underwent morphological evolution, for all three cosmological models considered (the percentages might be smaller for elliptical than S0 galaxies). These numbers assume a morphological evolution process which converts one spiral galaxy into either a S0 or an elliptical galaxy. If the morphological evolution process involves mergers of spiral galaxies, these numbers be would closer to $85\\%$ and $15\\%$, respectively. We conclude that most galaxies did not undergo morphological evolution, but a non-negligible fraction did. ", + "introduction": "\\subsection{Morphological Types} Galaxies exist in several forms, elliptical, lenticulars, spirals, and irregulars, usually referred to as {\\it morphological types}. Elliptical galaxies are featureless, ellipsoidal stellar systems composed of old Population II stars, with no appreciable amount of cold gas or dust. In addition, many of them are known to contain also a disk. Ellipticals galaxies are labeled as E0, E1, and so on, according to their ellipticity. Lenticular galaxies have a prominent, featureless disk, that contains no appreciable amount of cold gas or dust, and no spiral arms. They are very similar to the most elongated, E7 elliptical galaxies. These galaxies are labeled as S0. Spiral galaxies are composed of a disk of Population~I stars, cold gas, and dust, arranged in a pattern of spiral arms, and a central bulge of population~II stars which resemble small elliptical galaxies. The spiral arms are the site of active star formation, and contain a large number of young O and B stars. Spiral galaxies have flat rotation curves that extend to radii well beyond the visible edge of the galaxy, thus implying that these galaxies are imbedded into large dark matter halo. Spiral galaxies are labeled as Sa, Sb, Sc, and Sd galaxies according to their disk-to-bulge luminosity ratio (D/B), with the bulge dominating the luminosity for Sa galaxies, and the disk dominating for Sd galaxies. Galaxies that do not belong to any of these categories are classified as irregular galaxies. Some irregular galaxies result from collision and merging between galaxies, but the majority of irregular galaxies are small, gas rich galaxies similar to the Magellanic clouds. We label these galaxies as Im. All the galaxy types described above can be combined into a single sequence, $\\rm E0\\rightarrow E1\\rightarrow\\ldots\\rightarrow E7 \\rightarrow S0\\rightarrow Sa\\rightarrow Sb \\rightarrow Sc\\rightarrow Sd\\rightarrow Im$, called the {\\it Hubble sequence} \\footnote{The Hubble sequence is actually a ``tuning fork'' with two branches, one for {\\it unbarred} spirals and one for {\\it barred} spirals. In this paper, we ignore the difference between barred and unbarred spirals, thus collapsing the tuning fork into a rod. Hence, ``Sa'' designates both Sa and SBa galaxies, and so on.}. Near the start of the sequence, galaxies are mostly composed of old Population~II stars, with no dust and no cold gas, and therefore no active star formation, and a small disk-to-bulge ratio. As we move along the sequence, the preponderance of Population~II stars decreases in favor of younger, Population~I stars. The amount of dust and cold gas increases, D/B increases, and star formation becomes important. A successful theory of galaxy formation must be able to explain the existence of the Hubble sequence, the origin of each morphological type, their relative abundance, and their clustering properties. To achieve this goal, we must first identify and understand the physical processes that are involved in the galaxy formation process, as well as the processes that might subsequently alter the structure of galaxies after they are formed. The most important clue for understanding the galaxy formation process is the existence of a {\\it Morphology-Density Relationship} relating the likelihood of any given galaxy to have a particular morphological type to the {\\it local} density of the environment in which that galaxy is located. \\subsection{The Morphology-Density Relation at Present} There is a significant difference between the galaxy populations of nearby low-density fields and in the densest regions inside nearby clusters of galaxies. Though all morphological types are present both in clusters and in the field, field galaxies are predominantly spirals, while clusters of galaxies contain a much larger proportion of elliptical and S0 galaxies. Furthermore, population gradients are found inside clusters. Melnick \\&~Sargent (1977) showed that the proportion of spiral galaxies increases as a function of the distance from the cluster center, with a corresponding decrease in the proportion of S0 and elliptical galaxies. Dressler (1980) argued that this morphology-radius relation is applicable only to regular, spherical clusters with a well-defined center. Most clusters are highly irregular, and often contain several high density concentrations, or lumps. The distribution of the various morphological types inside these lumps is similar to the one in the center of the regular, spherical clusters. Dressler (1980) concluded that the correct way to describe the distribution of morphological types is in terms of the local number density of galaxies, and not the distance from the cluster center. Using a sample of 55 rich clusters, he showed that the fraction of elliptical and S0 galaxies increases with increasing surface number density of galaxies, with a corresponding decrease in the fraction of spiral galaxies, over 3 orders of magnitude in surface number density. The lowest density regions in the sample are composed of 80\\% spirals, 10\\% S0's, and 10\\% ellipticals, while the densest clumps are composed of 10\\% spirals, 40\\% S0's, and 50\\% ellipticals. Subsequent studies (Bhavsar 1981; de~Souza et al 1982; Postman \\& Geller 1984) confirmed the relations derived by Dressler (1980), and extended them to the low-density field. All these results are summarized in Dressler (1984). The morphology-density relation extends over 5 orders of magnitude in volume number density (Postman \\&~Geller claim 6 orders of magnitude), and is a slowly varying, monotonic relation. The lowest-density regions are composed of 80--90\\% spirals, while the highest-density regions are composed of 80--90\\% ellipticals and S0's. (Notice that a recent paper by Whitmore, Gilmore, \\&~Jones [1993] challenges the existence of the morphology-density relation, and claims that the morphology-radius relation is actually the correct one.) Notice that these various determinations of the morphology-density relation were all based on observations of relatively nearby galaxies. Therefore, this relation is valid only {\\it at present}. More recent observations with the Hubble Space Telescope (HST) suggest that the morphology-density relation evolves with time, and this actually supports the results we present in this paper. A discussion of the HST results is presented in \\S9. \\subsection{The Origin of the Morphological Types} Several galaxy formation models have been suggested to explain the origin of the Hubble sequence and the existence of the morphology-density relation. Dressler (1984) has grouped these various models into three classes, based on the relative importance of initial conditions and evolution processes in determining the final morphological type of galaxies. We shall follow the same classification here. \\subsubsection{Morphological Evolution} Models that belong to the first class all assume that galaxies form in similar environments, and therefore the existence of different morphological types does not result from different initial conditions, but instead from evolutionary processes happening after the galaxies have formed. Several models have been suggested to explain the abundance of S0 galaxies and deficiency of spiral galaxies in dense regions. These models all assume that S0 galaxies are spiral galaxies that have lost their gas and dust as a result of some evolutionary process taking place in the dense environments of cluster cores. The various possible physical mechanism for gas stripping include direct collisions (Spitzer \\&~Baade 1951) ram-pressure stripping (Gunn \\&~Gott 1972) and gas evaporation by a hot intracluster gas (Cowie \\& Songaila 1977). Dressler (1980) pointed out a major problem with these models: the various physical mechanisms suggested are efficient only in the densest regions, inside cluster cores. Though the {\\it fraction} of S0 galaxies is largest in these regions, the {\\it actual number} of S0's galaxies in these regions is small. About 80\\% of S0 galaxies are located in intermediate-density regions. Spiral galaxies in intermediate-density regions are deficient in gas by a factor of 2-3 relative to field spirals, indicating that gas ablation is important in these regions as well (Giovanelli, Chincarini, \\& Haynes 1981; Bothun, Schommer, \\&~Sullivan 1982; Kennicutt 1983). However, this effect is much too weak to explain the presence of S0 galaxies, which are gas deficient by a factor of 100 relative to field spirals. \\subsubsection{Initial Conditions Combined with Morphological Evolution.} The second class of models comprises all models in which both initial conditions and morphological evolution play an important role in determining the morphological types of galaxies. Kent (1981) had suggested that the morphology-density relation originates from the ``fading'' of disks in high density regions. In this model, initial conditions are assumed to be responsible for determining the initial morphological type of disk galaxies, such that disk galaxies with large D/B become predominantly spirals, while disk galaxies with small D/B become predominantly S0's. The model then assumes that the disks of spiral and S0 galaxies are fainter in high density regions than in low density region (this could result from the dissipation of the disk by tidal interaction, or, if the disks are still in the process of forming, then a large density environment might disrupt this process). The fading of disks causes some spiral galaxies to become too faint to be observable, and others to be identified as S0 galaxies. Furthermore, the fading of the disk of S0 galaxies causes some of these galaxies to be identified as ellipticals. With an appropriate choice of parameters, this model can successfully reproduce all the relations given in Dressler (1980). Larson, Tinsley, \\&~Caldwell (1980) have proposed a similar model, based on the time scale for gas exhaustion via stellar evolution in disks. In their model, the gas exhausted by star formation is constantly replaced by gas infalling from a gaseous envelope surrounding the galaxy. In high-density regions, tidal encounters would disrupt this envelope, resulting in a progressive fading of the disk as stellar evolution proceeds. The various gas-stripping processes mentioned in \\S1.3.1 could be responsible for transforming spirals galaxies into S0's inside cluster cores (even though they cannot account for the existence of field S0 galaxies). Byrd \\&~Valtonen (1990) have argued that the interaction of spiral galaxies with the tidal field of the cluster is a more efficient process than ram pressure stripping in depleting these galaxies of their interstellar gas, and eventually turning them into S0 galaxies (but not ellipticals). Their model is supported by the abundance of barred spiral galaxies in the core of the Coma cluster, since the formation of a bar in a normal spiral galaxy can also result from strong tidal interaction. If the galactic disks are ``faded'' in high density regions, as these models assume, then the luminosity function inside dense clusters should differ significantly from the one in low density clusters and in the field. However, observations show that the luminosity functions in low- and high-density regions are essentially identical (Dressler 1984, and references therein), though Biviano et al. (1995) recently suggested that this might not be the case for the Coma cluster. In order to maintain the luminosity function unchanged in high-density regions, any ``fading'' of the disk must then be accompanied by a corresponding brightening of the bulge. Mergers could be responsible for building up large galactic bulges in high-density regions. It has been suggested that elliptical galaxies result from the merging of spiral galaxies (Toomre \\&~Toomre 1972; Toomre 1977; White 1978; 1979; Fall 1979). Ross (1981) has suggested that galaxies form mainly as stellar disks, and that galactic bulges are formed by merging, for all galaxy types. This could explain the fact that the angular momenta of disk and bulge in disk galaxies are almost perfectly aligned (Gerhard 1981). Numerical simulations of galaxy mergers by Mihos \\&~Hernquist (1994a, 1994b) support this model, by showing that mergers trigger infall of material toward the center of the system. This model, if correct, would explain the abundance of {\\it both} S0 and elliptical galaxies in high-density regions. Numerical simulations (Efstathiou \\&~Jones 1979; Aarseth \\& Fall 1980) have shown that mergers of galaxies on highly eccentric orbits result in the slow-rotating systems, consistent with measurements of the spin parameter for elliptical galaxies. Merging events, however, are not expected to occur inside rich clusters, where most ellipticals are found. The velocity dispersion in these regions is quite high, resulting in a significant reduction of the gravitational cross sections of galaxies. More likely, mergers occur inside small groups of galaxies where the velocity dispersion is smaller, and later these groups assemble into clusters (see, e.g., Geller \\&~Beers 1982). Numerical simulations (Aarseth \\&~Fall 1980; Negroponte \\&~White 1983; Noguchi 1988; Barnes 1989; Barnes \\&~Hernquist 1991; 1992) show that galaxy mergers occur naturally inside small groups, and that such mergers result in the formation of spheroidal galaxies with essentially no disk (Barnes \\&~Hernquist 1992). Baugh, Cole, \\& Frenk (1996) have used a semi-analytical, Monte Carlo model to describe galaxy mergers in a standard Cold Dark Matter (CDM) universe. Their model produces a distribution of D/B which are consistent with observations, when the values of D/B are used to ascribe morphological types. Moore et al. (1996) have suggested that morphological evolution of spirals actually occurs inside dense clusters, in spite of the large velocity dispersion. In their model, called ``galaxy harassment,'' spiral galaxies are disrupted by the cumulative effect of several high velocity close encounters with other galaxies. The various studies of mergers described above consider the merging of two or more galaxies of comparable size. A completely separate problem is the merging of a disk galaxy with a satellite galaxy of much smaller mass. These merging events can modify the structure of the disk, but the effect is too small to result in actual morphological evolution (that is, spiral galaxies will remain spiral after ``swallowing'' a satellite). Numerical simulations (Quinn \\&~Goodman 1986; Quinn, Hernquist, \\& Fullagar 1993; T\\'oth \\&~Ostriker 1992) have shown that a merger between a disk galaxy and a satellite having a mass equal to 1/10 of the mass of the disk results in a important thickening of the disk, which is ruled out by observations. However, these simulations ignored the possibility that the satellite might dissolve significantly before the actual merging takes place. More recent simulations (Carlberg 1995; Huang 1995) have suggested that the main effect of these mergers is a tilt of the disk, accompanied by a transient warp, with no substantial thickening. There are several problems with models involving mergers. Elliptical and spiral galaxies have different globular cluster luminosity functions (Harris 1981). Since merging events would unlikely affect the structure of globular clusters, this result argues against elliptical galaxies being formed from the merging of spirals, {\\it if the number of globular clusters remains constant during the merging process}. However, Ashman \\& Zepf (1992) have argued that the merging of two galaxies results in the formation of additional globular clusters. Also, dwarf ellipticals presumably {\\it do not} result from mergers, so the continuity of properties for dwarf ellipticals to regular ellipticals (Sandage 1983) suggests that large elliptical do not result from mergers either. Merging events would most likely ruin tight correlations existing among various parameters for elliptical galaxies, such as color and luminosity (Bower, Lucey, \\& Ellis 1992) and effective radius, central velocity dispersion, and mean surface brightness (the ``fundamental plane,'' Djorgovski \\& Davis 1987; \\hbox{J\\o rgensen}, Franx, \\& \\hbox{Kj\\ae rgaard} 1996). Another possible problem is that stars are much more strongly concentrated in elliptical galaxies than in spirals (Combes et al. 1995). However, recent SPH simulations of galaxy mergers (Steinmetz 1995; Barnes \\& Hernquist 1996 and references therein) show that the merger of two spiral galaxies often results in the formation of much denser systems, sometimes too dense to be ellipticals galaxies. \\subsubsection{Initial Conditions} The third class of models comprises models in which the initial conditions are primarily responsible for determining the morphological type of galaxies, with subsequent morphological evolution playing little role or no role at all. Numerous models have been proposed (see Dressler 1984, and references therein), in which the morphological type is determined either by the local density, or the local amount of angular momentum. Such models could successfully explain the observed morphology-density relation only if galaxies have formed near their present location. This could be the case in cosmological models which have more power at large scale than small scale. In such models, clusters would form first, and then fragment into individual galaxies, in which case galaxies could actually be located at present near the location were they where formed. The alternative is that galaxies, at the epoch of their formation, somehow ``know'' the kind of environment in which they will be located at the present. This can be achieved if there is some kind of coupling between the perturbations responsible for forming the galaxies and the ones responsible for forming the clusters in which these galaxies end up. The problem with these scenarios is that they all invoked cosmological models that are usually considered ``marginal.'' These models constitute interesting alternatives to the more standard CDM model with Gaussian initial conditions, but there is at present no strong, absolute evidence favoring such models over the standard ones. To our knowledge, the most serious alternatives, at present, to the standard CDM models are the models with Cold + Hot Dark Matter (CHDM), models with a nonzero cosmological constant, and models with a tilted power spectrum. None of these models feature coupling between long- and short-wavelength modes in their initial conditions. Hence, following Dressler (1984), we will regard these types of galaxy formation models as ``last resort.'' \\subsection{Past History of Galaxies} In order to identify the correct galaxy formation model, we must reconstruct the past history of the presently observed galaxies, and in particular we need to know the kind of environment in which galaxies were located at various epochs. We are assuming that galaxies at their formation epoch have no knowledge of the future environment in which they will end up. We are therefore rejecting all ``class three'' models, unless galaxies form near their present location. Hence, if we find that elliptical and S0 galaxies located in the dense cluster cores were always located in high density environment, at all epochs up to the galaxy formation epoch (redshifts $z$ of order 3--5), it would argue in favor of the initial conditions being responsible for determining the morphological type (class three models), and against morphological evolution. If, to the contrary, many of these elliptical and S0 galaxies are found at early time in low density environments, it would argue in favor of morphological evolution (class one or two models). The goal of this paper is to settle this question. ", + "conclusions": "We conclude that a small but non-negligible fraction (of order 10\\%--20\\%) of the S0 and elliptical galaxies we observe today in the dense parts of clusters were not formed as S0's and ellipticals, but rather as spiral galaxies, and underwent morphological evolution between $z=3$ and $z=0$, presumably during cluster formation and merging. Since the fraction of galaxies involved in morphological evolution is neither 0\\% nor 100\\%, initial conditions and morphological evolution processes must {\\it both} play an important role in determining the morphological type of galaxies. Our simulations predict that the proportion of spiral galaxies should increase form the present observed value of $\\sim50\\%$ to larger values as one looks back in time, that is, at larger redshifts. However, they cannot predict at what redshift this effect would manifest itself, and consequently we cannot predict the shape of the morphology-density relation at high redshift. To make a theoretical prediction, we need first to understand the details of the morphological evolution process. Also, the epoch of galaxy formation most certainly depends upon the cosmological model, so before we can make quantitative predictions, we first need to settle the question of which cosmological model properly describes the formation of large-scale structures in the universe. However, a large amount of observational evidence supporting the existence of morphological evolution in dense environments at redshifts $z<0.5$ has been accumulated in recent years. Butcher \\& Oemler (1978, 1984) discovered a large excess of blue objects in clusters located at redshift $z\\gtrsim0.4$. Subsequent ground-based observations (Dressler \\&~Gunn 1982, 1983; Couch et al. 1983; Couch \\& Newell 1984; Dressler, Gunn, \\&~Schneider 1985; Ellis et al. 1985; Lavery \\&~Henry 1986; Henry \\&~Lavery 1987; Couch \\&~Sharples 1987; MacLaren, Ellis, \\&~Couch 1988; Soucail et al. 1988; Arag\\'on-Salamanca, Ellis, \\&~Sharples~1991; Arag\\'on-Salamanca et al. 1993) have shown that this ``Butcher-Oemler effect'' results from short-lived bursts of star formation affecting a subset of the cluster members. These starbursts could be triggered by the ram pressure of the intracluster gas when a galaxy first enters the cluster, by violent interaction between galaxies, or by mergers, (see Bothun \\&~Dressler 1986, and references therein; Oemler 1992, and references therein; Mihos \\&~Hernquist 1994a, 1994b). Recent Hubble Space Telescope observations of high redshifts clusters $z\\sim0.3-0.5$ revealed that the blue starburst objects are low luminosity spiral galaxies, with as many as $\\sim50\\%$ of them being disturbed by what appears to be either tidal disruption or merging (Dressler et al. 1994a, 1994b; Couch et al. 1994; Barger et al. 1995). The galaxy populations of these clusters differ significantly form the ones of nearby clusters, and resemble the ones seen in the nearby small groups and field. The difference between the galaxy populations of high-redshift and low-redshift clusters and the importance of dynamical interaction in high-redshift clusters compared to low-redshift ones provide strong evidence that morphological evolution has occurred inside rich clusters. Studies of galaxy populations in the field (Colless et al. 1990; Griffiths et al. 1994; Mobasher et al. 1996) and in small groups (Allington-Smith et al. 1993), reveal that no significant morphological evolution has occurred in these environments between redshift $z=0.5$ and the present, at least among luminous galaxies. (Driver et al. [1995], however, found an excess of {\\it faint} late type galaxies in the field.) These results rule out any model in which the morphological evolution of a galaxy is driven by an internal physical process. The morphological evolution process depends upon the richness of the environment, and thus results in a steepening of the morphology-density relation with time. The most recent studies (Dressler \\& Smail 1996; Smail et al. 1997; Dressler et al. 1997, and references therein) of high-redshift clusters, which include 10 rich clusters ($0.36 40 M_{\\odot}$) evolve away from the Zero Age Main Sequence their photospheres become decreasingly stable against radiation pressure and ultimately reach a critical point where the radiation pressure and gravity are balanced, leading to large mass loss and ending the redward evolution. Because of metal line opacity, the luminosity at which a stellar photosphere becomes unstable is much lower than that predicted by the classical electron-scattering Eddington limit. Quantitative studies of a ``modified Eddington limit'' were performed by Lamers \\& Fitzpatrick (1988, hereafter LF) using low gravity, line-blanketed, plane-parallel, LTE model atmosphere calculations. By extrapolating from the low gravity models to a point at which radiation pressure balanced gravitational pressure, they determined that the modified Eddington limit was in reasonable agreement with the observed upper luminosity limit for hot stars ($>10,000$~K) in the LMC. Lamers \\& Noordhoek (1993, hereafter LN) extended this work to examine the metallicity dependence of the modified Eddington limit; Achmad, de~Jager, and Nieuwenhuijzen (1993) found that cool supergiants ($<10,000$~K) are observationally excluded from the region of luminosity/temperature space predicted to be unstable from the modified Eddington limit approach (\\cite{gus92}). Alternate explanations for the HD-limit include instabilities of radial modes in massive stars (\\cite{gla93}; \\cite{kir93}); turbulent pressure (e.g., de Jager 1984); and binary star models (e.g., Kenyon \\& Gallagher 1985). Humphreys \\& Davidson (1994) critically review all these proposed instability mechanisms and conclude that none, at least in the current state of development, is fully satisfactory. It is important to understand the nature of the mass loss and instability mechanisms operating in the upper HR diagram --- not only to complete the theoretical picture of stellar evolution, but also to aid in the interpretation of observations of massive stars. Perhaps the most obvious application of such an understanding would be to determine whether the brightest stars can be used as reliable distance indicators (e.g., \\cite{hum87}). In this paper, we revisit the modified Eddington limit scenario proposed by Lamers and collaborators. We utilize up-to-date stellar atmosphere calculations to evaluate the radiation pressure stability of low surface gravity stars and the most recent stellar evolution calculations to transform the stellar atmosphere parameters ($T_{\\rm eff}$ and $\\log g$) to the HR diagram ($T_{\\rm eff}$ and $L$). The model atmosphere calculations and the transformation to the HR Diagram are described in \\S~\\ref{lowgsec}. In \\S~\\ref{comobssec}, the modified Eddington limit is compared with the observed upper HR Diagram of the LMC. Concluding remarks are given in \\S~\\ref{concomsec}. ", + "conclusions": "} In summary, we have determined the location of the modified Eddington limit for stars in the LMC using the most recent atmosphere models combined with a precise mapping to the HR Diagram through up-to-date stellar evolution calculations. We find that the modified Eddington limit is actually {\\it a full magnitude higher} than the upper luminosity limit observed for LMC stars. The observed limit is consistent with atmosphere models in which the maximum value of the ratio of the radiation force outwards to the gravitational force inwards, $Y_{\\rm max}$, is 0.9; i.e., the photospheres of stars at the observed luminosity limit are bound. With some caution, we thus suggest that the simple picture in which a massive star evolves redward until its photosphere reaches the modified Eddington limit and becomes unbound is invalid. Although the stars do evolve from the Zero Age Main Sequence in the direction of increasing $Y_{\\rm max}$, an instability evidently sets in {\\it before} the atmospheres reach the formal modified Eddington limit at $Y_{\\rm max} = 1.0$. This conclusion is necessarily tentative since this analysis, like others before, relies on plane-parallel, hydrostatic atmosphere models, while the atmospheres of real stars near the observed luminosity limit are likely to share neither of these properties. It appears unlikely, however, to be a coincidence that the temperature dependence of the luminosity limit should so closely match that of the $Y_{\\rm max}$ curves seen in Figures 1--4, whose shapes are nearly invariant to metallicity or to the precise value of $Y_{\\rm max}$ itself. The degree of stability against radiation pressure of the photospheres clearly plays an important role in shaping the upper stellar luminosity limits, although the current characterization of that stability may leave something to be desired. The $Y_{\\rm max}$ parameterization may well turn out to correlate with some more critical property, such as the depth of the ``boundary'' between a stellar wind and the underlying photosphere. A firm understanding of the upper luminosity limits and of the outbursts in LBVs will almost certainly require a melding of stellar wind, stellar photosphere, and stellar evolution calculations. Fortunately, progress in this area is being made (e.g., \\cite{sel93}, Schaerer~et~al.~1996)." + }, + "9708/astro-ph9708046_arXiv.txt": { + "abstract": "We present the results from a detailed analysis of the 0.6 - 10 keV spectra of 23 {\\it ASCA} observations of 18 objects. We find that in most cases the underlying continuum can be well-represented by a powerlaw with a photon index $\\Gamma \\sim 2$. However we find strong evidence for photoionized gas in the line-of-sight to 13/18 objects. We present detailed modelling of this gas based upon the {\\tt ION} photoionization code. Other studies have been made of the 'warm absorber' phenomenon but this paper contains the first consideration of the importance of the covering-fraction of the ionized gas and a direct comparison between models of attenuation by ionized versus neutral material. We find the X-ray ionization parameter for the ionized material is strongly peaked at $U_X \\sim 0.1$. The column densities of ionized material are typically in the range $N_{H,z} \\sim 10^{21}$--$10^{23} {\\rm cm}^{-2}$, although highly ionized (and hence psuedo-transparent) column densities up to $10^{24} {\\rm cm}^{-2}$ cannot be excluded in some cases. We also investigate the importance of the emission-spectrum from the ionized gas, finding that it significantly improves the fit to many sources with an intensity consistent with material subtending a large solid angle at the central source. Allowing a fraction of the continuum to be observed without attenuation also improves the fit to many sources, and is definitely required in the case of NGC~4151. A deficit of counts is observed at $\\sim 1$~keV in the sources exhibiting the strongest absorption features. We suggest this is likely to be the signature of a second zone of (more highly) ionized gas, which might have been seen previously in the deep Fe $K$-shell edges observed in some {\\it Ginga} observations. We find evidence that the ionized material in NGC~3227 and MCG-6-30-15 contains embedded dust, whilst there is no such evidence in the other sources We discuss these results in the context of previous studies and briefly explore the implications in other wavebands. ", + "introduction": "\\label{sec:intro} Prior to the launch of {\\it ASCA}, the paradigm for Seyfert 1 galaxies, was that the 2--10~keV regime can be well represented by a powerlaw continuum of (photon) index $\\Gamma \\sim 1.9$ (e.g. Nandra \\& Pounds 1994). Spectra obtained using the Large Area Counter (LAC) onboard {\\it Ginga} also revealed emission line features superimposed on this underlying 'primary' continuum attributable to Fe $K$-shell fluorescence in the 6.4--7.1~keV band. A flattening of the spectra was observed above $\\sim$ 10 keV thought to be the 'Compton-reflection hump'. Evidence for a $K$-shell absorption edge in the 7.1--8.9~keV band due to highly-ionized Fe was also reported in a number of cases (Nandra \\& Pounds 1994). The combination of these features offers an explanation as to why simple powerlaw models to the data obtained by earlier instruments typically revealed flatter ($\\Gamma \\sim 1.7$) spectra (e.g. Turner \\& Pounds 1989). There are a few exceptions to this rule, most notably NGC~4151 which exhibits a continuum noticeably flatter than the average ($\\Gamma \\sim 1.5$) and absorption within a substantial column of gas ($\\sim 10^{22}\\ {\\rm cm^{-2}}$) along the line of sight covering most (but possibly not all) of the central continuum (e.g. Yaqoob, Warwick, Pounds 1993 and references therein). This picture has been generally supported by {\\it ASCA} observations, with the addition that the higher spectral resolution afforded by the onboard instruments have shown the Fe $K\\alpha$ line to be significantly broadened (Mushotzky et al. 1995; Tanaka et al 1995). However, the situation in the soft ($< 1$~keV) X-ray regime is somewhat less clear. Historically, a confusing variety of spectral and temporal properties have been seen and/or implied in Seyfert 1 galaxies in the soft X-ray regime. A number of sources seemed to show a steep and rapidly variable 'soft excess' which could be characterised by either a low-temperature thermal component (e.g. Arnaud et al. 1985) or a steep second powerlaw (e.g. Turner \\& Pounds 1989; Turner, George \\& Mushotzky 1993). Data obtained using the solid-state spectrometer (SSS) onboard the {\\it Einstein Observatory} showed evidence for a complex spectral form in the soft X-ray regime (Turner et al 1991), in some cases this was thought to be due to leakage of the primary continuum through a patchy absorber (e.g. Reichert et al. 1985). Evidence for absorption by ionized material along the line-of-sight to active galactic nuclei (AGN) was first obtained in observations of the QSO MR2251-178. Using data from the Monitor Proportional Counter (MPC) onboard the {\\it Einstein Observatory}, Halpern (1984) reported the column density in this source to have increased by a factor $\\gtrsim4$ between two observations separated by $\\sim1$ yr. The excess flux observed in a subsequent observation using the High Resolution Imager (HRI) on the {\\it Einstein Observatory} led Halpern to suggest an explanation in terms of material with a column density $\\sim 10^{22}\\ {\\rm cm^{-2}}$, but photoionized such as to be transparent below the $K$-edge of O{\\sc vii} at 740~eV. Such an interpretation was strongly supported by {\\it EXOSAT} observations of MR2251-178. Pan, Stewart \\& Pounds (1990) found an inverse correlation between the inferred column density and the source flux, and demonstrated that this was consistent with the behaviour expected from photoionized material along the line-of-sight. Similar behaviour was observed in NGC~4151 (Fiore, Perola \\& Ramano 1990; Yaqoob \\& Warwick 1991) and MCG-6-30-15 (Nandra, Pounds \\& Stewart 1990), establishing the effect for low and high luminosity sources. Supporting evidence for such ``warm absorbers'' was provided by {\\it Ginga} observations of emission-line AGN which revealed iron $K$-shell features in some sources, suggestive of an origin in ionized material with columns densities $\\gtrsim 10^{23}\\ {\\rm cm^{-2}}$. Nandra \\& Pounds (1994) find this to be a common occurrence, with 12 out of 27 sources in their sample of Seyfert 1 galaxies and Narrow Emission Line Galaxies (NELGs) showing such a component. However, the low spectral resolution of {\\it Ginga} meant that the results were likely to be sensitive to the assumptions concerning the adjacent Fe $K$-shell emission line. Confirmation of the existence of warm absorbers came from {\\it ROSAT} Position Sensitive Proportional Counter (PSPC) observations, which revealed absorption edges attributable to ionized oxygen in numerous sources (e.g. MCG-6-30-15, Nandra \\& Pounds 1992; 3C351, Fiore et al. 1993; NGC~3783, Turner et al. 1993a). Early {\\it ASCA} observations (e.g MCG-6-30-15, Fabian et al. 1994, Reynolds et al. 1995; NGC~3783, George, Turner \\& Netzer 1995) immediately confirmed the PSPC results and showed evidence for warm absorbers in many other sources (e.g. NGC~3227, Ptak et al 1994; NGC~4051, Mihara et al 1994; Guainazzi et al 1996; EXO~055620-3820.2, Turner, Netzer \\& George 1996; Mrk~290, Turner et al 1996; PG1114+445, George et al. 1997a). Furthermore Reynolds (1997) finds evidence for such a component in at least half of a sample of {\\it ASCA} observations of 24 type-1 AGN. The presence of highly-ionized, absorbing gas instrinsic to Seyfert 1 nuclei is further confirmed by the presence of 'associated absorbers' seen in the UV in many sources. Discussion of these observations and the attempts to relate the UV and X-ray absorbers made to-date is postponed to \\S\\ref{Sec:XUVabso-disc}. Comparison of {\\it ASCA} data with the predictions of detailed photoionization calculations have revealed spectroscopic evidence for a second ionized absorber in NGC~3516 (Kriss et al. 1996a). In addition the absorption edges due to ionized oxygen appear to vary in a way which is inconsistent with the behaviour of a single absorber in at least some sources. In MCG-6-30-15, the depth of the O{\\sc viii} edge changes on a timescale $\\sim 10^4$~s with no corresponding change in the O{\\sc vii} edge (Reynolds et al 1995; Otani et al. 1996), whilst the converse is true in NGC~4051 (Guainazzi et al. 1996). The presence of emission lines at energies $\\lesssim 1$~keV has been suggested in a number of previous observations (e.g. Turner et al 1991). A blend of emission lines due to O{\\sc vii} (0.568--0.574~keV) has been claimed in the {\\it ASCA} data from NGC~3783 (George, Turner \\& Netzer 1995) and potentially seen in other sources (e.g. IC~4329A, Cappi et al. 1996). However the reality of such features has since been disputed, because of the realization of specific inadequacies in the low energy calibration of the {\\it ASCA} instrumentation (see \\S\\ref{Sec:calib_uncert}). Here we present the third in a series of papers describing the X-ray properties of a sample of 18 Seyfert 1 galaxies, using data obtained by {\\it ASCA}. The sample is comprised of 23 {\\it ASCA} observations of Seyfert class 1.0--1.5 galaxies (as defined by Whittle 1992) performed prior to 1994 May 01. In Nandra et al. (1997a, hereafter Paper~I), we presented imaging and timing data for this sample. In that paper we demonstrate that all of the sources show evidence for variability on timescales of minutes-hours, with an amplitude anti-correlated with the source luminosity, confirming previous results. We also showed that for a number of sources the variability amplitude is greater in soft X-rays ($\\lesssim 2$~keV) than in hard X-rays (for 10 of the 18 sources on timescales of $\\sim1$~hr). In Nandra et al. (1997b, hereafter Paper~II), we presented the 3--10 keV spectra of the sample sources, and found 14 out of the 18 sources to contain an iron ${\\rm K}\\alpha$ line which is resolved in the {\\it ASCA} SIS, with mean width $\\sigma \\simeq 0.4\\pm 0.1$~keV for a Gaussian profile. However, we found many of the line profiles to be asymmetric, suggesting an origin from the innermost regions of the putative accretion disk. In this paper, we present the 0.6-10.0 keV spectra of the sample sources. The extension of the spectral analysis down to 0.6 keV requires consideration of the properties of the circumnuclear absorbing material as well as any steepening of the spectrum at low energies. We present detailed modelling of absorption by photoionized material and discuss likely emission contributions as well as combinations of several emission and absorption components. We present the range of properties applicable to our sample, and discuss the characteristics and statistics of the ionized absorbers. We also consider the the importance of allowing a fraction of the underlying continuum to be unattenuated by ionized gas, and the inclusion of the emission features expected from the ionized material. As will become apparent below, there is an obvious variation in signal-to-noise ratio between different datasets and hence there are variations in detectability of the ionized gas. As in the two previous papers our intention is to study the mean properties of the sample and we only include detailed notes on individual sources as an Appendix. We describe the observations and data reduction in \\S\\ref{sec:anal}, and the photoionization models in \\S\\ref{Sec:ion_model}. Our spectral fitting procedure is introduced in \\S\\ref{Sec:procedure}, and the results presented in \\S\\ref{Sec:basic_models} and \\S\\ref{Sec:additional_models}. In \\S\\ref{Sec:disc-multi} we discuss the objects for which there in more than one observation within our sample. Our results are discussed within the general context of the physical properties AGN in \\S\\ref{Sec:Discussion}. We sumarize our overall findings in \\S\\ref{Sec:open_issues} along with outlining a number of important issues that remain to be resolved. ", + "conclusions": "\\label{Sec:Discussion} As discussed in the introduction, it is becoming increasingly clear that the spectra of a number of individual AGN contain features indicative of absorption by ionized gas within the cylinder-of-sight. It seems highly likely that this gas is photoionized by the intense radiation field produced by the central object. The aim of the work presented here was to determine the frequency and characteristics of such gas based upon new, self-consistent modelling using the {\\tt ION} photoionization code. In the preceeding sections we have presented the results from 23 {\\it ASCA} observations of 18 objects. When considered together these objects certainly do not constitute a complete sample. Nevertheless, in at least some aspects, we do believe our results to be representative of the variety of properties exhibited by the type-1 AGN as a whole. In \\S\\ref{Sec:basic_models} we considered models assuming the underlying continuum in the 0.6--10~keV band of our sources is well-represented by a single powerlaw. A number of models were considered, starting with the case in which the spectrum emerging from the source passes through two screens of neutral material fully covering the source: one at the redshift of the source and a second at zero redshift (with column densities $N_{H,z}$ and $N_{H,0}$ respectively). Models of increasing complexity were then considered by allowing a fraction ($D_f$) of the underlying continuum to be observed {\\it without} attenuation by $N_{H,z}$, allowing the gas represented by $N_{H,z}$ to be photoionized (with an ionization parameter $U_X$), and finally by also including the emission spectrum expected from the ionized gas (assuming it subtends a solid angle $\\Omega$ at the central source). In \\S\\ref{Sec:additional_models} the results of further spectral analysis were presented, including models with additional spectral components, with the primary goal of testing the robustness of the derived characteristics of the ionized gas. \\subsection{General Considerations} \\label{Sec:disc-general} All models considered in this work are single-zone models, and given the assumed gas density $n_H = 10^{10}\\ {\\rm cm^{-3}}$, the gas reacts instantaneously to continuum variations. However, the models are intended to represent a wide range of conditions including cases with much lower values of $n_H$. In such cases, the level of ionization represents the response of the gas to the long-term, average flux of the illuminating source. We have not considered all such cases here since we do not consider the data yet warrant this level of discussion. In \\S\\ref{Sec:basic_models} we found that one or more of the models considered offer an explanation of 17/23 datasets within the formal criteria outlined in \\S\\ref{Sec:acceptability}. The criteria used to determine the acceptability of model as a description of the data is of course a somewhat subjective decision. The primary criterion used here that $P(\\chi^2 \\mid dof)\\leq 0.95$ (i.e. that the model is a true representation of the data with a probability of better than only 1 chance in 20) might be considered relatively lax by some readers. Furthermore, as can been seen from Tables~3--8, a 'perfect' parameterisation of the data ($P(\\chi^2 \\mid dof) \\simeq 0.5$ corresponding to a reduced--$\\chi^2$ value of unity) is only obtained in a small number of cases. Relatively high values of $P(\\chi^2 \\mid dof)$ imply that there are residuals in the data not modelled adequately. However, whilst some of these residuals might be carrying physical information associated with the source, others might be the result of inaccuracies in the calibration of the instruments (see also Fig.~13). The latter is currently under intense study by the instrument teams in anticipation of being able to assign unmodelled residuals to the former explanation with a higher degree of confidence. However our intention in this paper is to explore just how well the data from the sample can be described by our relatively simple models for the ionized gas. We believe this approach compliments that taken by some previous workers where the absorption features introduced by the ionized material are modelled as a series of absorption edges. Throughout the analysis we have not assumed any systematic errors, and have implicitly assigned the relatively poor values of $P(\\chi^2 \\mid dof)$ to be the result of problems with the calibration in the case of 4 additional datasets in \\S\\ref{Sec:basic_models} (Table~12). Furthermore, beyond the limited number considered in \\S\\ref{Sec:additional_models}, we have not explored more complex models in detail. For instance we have not considered models in which the ionized material has a significant ($\\gtrsim 10^4\\ {\\rm km\\ s^{-1}}$) velocity dispersion along the line-of-sight. Given the spectral resolution of the SIS detectors, dispersions smaller than this value are impossible to constrain with the signal-to-noise ratio of the data presented here (see also \\S\\ref{Sec:redshift}). We have not explored different abundance ratios, or multi-component models in those cases where we have indications for their existence (see \\S\\ref{Sec:disc-2ion}). Our intention in the present work is merely point the reader to the need of such analysis and defer it for future discussion. \\subsection{The form of the underlying continuum} \\label{Sec:disc-2plaw} It is well established that many Seyfert~1 galaxies exhibit so-called 'soft-excesses' and 'Compton-reflection' tails (e.g. Turner \\& Pounds 1989; Nandra \\& Pounds 1994). After accounting for these effects, and those of the ionized-absorber, {\\it Ginga} observations showed that the underlying X-ray continua can be generally well represented by a powerlaw of photon index $\\Gamma \\sim 1.9$--2.0 (e.g. Nandra \\& Pounds 1994). In \\S\\ref{Sec:basic_models} we obtained acceptable fits for most of the datasets by one or more of the models considered, with the distribution of the photon index peaked at $\\Gamma \\sim 2$ (Fig.~2). In \\S\\ref{Sec:complex_cont} we found significant improvements in the goodness--of--fit were obtained for many of the datasets by the inclusion of additional spectral components. However, in the majority of cases, we again found the bulk of 0.6--10~keV bandpass to be consistent by an underlying powerlaw with $\\Gamma \\sim 2$. It should be noted that in a number of cases the best-fitting models imply that very little of the underlying continuum is observed directly (e.g. see NGC~3783(1,2) in Fig~6; NGC~4151(2,4,5) in Fig~9). In passing we note that in such cases, $\\Gamma$ can still be well constrained by the data presented here as the shape of the 'recovery' in the absorption features in the 1--5~keV band is well defined in the models and well constrained by {\\it ASCA} data. For comparison with other work, we have calculated the expectation value and dispersion of the distribution of spectral indices, employing the method of Maccacaro et al. (1988). Considering those datasets for which we obtained acceptable fits, we find $<\\Gamma> = 1.94 \\pm 0.08$, with a significant, intrinsic dispersion of $0.18\\pm 0.07$. These values are in excellent agreement with those found by {\\it Ginga} and in Paper~II, when the effects of Compton reflection are taken into consideration. In a number of cases, complex continua do seem to be present in the {\\it ASCA} bandpass. Mrk~335 does indeed appear to possess real curvature in the underlying continuum (Fig~15). Evidence was also found for curvature in the observed continuum in NGC~4051, IC~4392A \\& NGC~5548. However in at least NGC~5548 we consider the evidence for curvature in the {\\it underlying} continuum to be questionable and/or alternative explanations are available (see \\S\\ref{Sec:2plaw}). The observed spectrum of NGC~4151 has long been considered somewhat idiosyncratic amongst Seyfert 1 galaxies. However we do find acceptable solutions and consider the underlying continuum to indeed be a powerlaw, albeit flatter ($\\Gamma \\sim 1.5$) than that observed for the majority of sources (see \\S\\ref{Sec:disc-multi} and \\S\\ref{Sec:disc-Df}). \\subsection{The frequency of neutral absorbers} \\label{Sec:disc-freq-neut} Unlike Seyfert 2 galaxies, Seyfert 1 galaxies are generally not believed to contain significant column densities of neutral gas detectable in the 0.6--10~keV band (ie. within the range $10^{20} \\lesssim N_{H,z} \\lesssim 10^{25}\\ {\\rm cm^{-2}}$). In the case of models {\\it A(i)} and {\\it A(ii)}, the absorption at the redshift of the source (with column density $N_{H,z}$) is assumed to be due to neutral material. As described in \\S\\ref{Sec:zwabspo}, significant absorption is not detected in any of the datasets for which model {\\it A(i)} offers an acceptable solution (with $N_{H,z} \\lesssim$ few $\\times10^{20}\\ {\\rm cm^{-2}}$; see Fig~3a). Model {\\it A(ii)}, in which a fraction $D_f$ of the underlying continuum is able to escape without passing through $N_{H,z}$, does offer an improvement in the goodness--of--fit for many datasets, and a range of derived column densities for those in which an acceptable solution was obtained (Fig~3b). However, as discussed in \\S\\ref{Sec:zpcfpo}, these solutions are most likely due to a residual problem in the instrument calibration (e.g. 3C~120), lie in a region of the $N_{H,z}$,$D_f$ indicative of (relatively-subtle) spectral curvature (e.g. NGC~7469), and/or superior solutions are obtained assuming subsequent models (e.g. Mrk~335). We note that it has long been suggested that NGC~4151 might contain a significant column density attenuating $\\sim 95$\\% of the continuum. However as shown in \\S\\ref{Sec:basic_models}, we find acceptable solutions only for models in which the gas is ionized (see also \\S\\ref{Sec:disc-multi} and \\S\\ref{Sec:disc-Df}). We have not explicitly tested for {\\it additional} intrinsic column densities of neutral material in any of our models which contain ionized gas (i.e. in models {\\it B(i)}--{\\it C(ii)}). Such a situation might arise for example in the case where the primary continuum passes through a region of ionized gas, and then a further screen of neutral gas during its passage through the host galaxy. Nevertheless, all our models do contain a neutral column at zero redshift, $N_{H,0}$, to account for the Galactic absorption (and hence expected to be $\\simeq N_{H,0}^{gal}$). At the redshift of most of the sources in our sample, any absorption by neutral material intrinsic to the source would be indistinguishable from absorption at zero redshift. Thus any such absorption would result in a derived value of $N_{H,0} >> N_{H,0}^{gal}$. Generally we find no such evidence in our sample, with the only reliable detection in the case of IC~4329A (with a neutral column density $\\sim 4\\times10^{21}\\ {\\rm cm^{-2}}$). The host galaxy is this AGN is observed edge-on so this result is not surprizing (see Appendix). \\subsection{The frequency of ionized absorbers} \\label{Sec:disc-freq} In \\S\\ref{Sec:basic_models} we found 16/23 datasets (11/18 objects) are improved at $>99$\\% confidence over models {\\it A(i)} \\& {\\it A(ii)} if the absorbing material is assumed to be photoionized. Of these, 12/16 datasets (9/11 objects) can be adequately described by one or more of models including ionized gas presented in \\S\\ref{Sec:basic_models} (i.e. models {\\it B(i)}--{\\it C(ii)}), under the assumption of a single power-law continuum. Of the remaining 4/16 datasets, 3 datasets (only one of which being from an additional object: NGC~3516) show strong evidence for ionized gas (with the statistical poverty of the fit due to further spectral complexity). Thus we find clear evidence for absorption by ionized gas in 15/23 datasets (10/18 objects). However the presence of ionized gas is also strongly suspected in 3 additional datasets (3 additional objects: NGC~4051, IC~4329A, NGC~7469) from the analysis presented in \\S\\ref{Sec:additional_models}. \\begin{itemize} \\item\t{\\it Thus we conclude there is evidence for ionized gas in 18/23 datasets (13/18 objects) in our sample}. \\end{itemize} Reynolds (1997) has also presented results from a sample of 24 type-1 AGN using {\\it ASCA} data, using the same datasets as presented here for 15 sources (Table~2). Based on a search for O{\\sc vii} and O{\\sc viii} edges, Reynolds found strong evidence for ionized gas in 12/24 objects. These 12 include 9 sources also in our sample (Table~12), along with Mrk~290, 3C~382 and MR~2251-178. Reynolds also fitted these data with a single-zone photoionization model (based on the photoionization code \\verb+CLOUDY+ which is similar in approach and scope to \\verb+ION+ used here), assuming a photoionizing continuum consisting of a single powerlaw over the entire 0.0136--13.6~keV band. Converting their quoted ionization parameters ($\\xi$) to $U_X$ using Fig.1b, gives $1.5 \\lesssim \\Gamma \\lesssim 2.0$, $10^{21} \\lesssim N_{H,z} \\lesssim 2\\times10^{23}\\ {\\rm cm^{-2}}$ and $0.02 \\lesssim U_X \\lesssim 0.09$, broadly in agreement with the distributions found here (Figs~2, 3 and 7 respectively). A detailed source-by-source comparison between our results and those of Reynolds (1997) and other workers is provided in the Appendix. \\subsection{Column density and ionization parameter of the ionized gas} \\label{Sec:disc-hagai} We find no preferred value of $N_{H,z}$ under any of the models tested, with clear differences from object--to--object (e.g. Fig.~3). However, in each of the models including photoionized gas, we find the distribution in the ionization parameter strongly peaked at $U_X \\sim 0.1$ (Fig.~7). Considering those datasets for which we obtained acceptable fits (and again employing the method of Maccacaro et al. 1988), we find $<\\log U_X> = -0.92 \\pm0.21$. While the values cluster around this value, there is a highly significant intrinsic dispersion in $\\log U_X$ of $0.21^{+0.32}_{-0.09}$. Nevertheless this relatively narrow range in $U_X$ is our most interesting and unexpected result. As such it deserves some further discussion as the whether it is purely the result of selection effects, or has physical significance. We considering first values of $U_X >> 0.1$. With the signal--to--noise ratio typically afforded by {\\it ASCA} data, the presence of ionized gas is most readily inferred by deep ($\\tau \\gtrsim 0.2$) absorption edges which it introduces into the observed X-ray spectrum. For a given value of $N_{H,z}$, there is a maximum value of $U_X$ above which the gas is unstable and goes to the high electron temperature ($T_e \\sim 10^6$~K) branch of the two-phase curve (e.g. see Fig.~2 in Netzer 1996). The value of $U_X$ at which this thermal instability occurs is a function of the gas density and composition, along with the form of the ionizing continuum. Under the assumptions made here, the instability occurs for $U_X > 10^{-22} N_{H,z}$. Such gas is completely transparent, and the observed spectrum will be identical to the underlying (powerlaw) continuum, consistent with the data from a minority of objects in our sample. Therefore the presence of such material along the cylinder--of--sight is impossible to prove (or disprove) using the current {\\it ASCA} data alone, offering a potential explanation of the apparent lack of objects with $U_X >> 0.1$ in Fig.~7. For $U_X << 0.1$, different ions contribute different continuum opacity, depending on the gas composition and level of ionization. For $10^{21} < N_{H,z} < 10^{24}\\ {\\rm cm^{-2}}$ and $0.03 10^{22}\\ {\\rm cm^{-2}}$. There is no reason {\\it a priori} why there should not be objects with ionized material in this region. However, there is a potential selection-effect here whereby if such material lays along the cylinder--of--sight to the broad emission line region (BELR) and contains embedded dust, the broad emission lines will be absorbed by this material. This may lead to such objects being classified as something other than Seyfert 1 galaxies, and hence being excluded from the sample of sources presented here. Alternatively, should this explanation not be the case, then there must be some yet unknown mechanism that determines these unique conditions by either fixing the pressure or perhaps the dynamical state of the absorbing material. Future studies of a larger sample of objects are required to investigate these possibilities. We have compared the parameters of the ionized gas with the derived continuum luminosity, $L_X$, in the 0.1--10~keV band (after correcting for absorption). We find no clear relation between $N_{H,z}$ and $L_X$, but some indication that $U_X \\propto L_X$ (Fig.~17). If true, from the definition of $U_X$ (Eqn.~\\ref{eqn:U_X}), this implies the ionized gas has similar densities, $n_H$, and is at similar radii, $r$ in all sources. However, the majority of the sources considered here for which $U_X$ is well constrained lie within a restricted range of $L_X$ ($10^{43}$--$10^{44}\\ {\\rm erg\\ s^{-1}}$). A recent observation of the quasar PG~1114+445 ($L_X \\sim 6\\times 10^{44}\\ {\\rm erg\\ s^{-1}}$) also exhibits strong absorption features due to ionized gas with $N_{H,z} \\simeq 2\\times10^{22}\\ {\\rm cm^{-2}}$ and $U_X \\simeq 0.1$ (George et al. 1997b). Furthermore, Nandra et al (1997e) have recently reported the results from {\\it ASCA} observations of the RQQ MR~2251-178 ($z = 0.068$, $L_X \\simeq 2\\times 10^{45}\\ {\\rm erg\\ s^{-1}}$), finding $N_{H,z} \\sim 2\\times10^{21}\\ {\\rm cm^{-2}}$ and $U_X \\sim 0.07$. Inclusion of these objects on Fig.~17 argues against any obvious relation between $U_X$ and $L_X$. We found no evidence for any relationship between any other pairs of parameters derived from the X-ray results (e.g. $L_X$ vs $\\Gamma$; $\\Gamma$ vs $U_X$; $\\Gamma$ vs $N_{H,z}$). In most cases this is again primarily due to the limited range of $L_X$ and/or $\\Gamma$ exhibited by most the objects in the sample. Thus, further progress requires high signal-to-noise observations of similar sources covering the high and low $L_X$ regimes. \\subsection{Constraints on ionized emitters} \\label{Sec:disc-emis} The dominant effects of the ionized gas on the observed spectrum from these sources is absorption by the material along the line-of-sight. However, ionized material out of the line-of-sight will give rise to emission lines and recombination continua, along with the (absorbed) continuum Compton-scattered back into the line-of-sight. The inclusion of such an emission component leads to an improvement at $>$99\\% confidence in 11 datasets (9 objects) in the full covering case (model {\\it C(i)}), and 8 datasets (7 objects) in the partial covering case (model {\\it C(ii)}). However there is no case for which this component is {\\it required} (i.e. a satisfactory fit only obtained for model {\\it C}). It should also be noted that the inclusion of the emission component generally does not have a significant effect on the values derived for $U_X$ and $N_{H,z}$. For the value of $U_X$ and $N_{H,z}$ derived for the majority of sources, the emitted spectrum contains significant line emission in the 0.6--2.0~keV band, but constitutes only a relatively small fraction ($\\lesssim 10$\\%) of the total flux observed in this band, even when the ionized gas subtends a large solid angle at the central source. The best-fitting values of $\\Omega$ are ill-constrained, but in most cases consistent with $\\Omega = 4\\pi$, and there is no case where $\\Omega > 4\\pi$ at $>95$\\% confidence. Formally we find $< \\log \\Omega/4\\pi >$ $= 0.23^{+0.24}_{-0.23}$ for the datasets for which we obtained acceptable fits, In cases where the ionizing continuum source is constant, then the normalization of the emission spectrum is a direct measure of the geometry of the emitting gas. However in cases where the ionizing continuum varies, then the normalization of the emission spectrum will follow the luminosity history of the continuum source, with a lag due to light travel-time effects. In the latter case we are therefore offered the opportunity to determine the location of the ionized gas via the delay in the intensity of the emission features following variations in the ionizing continuum. Unfortunately however, the signal-to-noise ratio is too low for the few objects presented here for which there are multiple observations (\\S\\ref{Sec:disc-multi}) to place even crude constraints on whether the intensity of the emission changed in an appropriate manner. \\subsection{Evidence for an unattenuated component} \\label{Sec:disc-Df} In \\S\\ref{Sec:basic_models} \\& \\S\\ref{Sec:additional_models}, we find that allowing a fraction $D_f$ of the powerlaw continuum to be observed without suffering attenuation by the column density $N_{H,z}$ significantly improves the fit (compared to the corresponding model with $D_f=0$) for a large number of dataset/model combinations. Indeed, inspection of $F(\\frac{Aii}{Ai})$ (Table~4) reveals all but 3 datasets are improved in the case of model {\\it A(ii)}, and inspection of $F(\\frac{Bii}{Bi})$ and $F(\\frac{Cii}{Ci})$ (Tables~6 and 8) reveals that 8 datasets are improved in the case of models {\\it B(ii)} and {\\it C(ii)}. However, in many of these cases the best-fitting solutions have $D_f \\sim 1$ (Fig.~5) and as described in \\S\\ref{Sec:complex_cont} are an artifact of spectral curvature. Indeed, as shown in Table~12, only in the case of 2 datasets (1 object: NGC~4151) are solutions which satisfy our criteria for acceptability obtained {\\bf only} if $D_f>0$. Thus we conclude most of the sample show solutions close to the full-covering case and allowing $D_f$ as a free parameter does not significantly change the properties of the sample. As noted earlier, this model is appropriate to a partial covering of the cylinder--of--sight and to a geometry in which some fraction of the underlying continuum, initially emitted in other directions, is scattered back into the line-of-sight. As discussed in \\S\\ref{Sec:disc-multi}, a possible explanation for the spectral variability observed in NGC~4151 consists of clumps of photoionized gas (of differing column density) moving through the cylinder--of--sight on timescales $\\lesssim 1$~day. Under such an hypothesis, the fact that $D_f$ is similar for all the observations (at $\\sim$5\\%) suggests scattering as the most likely explanation. Circumstantial support is provided by the observation of a similar component in many type-2 AGN, with implications for unification schemes (e.g. Turner et al 1997a,b). Indeed the optical spectrum of NGC~4151 has been observed to exhibit both type-1 and type-2 characteristics at various times. It should be noted that the parameterization used here approximates the scattered continuum as a simple powerlaw as would be appropriate in the case of pure electron scattering in a very highly-ionized medium. In a more realistic treatment, the ionization state of the scattering gas would be taken into account, which under a wide range of parameter space would lead to emission and absorption features further complicating the spectrum $\\lesssim 2$~keV. It should be noted that the fact that such a component is unambiguously observed only in NGC~4151 amongst our sample is simply due to the suppression of the transmitted continuum by the high column density of gas within the cylinder--of--sight of this source (with the gas having $U_X$ such that there's still significant opacity) \\subsection{Implications of the '1~keV deficit'} \\label{Sec:disc-2ion} A number of sources have evidence for a deficit of counts at $\\sim$1~keV compared to the predictions of our photoionization models. As noted above, such a deficit occurs in those sources exhibiting the strongest absorption features. In \\S\\ref{Sec:2ion} we showed that this could be modeled with an additional edge (consistent with O{\\sc viii}) in the rest frame of the source. There are several possible explanations of such a feature. First, it could be indicative that the cylinder-of-sight contains at least two absorption systems, each of different density, $n_H$, and column density, $N_{H,z}$. Such an hypothesis is supported in the case of MCG-6-30-15 by the short-timescale variability observed in the depth of the O{\\sc viii} edge whilst the depth of the O{\\sc vii} edge remains constant (\\S\\ref{Sec:disc-multi}). Circumstantial evidence for such a possibility also comes from the {\\it Ginga} observations that a large fraction of Seyfert 1 galaxies (including some of those considered here) contain an Fe $K$-shell absorption edge, far deeper than that predicted by the models considered here (Nandra \\& Pounds 1994; \\S\\ref{Sec:Fe-absorption}). Unfortunately the geometry of the absorption systems cannot be constrained by the current data. One possible geometry is that both absorbing systems consist of complete screens, but at different radii from the central source. A full and correct treatment of such a scenario requires the ionization-equilibrium of the second screen of gas to be calculated in the same way as the first screen, but with the photoionizing continuum appropriate for that which has already passed through the first screen. Such a treatment is beyond the scope of the current paper. An alternative geometry is that the absorbing screen is clumpy on scale-sizes smaller than the emission region, resulting in the observed spectrum being the sum of the spectra transmitted by regions of different density and $N_{H,z}$. Yet another possibility is that there is indeed only a single absorbing cloud along the line-of-sight at any instant, but that different clouds (of different densities \\& $N_{H,z}$) move across the line-of-sight on timescales much shorter than the typical {\\it ASCA} observation. In this case our analysis reveals time-averaged parameters of the absorber. Finally, we note that some of our detailed assumptions regarding atomic processes may be inapplicable to AGN, such as other processes important in determining the ionization equilibrium for O{\\sc vii}/O{\\sc viii}. Given the uncertainty in the form of the 'primary' continuum, the spectral resolution and typical signal-to-noise ratios typically achieved, it appears to be extremely difficult to distinguish between these possibilities using {\\it ASCA} data. \\subsection{The location of the ionized gas} \\label{Sec:disc-location} The location and geometry of this ionized gas is currently unclear, although there are prospects for future progress. Of course if one knows the density, $n_H$, of the ionized gas, $U_X$ and the intensity of the photoionizing source, one can use eqn.~\\ref{eqn:U_X} to derive its radius as \\begin{equation} r_{ld} \\simeq 3\\times10^{5} (\\frac{L_{X44}}{U_X n_H})^{0.5}\\ {\\rm light-days} \\end{equation} Assuming $n_H = 10^{10}\\ {\\rm cm^{-3}}$ (as used for the model calculations), $U_X \\simeq 0.1$ and $L_{X44}\\sim 0.1$--1 (as appropriate for the majority of the objects presented here), one obtains $r_{ld} \\sim$few light-days. Such a location is comparable to that of the BELR in these objects. However as noted in \\S\\ref{Sec:ion_model}, our models are insensitive to densities in the range $10^{4} \\lesssim n_H \\lesssim 10^{10}\\ {\\rm cm^{-3}}$, thus values of $r_{ld} \\lesssim$few light-years cannot be excluded on this basis. As described in \\S\\ref{Sec:disc-multi}, the apparent variations in the column density of the ionized gas seen in the X-ray band can be used to place upper limits on the radius of the gas if one assumes these variations are due to inhomogeneities in the matter passing through the cylinder--of--sight due to {\\it only} Keplerian motion. For the few cases observed to-date, such an assumption also implies a location close to the BELR and hence is consistent with all/part of the same gas being responsible for the absorption features seen in the optical/UV. However such estimates are currently extremely crude, and it should be stressed that the bulk motion of the ionized gas is highly uncertain and could easily be dominated by other kinematics components. Constraints on the location could be placed by applying (under certain assumptions) photoionization models to both the UV/optical and X-ray absorption features. As will be described in \\S\\ref{Sec:XUVabso-disc}, the results from current studies are mixed and future progress requires simultaneous, high-resolution observations in all three wavebands. Reverberation mapping using the emission features produced within the ionized gas is another means whereby the location of the material could be determined. However as described in \\S\\ref{Sec:disc-emis}, current results are inconclusive. However we believe that in all likelihood there is ionized gas throughout the nuclear region in AGN. Multiple components, separated in at least velocity-space are seen explicitly in the optical/UV absorption features of some objects. At least two components are implied spectroscopically and/or from temporal variations in the X-ray absorption features of some objects. In addition, highly ionized gas (occupying a different region of $U_X$, $N_{H,z}$ parameter-space than that responsible for the absorption on Seyfert 1 galaxies) is commonly invoked to explain the 'scattered' radiation in Seyfert 2 galaxies (e.g. see Turner et al 1997b,c). \\subsubsection{The possible effect on UV--IR observations} \\label{Sec:disc-dust} Should the volume of ionized gas contain embedded dust, and is located outside the appropriate emission regions, then one might expect a correlation between the column density $N_{H,z}$ derived from X-ray observations and the various reddening indicators at longer wavebands. Such studies therefore have the potential for further constraining the location of absorbing material and/or geometry of the nuclear regions. Brandt, Fabian \\& Pounds (1996) have recently suggested that the infrared-bright quasar IRAS~13349+2438, which exhibits a large degree of reddening in the optical, might contain photoionized gas with internal dust. Reynolds (1997) noted that MCG-6-30-15, whose X-ray spectrum is clearly attenuated by ionized gas, also exhibits a large degree of reddening at longer wavelengths. The study of the creation, survival and implications of dust within the nuclear regions of AGN has received much attention for several decades (e.g. see Laor \\& Draine 1993 and references therein). Unfortunately the use of the various reddening indicators in order to determine the column density of dust along the line--of--sight (and which can then be compared to that observed in the X-ray regime) rely upon making one or more assumptions. The primary assumption is of course that one knows the intrinsic value of the quantity from which the observed reddening is determined. If one knew the form of underlying continuum, multiwaveband continuum measurements could be applied in the optical--UV . However such a technique often suffers from a lack simultaneous observations (as well as observational 'contamination' of other spectral components). Broad emission line ratios, in particular broad hydrogen lines, are poor reddening indicators since the lines are significantly affected by optical depth effects and there is no satisfactory theory to predict their intrinsic ratios (Netzer 1990). Clearly all techniques also require assumptions concerning the composition (chemical and grain-size distribution) of the dust itself, and well as the dust/gas mass ratio. With these points in mind, we consider only a single reddening indicator here (resonance-absorption line studies in the UV/optical are also considered in \\S\\ref{Sec:XUVabso-disc}, but there do not provide a direct diagnostic of the dust-phase material). Fig.~18 shows $(f_{125}/f_{220})_{obs}$ (defined as the mean ratio of the observed flux at 125~nm to that at 220~nm, determined as described in \\S\\ref{Sec:sample} and listed in Table~1) is shown against $N_{H,z}$ for the datasets for which model {\\it C(ii)} was considered acceptable (\\S\\ref{Sec:ion_pc_emis}). It can be seen that the majority of sources are consistent with $2 \\lesssim (f_{125}/f_{220})_{obs} \\lesssim 5$, but that 3 datasets (NGC~3227, MCG-6-30-15(1,2)) have values a factor $\\sim 10$ lower. Thus it is tempting to make the assumption that the intrinsic flux ratio\\footnote{It should be stressed that line emission may contribute to $f_{125}$ and (to a lesser extent) $f_{220}$, and thus $(f_{125}/f_{220})_{int}$ is not necessarily a good indication of the underlying UV continuum. However, the use of the color here is only based on noting $(f_{125}/f_{220})_{obs}$ appears to be crudely constant (within a factor $\\sim 2$) for the majority of the objects considered.} is in the range $2 \\lesssim (f_{125}/f_{220})_{int} \\lesssim 5$, with the $(f_{125}/f_{220})_{obs}$ lower in the case of NGC~3227 \\& MCG-6-30-15 due to reddening. Also shown in Fig.~18 is the predicted $(f_{125}/f_{220})_{obs}$ assuming $(f_{125}/f_{220})_{int} = 3.1$ for various values of the {\\it difference} between the optical depths at 125~nm and 220~nm (where $\\Delta \\tau = \\tau_{125} - \\tau_{220}$). It can be seen that NGC~3227 and MCG-6-30-15(1,2) are consistent with $(f_{125}/f_{220})_{int}$ similar to the other objects if their fluxes at 125~nm and 220~nm are absorbed by a column density similar to that observed in the X-ray band {\\it and} $\\Delta \\tau = 0.5 N_{H,z}/10^{21}$ (dashed line). Interestingly, the 'standard Galactic' extinction curve (for a 1:1.12 silicate:graphite mixture of dust grains with sizes $5 \\leq a \\leq 250$~nm, number density $\\propto a^{-3.5}$ and total dust/gas mass ratio of $10^{-2}$ --- see Laor \\& Draine 1993 and references therein) gives $\\Delta \\tau$ close to this value ($\\Delta \\tau = 0.42 N_{H,z}/10^{21}$). Thus in the case of NGC~3227 and MCG-6-30-15(1,2), and assuming a Galactic dust/gas mass ratio, this UV color is consistent with reddening by a column density of dusty material similar to that inferred by the X-ray observations. There are several implications of these results. First, we consider the implications on the location of the dust and gas in the case of NGC~3227 and MCG-6-30-15. The consistency between the $N_{H,z}$ derived from the X-ray absorption studies and the UV-reddening in the case of these two sources, assuming a 'reasonable' dust/gas mass ratio, implies that the dust- and gas-phase material might occupy the same volume. If true, this allows us to place constraints on the location of this material based upon the survival of the dust. For a typical AGN continuum of bolometric luminosity $L_{bol} = f_{bolX} L_{X44}$ (as defined in \\S\\ref{Sec:bolX}), and assuming maximum equilibrium temperature of $\\sim 1750$~K, Laor \\& Draine find sublimation radii of $r_{sub} \\simeq 20 (f_{bolX} L_{X44})^{1/2}$~light-days and $\\simeq 150 (f_{bolX} L_{X44})^{1/2}$~light-days for graphite grains with $a = 10^{-3}$~cm and $5\\times10^{-7}$~cm respectively. (Silcate grains, with a slightly lower maximum equilibrium temperature ($\\sim 1400$~K), have slightly larger values of $r_{sub}$.) Assuming $f_{bolX} \\simeq 4$, $r_{sub} \\sim 10$~light-days and $\\sim 5$~light-days for large graphite grains in MCG-6-30-15 ($L_{X44} \\simeq 0.3$) and NGC~3227 ($L_{X44} \\simeq 2\\times10^{-2}$) respectively. The corresponding value in the case of NGC~5548 ($L_{X44} \\simeq 1$) is $r_{sub} \\sim 40$~light-days for large graphite grains. The latest results from the intensive reverberation mapping of the broad emission line region (BELR) in NGC~5548 have suggested lags of $\\lesssim 2 \\sim$14~days, with the lowest lags seen in the highest ionization lines (e.g. Korista et al 1995 and references therein). As first pointed out by Netzer \\& Laor (1993), since both the radius of the BELR and $r_{sub}$ scale as $L_{bol}^{1/2}$ we conclude dust can only survive (i.e. in {\\it equilibrium}) at radii just beyond the BELR in these objects. If the dust covers a substantial fraction of the BELR, then it will attenuate the broad emission lines in these objects. For instance, assuming the same 'standard Galactic' extinction curve used above, $\\tau({\\rm C{\\sc iv}}) = 9 N_{H,z}/10^{22}$ at the energy of C{\\sc iv} emission. Thus from the values of $N_{H,z}$ found in MCG-6-30-15 and NGC~3227, a broad C{\\sc vi} emission line is predicted to be totally absent in both sources, consistent with observations (see e.g. Courvoisier \\& Paltani 1992). The extinction in the optical is lower (e.g. $\\tau({\\rm H}_{\\beta}) = 3 N_{H,z}/10^{22}$, $\\tau({\\rm H}_{\\alpha}) = 2 N_{H,z}/10^{22}$) but sufficient to attenuate $\\sim$80\\% and $\\sim$95\\% of the broad ${\\rm H}_{\\beta}$ line, and $\\sim65$\\% and $\\sim$85\\% of the broad ${\\rm H}_{\\alpha}$ line in NGC~3227 and MCG-6-30-15 respectively (again assuming the dusty absorber completely covers the entire BELR). However it is clear from the value of $r_{sub}$ above and the discussion in \\S\\ref{Sec:multi-mcg6} that dusty clouds cannot survive at a sufficiently small radius that Keplerian motion alone can offer an explanation of the short-timescale variability ($\\sim 10^4$~s) in the O{\\sc viii}~edge observed in MCG-6-30-15. Thus either the dusty absorber is extremely non-uniform and/or its kinematics are extremely non-Keplerian, or (far more likely) the gas giving rise to the bulk of the O{\\sc viii}~edge is at much smaller radii than that with embedded dust. Interestingly NGC~3227 and MCG-6-30-15 have the lowest axial ratios ($a/b$, Table~1) in our sample, with the exception of the well-known, edge-on galaxy IC~4329A. We now consider the majority of the objects on our sample, for which $(f_{125}/f_{220})_{obs}$ provides no evidence for reddening. There are two obvious explanations. First, the column density of ionized gas within the cylinder--of--sight may simply not contain any embedded dust in these objects. This might arise from differences in the origin of the material, or as a result of the material in these objects being within $r_{sub}$. Alternatively the dust may be present yet unobservable. This will true if the dust is of a different composition in which $\\Delta \\tau \\sim 0$ (such as is the case if the grains are large), or if the dust clouds cover only a fraction of the UV emission region. \\subsubsection{Relationship to 'associated absorbers' in the UV/optical} \\label{Sec:XUVabso-disc} The presence of highly-ionized gas, instrinsic to the Seyfert 1 nuclei is also obtained from observations carried out in the UV. {\\it IUE} observations revealed only a small fraction ($\\sim$3\\%) of sources to exhibit narrow absorption lines close to the systemic velocity (e.g. Ulrich 1988). However, recent results taking advantage of the higher spectral resolution and signal-to-noise ratio available with the FOS and GHRS instruments onboard {\\it HST} have revealed that up to 50\\% of Seyfert 1 galaxies may in fact contain such absorption features (e.g. Crenshaw 1997). Such features have also been reported in {\\it HUT} observations of NGC~4151 (Kriss et al. 1992) and NGC~3516 (Kriss et al. 1996b). Furthermore variability in the UV absorption features have been observed on timescales of weeks--years in NGC~4151 (Bromage et al. 1985), NGC~3783 (Maran et al 1996), and NGC~3516 (Koratkar et al. 1996), and multiple radial velocity components have been reported in Mrk~509 (Crenshaw, Boggess \\& Wu 1995) and NGC~3516 (Crenshaw, Maran \\& Mushotzky 1997). As noted by Crenshaw (1997), there is a reassuring trend whereby Seyfert 1 galaxies which have evidence for strong absorption features due to ionized gas in the X-ray band also tend to have strong absorption lines in the UV/optical. This raises the question as to whether the associated absorbers seen the UV/optical are caused by (part or all of) the same ionized material which gives rise to the absorption features observed in the X-ray band. The UV absorption components are often blueshifted with respect to the emission lines, but with velocity shifts (typically $\\sim$few$\\times 10^{2}\\ {\\rm km\\ s^{-1}}$) too small to be detectable with current X-ray observations. Thus direct tests for kinematic consistency are not yet possible. However various attempts have been made to relate the UV/optical and X-ray absorption features using photoionization calculations (e.g. Mathur 1994; Mathur et al. 1994; Mathur, Wilkes \\& Elvis 1995; Shields \\& Hamann 1996; Kriss et al 1996a,b). These have met with varying degrees of success, which is hardly surprizing given the uncertainties in the location and kinematics of the gas and the spectral form of the photoionizing continuum. Nevertheless there have been a number of impressive successes linking the UV/optical and X-ray absorbers (e.g. see Mathur 1997). The observed UV/optical absorption features are often superimposed on the wings of broad emission lines. This clearly places the gas responsible for these 'associated absorbers' at radii greater than that where the bulk (of the blue-wing) of that emission line originates ($\\lesssim 2 \\sim$14~light-days in the case of NGC~5548). However ionized gas could exist at a wide range of radii, with only part of the gas responsible for the X-ray absorption within the cylinder--of--sight to the UV/optical emission line region(s). Indeed the presence of multiple components to the UV/optical absorption features (and the indication of additional absorption zones in the X-ray --- e.g. \\S\\ref{Sec:disc-2ion}) clearly indicates a single screen of gas to be an over-simplification. Thus the column densities implied by the UV/optical absorption features may be small (and dominated by different ionization states) compared to those implied from the X-ray features. Furthermore, the emission and absorption features predicted in the UV/optical from the ionized gas are somewhat dependent on the form of the photoionizing continuum in the UV--X-ray band. This is poorly determined for all objects, leading to further ambiguities. In addition, in at least some objects the absorption features in both the UV and X-ray are known to vary with time. To date, most comparisons between the absorption features have been performed using non-simultaneous data, leading to obvious risk of deception. Finally, there is the usual irritation of uncertainties in the abundances within the absorbing material. The best hope for progress is provided by intense, simultaneous monitoring of the UV/optical and X-ray absorption features." + }, + "9708/hep-ph9708215_arXiv.txt": { + "abstract": "We present a full characterization of the phase transition in U(1) scalar field theory and of the associated vortex string thermodynamics in 3D. We show that phase transitions in the string densities exist and measure their critical exponents, both for the long string and the short loops. Evidence for a natural separation between these two string populations is presented. In particular our results strongly indicate that an infinite string population will only exist above the critical temperature. Canonical initial conditions for cosmic string evolution are show to correspond to the infinite temperature limit of the theory. ", + "introduction": " ", + "conclusions": "" + }, + "9708/hep-ph9708308_arXiv.txt": { + "abstract": "\\thispagestyle{empty} \\begin{minipage}{5in} \\thispagestyle{empty} \\baselineskip 16pt Assuming three flavour neutrino mixing takes place in vacuum, we investigate the possibility that the solar $\\nu_e$ take part in MSW transitions in the Sun due to $\\Delta m^2_{31} \\sim (10^{-7} - 10^{-4})~eV^2$, followed by long wave length vacuum oscillations on the way to the Earth, triggered by $\\Delta m^2_{21}$ (or $\\Delta m^2_{32}$) $\\sim (10^{-12} - 10^{-10})~eV^2$. The solar $\\nu_e$ survival probability is shown to be described in this case by a simple analytic expression. New ranges of neutrino parameters which allow to fit the solar neutrino data have been found. The best fit characterized by the minimum $\\chi^2$ is extremely good . This hybrid (MSW+vacuum oscillations) solution of the solar neutrino problem leads to peculiar distortions of energy spectrum of the boron neutrinos which can be observed by the SuperKamiokande and SNO experiments. Other flavour scheme (e.g. 2 active $\\nu$s + 1 sterile $\\nu$) can provide MSW+vacuum solution also. \\end{minipage} ", + "introduction": " ", + "conclusions": "" + }, + "9708/astro-ph9708251_arXiv.txt": { + "abstract": "The low frequency magnetic field spectrum in the primordial plasma is of particular interest as a possible origin of magnetic fields in the universe (e.g., Tajima et al. 1992 and Cable and Tajima 1992). We derive the magnetic field spectrum in the primordial plasma, in particular, at the epoch of primordial nucleosynthesis. The pioneering study of Cable and Tajima (1992) of the electromagnetic fluctuations, based on the {\\it Fluctuation-Dissipation Theorem}, is extended. Our model describes both the thermal and collisional effects in a plasma. It is based on a kinetic description with the BGK collision term. It is shown that the zero-frequency peak found by Cable and Tajima (1992) decreases. At high frequencies, the blackbody spectrum is obtained naturally without the necessity of the link procedure used by them. At low frequencies ($\\omega \\leq 4\\omega_{pe}$, where $\\omega_{pe}$ is the electron plasma frequency) it is shown that the magnetic field spectrum has more energy than the blackbody spectrum in vacuum. ", + "introduction": "\\label{sec:In} Although plasma is the main constituent of the primordial universe, very few previous studies deal directly with plasma phenomena. The effect of a plasma in cosmology normally has been studied with respect to the origin of the magnetic field. For example, the study of Harrison \\cite{har} elaborates a model of the origin of the magnetic field due to turbulence in the primordial plasma. There have been studies \\cite{mag} that analyze the effect of a magnetic field on primordial nucleosynthesis. Other studies, like that of Halcomb et al. \\cite{hal}, deduced the dispersion relation of waves, taking into account the expansion of the universe. There still lacks a general study of plasma phenomena related to cosmology. A plasma in thermal equilibrium, sustains fluctuations of the magnetic field (even for a non-magnetized plasma). This study is concerned with the study of the magnetic field spectrum in the primordial plasma. The electromagnetic fluctuations in a plasma has been made in numerous works, including those of Dawson \\cite{daw}, Rostoker et al.\\cite{ros}, Sitenko et al. \\cite{sit1}, and Akhiezer \\cite{akh1}. Most of the results are compiled in the books of Sitenko and Akhiezer et al. \\cite{sit2,akh2}. Little attention has been given to the question of how the magnetic field spectrum looks in a plasma. A naive answer to this question might be that it is a blackbody spectrum with a cut-off at the plasma frequency, knowing that photons only propagate in a plasma for $\\omega > \\omega_{p}$, where $\\omega_{p}$ is the plasma frequency. This is not true however, due to the magnetic fluctuations of the plasma. Cable and Tajima and Tajima et al. \\cite{tc1,tc2,tc3} performed a broad study of the magnetic field fluctuations in a plasma. They based their analyses on the {\\it Fluctuation-Dissipation Theorem}. They were concerned, in particular, with the low-frequency spectrum of fluctuations, because, as they pointed out, no expression exists. The {\\it Fluctuation-Dissipation Theorem} \\cite{sit2,akh2} predicts the intensity of electromagnetic fluctuations. The intensity of such fluctuations is highly dependent on how the plasma is described, in particular, on the dissipation mechanisms used. Cable and Tajima and Tajima et al. \\cite{tc1,tc2,tc3} studied the magnetic field fluctuations, for several cases. Two of their descriptions concern the primordial plasma which we are interested in, which is an isotropic, non-magnetized and non-degenerate plasma: a) a cold, gaseous plasma and b) a warm, gaseous plasma described by kinetic theory. In their study, Cable and Tajima (hereafter CT) in case (a) used the {\\it cold plasma} description with a constant collision frequency. In case (b) they analyzed the spectrum of fluctuations only for low frequencies, with the {\\it warm plasma} description for phase velocity $\\omega/k$ less or equal to the thermal velocity of the electrons, $v_{e}$ and the ions, $v_{i}$ in a collisionless description. For the {\\it cold plasma} description the spectrum that they obtain has a large zero-frequency peak. As the frequency is increased, the spectrum first drops below the blackbody spectrum in vacuum, then becomes the blackbody spectrum at high frequencies. In case (b), for the {\\it warm plasma} description, the analyses was made only for the low frequency regime and they argued that the zero-frequency peak is present as well. They argue that the energy contained in the peak is approximately equal to the energy {\\it lost} by the plasma cut-off effect. In order to obtain a correct magnetic field spectrum, it is necessary to describe the plasma in the most complete way as possible, taking into account thermal and collisional effects in a unified description. In this study we extend the pioneering work of CT, presenting a model that includes, in a unified description, collisional and thermal effects. Our model is based on kinetic theory incorporating thermal effects for all frequencies and wave numbers (not only for $\\omega/k \\leq v_{e},~v_{i}$). In order to describe the collisions that exist in the plasma, we used a model collision term. This collision term describes binary collisions, as used in the work by CT. In this way, we extend the previous model describing thermal and collisonal effects for all frequencies and wave numbers. Their description, the {\\it cold plasma} and the {\\it warm plasma} description in the collisionless case, are special cases of this model. (A pure collisionless treatment is unreal, such as case (b) extended to all frequencies, since if there were no collisions, then only Cherenkov emission could produce fluctuations, and there are no particles traveling fast enough to emit light waves.) However, for a fully ionized plasma as is our case, a treatment that takes into account collisions in a more complete way is necessary. Our model, an extention of the CT model, describes the basic features of a kinetic description. We present in Section \\ref{sec:mag} the general expressions of the magnetic field fluctuations based on the {\\it Fluctuation-Dissipation Theorem}. We review the {\\it cold plasma} description in Section \\ref{sec:co} and the {\\it warm plasma} description in the collisionless case, in Section \\ref{sec:war}. In Section \\ref{sec:dis} we present a general discussion and criticism of the assumptions made by CT. In Section \\ref{sec:mod} we present our model. Finally, in Section \\ref{sec:con}, we discuss the results and present our conclusions. ", + "conclusions": "\\label{sec:con} The magnetic field spectrum can be deduced from the fluctuations of the magnetic field described by the {\\it Fluctuation- Dissipation Theorem} and it is highly dependent on the way the plasma is described. We discussed the {\\it cold plasma} description and the {\\it warm plasma} description in the collisionless case studied by CT. We showed that $x_{cut}$ is much smaller than $x_{max}$ and $x_{lim}$, where $x_{lim}=k_{lim}c/\\omega_{pe}$ used in Sec. VII of CT (arguing that for $k > k_{lim}$ the plasma has a negligble effect). $x_{max}$ is the cut-off used in treating binary collisions, $(x_{max})^{-1}$ being the distance of closest approach between a test particle and an electron in a plasma (divided by $c/\\omega_{pe}$). We also showed that the {\\it Fluctuation-Dissipation Theorem} contains the eigenfrequencies of the plasma (in the transverse case), the photons. Using the {\\it cold plasma} description with the upper limit $x_{upper}=x_{max}$, we obtain the blackbody spectrum at high frequencies naturally, without the necessity of a {\\it link} procedure used by CT. For this case, the valley disappears and the curve is above the blackbody spectrum in vacuum. The calculations were made for two types of primordial plasmas: The electron-positron plasma at the beginning of the Big Bang nucleosynthesis; and the electron-proton plasma at lower temperatures. The manner to obtain the entire magnetic field spectrum is analyzing the magnetic field fluctuations. This is the only way to obtain information, not only about modes that propagate (i.e., photons), but also modes that do not propagate. The modes that do not propagate appear not only at low frequencies but also at high frequencies due to the correlations in the plasma. Only at very high frequencies does the photon contribution dominate the magnetic field spectrum. The argument used by CT, that the energy under the peak is almost equal to the energy ``stolen'' by the plasma cut-off effect of the blackbody in vacuum, is incorrect. There is no reason why we have to have the same energy as the blackbody spectrum in vacuum for photons for $\\omega < \\omega_{pe}$, since the photons have a different dispersion relation than the fluctuations in the plasma. In fact, using a upper limit $x_{upper}=x_{max}$ in the {\\it cold plasma} description, for example, the spectrum obtained is above the blackbody spectrum in vacuum. The reason why the collective modes of the plasma can have more energy for $\\omega \\leq \\omega_{p}$ than the photons in vacuum, can be understood as follows. Photons are massless bosons with the dispersion relation $\\omega^{2}=k^{2}c^{2}$. For the energy interval, $0 \\leq \\omega \\leq \\omega_{p}$, the wave number interval is $k=0$ to $k$ equal to $\\omega_{p}/c$. A relatively small amount of phase space is involved. For the collective motions of the plasma, in general, we have a larger amount of phase space. For example, for plasmons with energy $\\omega \\sim \\omega_{p}$, the amount of phase space extends to a maximum $k$ of $k_{D} \\cong \\omega_{p}/v_{t}$, which is greater than $\\omega_{p}/c$ for the photons. In general, for a given frequency for $\\omega < \\omega_{p}$, the greater phase available to the collective modes of the plasma (than that of the photons) implies more energy, or a higher spectrum. We presented a model that incorporates, in the same description, thermal and collisional effects. We used the Vlasov equation with the BGK collision term. This collision term describes the binary collisions in the plasma. A model that takes into account collisions in a more complete way is necessary. For a fully ionized plasma it is necessary to use the Fokker-Planck collision term that takes into account the effect of the microscopic fields. Due to the complexity of the solution of the kinetic equation with such a collision term, we used the BGK collision term. This model, an extention of the CT model, describes the basic features of a kinetic description. As we noted before, the results are very dependent on the cut-off chosen. Using $x_{max}$ as the cut-off, consistent with the collision term used, we obtain results that differ from the CT results. The final magnetic spectrum of a non-magnetized plasma in thermal equilibrium has the following characteristics: a) The peak intensity found by CT for frequencies $\\omega \\sim 0$ decreases; b) The blackbody is obtained naturally for high frequencies; and c) The magnetic spectrum has more energy than the blackbody photon spectrum in vacuum, in particular, for frequencies, $\\omega \\leq 4\\omega_{pe}$, where $\\omega_{pe}$ is the electron plasma frequency. \\begin{center} {\\bf ACKNOWLEDGMENTS} \\end{center} The authors would like to thank Swadesh Mahajan for useful suggestions, especially concerning the BGK collision term. The authors also would like to thank Arthur Elfimov for helpful discussions and the anonymous referee for useful comments and suggestions. M.O. would like to thank the Brazilian agency FAPESP for support and R.O. would like to thank the Brazilian agency CNPq for partial support." + }, + "9708/astro-ph9708121_arXiv.txt": { + "abstract": "We have evaluated a systematic effect on counts-in-cells analysis of deep, wide-field galaxy catalogues induced by the evolution of clustering within the survey volume. A multiplicative correction factor is explicitly presented, which can be applied after the higher order correlation functions have been extracted in the usual way, without taking into account the evolution. The general theory of this effect combined with the ansatz describing the non-linear evolution of clustering in simulations enables us to estimate the magnitude of the correction factor in different cosmologies. In a series of numerical calculations assuming an array of cold dark matter models, it is found that, as long as galaxies are unbiased tracers of underlying density field, the effect is relatively small ($ \\simeq 10\\%$) for the shallow surveys ($z <0.2$), while it becomes significant (order of unity) in deep surveys ($z \\sim1$). Depending on the scales of interest, the required correction is comparable to or smaller than the expected errors of on-going wide-field galaxy surveys such as the SDSS and 2dF. Therefore at present, the effect has to be taken into account for high precision measurements at very small scales only, while in future deep surveys it amounts to a significant correction. ", + "introduction": "Cosmological observations are necessarily carried out on a null hypersurface or a light-cone. At low redshifts ($z<0.1$), this can be regarded as to provide information on the constant-time hypersurface ($z=0$) which is a quite conventional implicit approximation underlying cosmological studies using the galaxy redshift surveys. When the depth of the survey volume exceeds $z\\sim 0.1$, however, this approximation breaks down, and one should simultaneously take account of the intrinsic evolution of galaxy clustering and the light-cone effect in addition to any other selection effect in interpreting the data. This is indeed the case for the on-going wide-field surveys of galaxies including 2dF (2-degree Field Survey) and SDSS (Sloan Digital Sky Survey). To our knowledge, the first quantitative consideration of the light-cone effect is made by Nakamura, Matsubara, \\& Suto (1998) who derived the systematic bias in the estimate of $\\beta \\approx \\Omega_0^{0.6}/b$ from magnitude-limited surveys of galaxies combining the cosmological redshift distortion effect (Matsubara \\& Suto 1996) and the evolution of galaxy clustering within the survey volume. In this paper, we examine the light-cone effect on higher-order statistics of galaxy clustering, considering counts-in-cells analysis specifically. Let us consider first the higher-order statistics on the idealistic constant-time hypersurface. Denote the volume averaged $N$-th order correlation functions at a redshift $z$ by $\\overline{\\xi}_N(R;z)$, where $R$ is the comoving smoothing length, and introduce the normalized higher-order moments $S_N(R;z) \\equiv \\overline{\\xi}_N(R;z)/[\\overline{\\xi}_2(R;z)]^{N-1}$. The hierarchical clustering ansatz states that $S_N(R;z)$ is constant and independent of the scale $R$. This is a good approximation in nonlinear regimes, although small but definite scale-dependence is clearly detected from N-body experiments (Lahav et al. 1993; Suto 1993; Matsubara \\& Suto 1994; Suto \\& Matsubara 1994; Jing \\& B\\\"{o}rner 1997). In addition, perturbation theory predicts that $\\overline{\\xi}_N(R;z)$ evolves in proportion to $\\left[\\overline{\\xi}_2(R;z)\\right]^{N-1}$, and therefore $S_N(R;z)$ is independent of time, i.e. it is constant with respect to $z$. The next section describes the general theory of the light cone effect on $S_N(R;z)$ defined above. Using the ansatz by Jain, Mo, \\& White (1995; hereafter JMW), \\S3 evaluates the appropriate correction in an array of cold dark matter (CDM) models. Finally, \\S4 summarizes the results and discusses the implications for redshift surveys. ", + "conclusions": "" + }, + "9708/astro-ph9708009_arXiv.txt": { + "abstract": "We examine the effect of different dark matter halo potentials on the morphology and kinematics of tidals tails in a merger model of NGC 7252. We find that models of merging galaxies with low halo masses of $M_h \\sim 4-8 M_{disk+bulge} (M_{db})$ can fit the observed morphology and kinematics of the NGC 7252 tails while galaxies with high mass halos ($M_h \\sim 16-32 M_{db}$) fail in this respect. In high mass models, the deep potential only allows weakly bound disk material (stars or gas) at $R$ \\gtsima 5 disk scale lengths to be ejected in tidal tails which tend to fall back onto the parent galaxies before the final merger. Galaxies with massive, low density halos are somewhat more successful at ejecting tidal debris during mergers, but still have difficulties recreating the thin, gas-rich tails observed in NGC 7252. Our models suggest upper limits for the dark halo masses in the NGC 7252 progenitor galaxies of roughly $M_h$ \\ltsima 10 $M_{db}$. We note, however, that our calculations have focused on the rather idealized case of the isolated merging of galaxies with distinct dark matter halos; calculations which employ more realistic (``cosmological'') initial conditions are needed to fully explore the use of tidal tails in constraining dark matter in galaxies. ", + "introduction": "While the existence of dark matter halos around galaxies seems well demonstrated through such diverse kinematic tracers as disk galaxy rotation curves (\\eg Rubin \\etal 1982, 1985; Kent 1987), satellite galaxies and globular clusters (Zaritsky \\etal 1989; Zaritsky \\& White 1994; Kochanek 1996), and hot gas around ellipticals (\\eg Forman, Jones, \\& Tucker 1985), the radial extents and total masses of these halos remains poorly constrained. The rotation curves of spiral galaxies generally probe the mass distribution out to only $\\sim$ 10 disk scale lengths, while estimates based on more distant satellites are statistical in nature, and sensitive to selection effects and assumptions about orbital kinematics (see, \\eg Zaritsky \\& White 1994; Kochanek 1996). Taken together, these lines of argument generally suggest that galaxies with circular velocities similar to the Milky Way have halos with masses M$_{\\rm halo} \\sim 10^{12}$ M$_{\\sun}$ and extend beyond $\\sim 100$ kpc. In an attempt to constrain dark matter halos in an independent manner, Dubinski, Mihos, \\& Hernquist (1996; hereafter DMH) showed that the morphology of tidal tails produced in galaxy collisions depends sensitively on the potential of the galaxies. The use of tidal debris to probe dark matter halos was originally proposed by Faber \\& Gallagher (1979), and later emphasized by White (1982) and Negroponte \\& White (1983), who argued that galaxies with massive dark halos might have difficulty forming long tidals due to their deeper potential wells. Barnes (1988) tested these ideas using self-consistent models and noted a weak anticorrelation between the masses of the dark halos of the colliding galaxies and the amount of material ejected in the tidal tails. However, Barnes used galaxies with relatively low mass halos (\\hdbs of 0, 4, and 8:1) and concluded that tidal tails are generically easy to produce. Employing halos much more massive than those used by Barnes, DMH demonstrated that if one considers halos as massive and as extended as some observations suggest, the formation of long tidal tails is sharply curtailed. Given that a number of merging galaxies display long tidal tails (\\eg NGC 4038/39, NGC 7252, the Superantennae), DMH argued that such galaxies must have \\hdbs on the order of 10:1 or less. DMH's study focussed primarily on the {\\it morphology} of tidal tails produced in various galaxy encounters. However, the kinematics of tidal debris may also provide additional constraints which can be compared directly to observed HI kinematics of merging galaxies (\\eg Hibbard 1994; Hibbard \\& Yun 1997). The kinematics of tidal debris trace the encounter by following trajectories determined in large part by the orbital energy and interaction geometry. Hibbard \\& Mihos (1995; hereafter HM) used the morphology and kinematics of the extended tidal tails around NGC 7252 to reconstruct the dynamical history of this merger. Their model constrained the orbital geometry and viewing angle of the encounter as well as the merging timescale, and predicted future infall rates of material currently populating the tidal tails. However, HM used a single \\hdb of 5.8:1 in their simulations and did not investigate in any detail the sensitivity of their results to the internal structure of the merging galaxies. In what follows, we consider both the morphology and kinematics of tidal tails formed from collisions of galaxies with various halo properties, to provide additional constraints on the amount of dark mass around galaxies as well as to understanding the long term evolution of the tidal debris. Our first step is to examine the kinematics of tidal tails in general, employing models with extended disks of material to trace the dynamics of the loosely bound material from which tidal tails are drawn. We then reanalyze NGC 7252, comparing the morphology and kinematics of the observed tidal tails to those produced in the models. Finally, we address the robustness of the DMH results by considering models with rotating halos and ones with high mass halos having lower central densities and shallower potentials. ", + "conclusions": "The models presented here expand on the work of DMH and HM in two respects. First, we have followed the evolution of material initially located at very large distances in the progenitor galaxies, allowing us to examine the detailed kinematics and morphology of this loosely bound material. In simulations involving mergers of galaxies with increasing halo mass, the tidal debris is drawn primarily from particles located at increasingly large radii within the progenitors, and more of this material remains bound to the merger remnant. The tidal tails which form immediately in mergers involving very massive halos quickly fall back into the galaxies, so that they are no longer visible by the time the galaxies merge. These models reinforce the claim of DMH that observed merger remnants with long tidal tails must have formed from progenitors with relatively small \\hdbs\\llap. Second, we have explored a variety of halo models in an attempt to reproduce the observed tidal tail morphology and kinematics of NGC 7252, placing some constraints on the dark matter distribution around merging galaxies. Again, the observations are best fit using mergers of galaxies with \\hdbs in the range of 4--8. However, we find that precise estimates are difficult because of degeneracies in the solution, as originally suggested by HM. The tendency of tail material to be drawn from larger initial radii with increasing halo mass (Figure \\ref{fig-rinitfin}) has implications for recent QSO absorption line studies. Because of abundance gradients in galactic disks, the ability of galaxy interactions to expel metal-rich material to large distances will be sensitive to the mass distributions of dark matter halos. Accordingly, the metallicity of tidal tails may be used as another constraint on the masses of galaxy halos. Absorption lines produced by tidal debris from intervening galaxies have been identified in several QSO spectra (\\eg Sargent \\& Steidel 1990; Norman \\etal 1996); while metallicity estimates are uncertain, if such systems prove to be reasonably metal-rich, it would support the idea that galaxy halos may be less massive than other observational estimates. This argument is similar to the suggestion that low-mass galaxies are more able to eject metal-rich material into the IGM through starburst-driven superwinds (\\eg Heckman, Armus, \\& Miley 1990); in this case, however, the energy involved in expelling metal-rich material comes from tidal encounters rather than starburst winds. With the extended disk models showing that galaxies with more massive halos may eject significant amounts of extended HI into tidal tails, the possibility arises that subsequent star formation could convert this gas into stars and produce the long {\\it optical} tidal tails observed in some merger remnants. However, to match observed tidal tails, which contain as much as 10--20\\% of the blue luminosity of merging galaxies, this star formation must be prodigious. For example, if the optical light in the tidal tails of NGC 7252 were to come from stars formed {\\it in situ}, then $\\sim$ 80\\% of the gas in the tails must have been converted into young stars at several M$_{\\sun}$ yr$^{-1}$ to reproduce the total blue luminosity and observed (remaining) gas content of the tails. While some star formation is observed in tidal tails, it typically occurs in a few star forming clumps rather than being smoothly distributed, and at much lower rates. Furthermore, the observed colors of tidal tails are more representative of material stripped from the inner disks of galaxies, rather than young stellar populations (Schombert \\etal 1990). We have also investigated interactions using galaxies containing halo models with different internal kinematics and mass distributions. Maximally rotating halos ($\\lambda = 0.20$) have no discernible effect on the evolution of a Model C merger and so the amounts of rotation inferred in halos from cosmological arguments ($\\lambda=0.05$) are unlikely to be important for determining the evolution of merging galaxies. Mergers of galaxies with high mass, extended halos (Model E) are able to eject more material into tidal debris, due to their shallower potential wells. However, the tidal debris is very diffuse and suffers significant orbit crossing, making it difficult to identify with the gas-rich tidal features in objects like NGC 7252. However, our models have only examined one representation of a low density halo model and a more systematic study of the effects of low density halos is warranted, especially in light of observational (Casertano and van Gorkom 1991; Persic \\etal 1996) and theoretical (Navarro \\etal 1996) results which suggest that the most luminous spiral galaxies may have declining rotation curves. The models described here address many of the loopholes left open by DMH, supporting the conclusion that long tidal tails are a signature of compact, low-mass halos in the progenitors to a merger. Nonetheless, one limitation of the models still remains: their rather idealistic initial conditions -- galaxies with individual, distinct dark matter halos merging in the absence of any background potential. Given recent cosmological simulations which show that galaxy halos often merge before their luminous galaxies do (\\eg Katz, Hernquist, \\& Weinberg 1992), our simulations may be an oversimplified version of galaxy mergers. In fact, dynamical friction against a background dark matter distribution may hasten merging, allowing massive galaxies to merge before their tidal tails have fallen back into the galaxies. The development of tidal tails will also be affected by the interaction between {\\it three} potential wells (two galactic and one background); the structure and kinematics of the resultant tidal features is difficult to assess without detailed modeling. The next consistency check on our results would therefore be to examine mergers in a more ``cosmological'' setting, in which galaxies merge in a more diffuse ``sea'' of dark matter. In principle, the statistics of tidal tails could be used to infer the properties of dark matter halos. In practice, however, this may be difficult to achieve. While long tidal tails suggest low mass halos, the converse may not necessarily be true -- the lack of observed tidal tails may have been the result of an unequal mass merger, an unfavorable orbital geometry (\\ie a retrograde merger), unsuitable progenitors (ellipticals or S0's), or rapid fading in surface brightness due to kinematic evolution of the tails (\\eg Mihos 1995). Furthermore, sample selection would be fraught with bias -- as mergers are generally identified through the presence of tidal debris, care would need to be taken to ensure the sample would not be skewed towards low mass systems with obvious tidal tails. While a statistical constraint on the dark matter content of galaxies using tidal tails may be problematic, the implications for individual systems seem more clear. For merging galaxies such as NGC 7252, the Antennae, and the Superantennae, the presence of long tidal tails is difficult to reconcile with massive dark matter halos, unless perhaps the halos are very extended and diffuse. While most kinematic probes of the mass distribution in galaxies (\\ie rotation curves, satellite kinematics) yield lower limits on halo masses, the results described here suggest some of the first {\\it upper limits} on the dark matter content of galaxies. As such, it is of immediate interest to test these concepts using both numerical simulation and detailed observational studies of the morphology, kinematics, and metallicity of tidal tails. As coherent kinematic tracers at the largest radius, tidal tails may yet unveil the dark matter halos in which galaxies live." + }, + "9708/astro-ph9708186_arXiv.txt": { + "abstract": "Models are presented for CO rotational line emission by high redshift starburst galaxies. The influence of the cosmic microwave background on the thermal balance and the level populations of atomic and molecular species is explicitly included. Predictions are made for the observability of starburst galaxies through line and continuum emission between $z=5$ and $z=30$. It is found that the Millimeter Array could detect a starburst galaxy with $\\sim 10^5$ Orion regions, corresponding to a star formation rate of about $30 \\rm M_\\odot yr^{-1}$, equally well at $z=5$ or $z=30$ due to the increasing cosmic microwave background temperature with redshift. Line emission is a potentially more powerful probe than dust continuum emission of very high redshift objects. ", + "introduction": "Searches for CO emission from cosmological objects have had some success. Examples include the IRAS source F10214+4724 at $z= 2.29$ (Solomon, Downes and Radford 1992), the clover-leaf galaxy at $z=2.6$ (Barvainis et al.~1994), and a quasar at $z=4.69$ (Omont et al.~1996). These searches establish that large amounts of molecular gas are present at high redshifts. This is to be expected in the light of the recent detections of high redshift Lyman break galaxies, e.g.~in the Hubble Deep Field, which are actively forming stars at $z\\sim 3-4$ (Steidel et al.~1996). Since star formation is ultimately driven by the collapse of cold molecular clouds, the occurrence of active star formation must be reflected by the physical structure of the interstellar medium (ISM). The detection of molecular gas and dust at high redshift therefore provides an excellent probe of the stellar processes occurring in cosmological objects. In fact, the metallicity and physical state of the high redshift ISM provides indirect constraints on the star formation rate and hence on models of galaxy formation. In the next decade, instruments will come on line to explore the infrared and millimeter regions of the spectrum with the goal of detecting objects at very high redshift: NGST, FIRST, and the MMA. The latter will have the ability to detect emission lines fluxes at the milliJansky level around wavelengths of a few millimeters, and will be ideal for a search for highly redshifted molecular lines, in particular for CO lines. In this work the excitation of the CO molecule will be investigated in the presence of a warm Cosmic Microwave Background (CMB) and the subsequently altered thermal balance. The aim is to determine which molecular lines are best suited for the detection of star-forming primordial galaxies, and to assess up to which redshift such measurements are feasible with the planned next generation of observatories. ", + "conclusions": "The total molecular gas mass in Orion is about $2\\times 10^5\\,\\rm M_\\odot$. The star formation rate in Orion is probably about $3\\times 10^{-4} \\rm M_\\odot \\, yr^{-1}$, according to recent observations of the Orion Nebula Cluster (Hillenbrand 1997). Hence our putative protogalaxy containing $3\\times 10^5$ Orions is forming stars at a rate of $90\\rm \\, M_\\odot \\, yr^{-1}$ from a gas reservoir of about $6\\times 10^{10} \\, \\rm M_\\odot$. These numbers are merely meant to be representative for a protospheroid of modest mass, amounting to only 10 percent or so of the characteristic stellar mass $M_\\ast$ (as defined by the galaxy luminosity function) of an elliptical galaxy if the star formation efficiency of the gas is about 50 percent. Our results can of course be trivially rescaled. The local ratio of FIR luminosity to molecular gas mass of about $20 \\rm L_\\odot/M_\\odot$ is in the range associated with luminous starbursts (Sanders, Scoville \\& Soifer 1991). This ratio will vary with location. A more typical value is about half of this, for then the FIR luminosity is inferred to be about $6\\times 10^{11} \\rm L_\\odot$ and scales appropriately with the global star formation rate for the Milky Way. Hence our model of $3\\times10^5$ Orions matches in luminosity, molecular mass, and star formation rate what would be expected from a luminous starburst. The column-averaged population distribution of CO is shown in Figure 1a (top panel) as a function of redshift. Line intensities averaged over the source are shown in Figure 1b. The peak in the level population shifts with increasing redshift to J=5-6, with critical densities $n_{\\rm c}\\sim 10^5$ cm$^{-3}$ and excitation energies of $\\sim 80-100$ K for the corresponding transitions, at $z\\sim 30$, before becoming thermalized by the CMB. Note that the line luminosities increase as the CMB becomes hotter because the level excitation increases with increasing J until a CMB temperature $T\\sim 90 \\rm K$ is attained and higher J lines cannot be pumped in an Orion-like environment due to their high critical density. The predicted line fluxes are shown in Figure 2. Because of the enhancement of line fluxes by the CMB, we find the remarkable result that beyond $z\\sim 5$, the predicted line intensities are almost independent of redshift. Solomon, Radford and Downes (1992) previously noted that the CO(J=$3-2$) line is always comparable in strength to the CO(J=$1-0$) line up to $z\\sim 2$ due to the warmer microwave background. We find that as higher rotational levels are populated at higher redshift, one can measure a starburst at $z\\sim 30$ as easily as at $z\\sim 5$. In fact, the measurement may be even easier since the emission peaks at longer wavelengths. The upper and lower grey curves in Figure 2 indicate the range in fluxes for Orion-like regions with metallicities which are 4 times higher and 4 times lower than solar. The overall effect is a factor of a few, which indicates that the effect of the CMB on the CO line intensities is robust. All lines of interest are in the millimeter range, and the milliJansky fluxes predicted for our $3\\times 10^5$ Orion model are within the capability of the MMA to eventually be measurable. Of course, one would have to search a considerable amount of sky. If the duration of a starburst is $\\sim 10^8\\alpha$yr, and $0.01\\beta$ is the fraction of early-forming spheroids of mass above $0.1 M_\\ast$ relative to present-day ellipticals, one would at most expect $\\sim 100 \\alpha\\beta$ per sq.~deg.~at $z\\sim 5$, and an order of magnitude fewer at $z\\sim 30$. It is of interest finally to compare the FIR emission with our predicted line fluxes. In Figure 3, we show the continuum fluxes estimated for the Orion dust model. As one proceeds to redshifts above $\\sim 5$ the dust emission becomes inreasingly harder to detect relative to the line emission. At 100 $\\mu$m, the typical continuum flux is about 1 mJy for the $z=5$ model. It will be a challenge even for FIRST to detect such a weak signal, requiring days of integration. Conversely, the MMA sensitivity limit for spectral lines is expected to be about 1mJy, as compared to our predicted fluxes of several mJy for our model of what is only in effect a modest starburst by ultraluminous infrared galaxy standards. We comment in closing that the main emphasis has been on the rotational lines of CO because of the fortunate energy level spacing for increasing redshifts. Other molecular species such as CS, HCO$^+$ and HCN will be bright emitters as well, but do not couple as favorably with the CMB. Therefore, their fluxes will be strongly reduced by cosmological redshift effects and not easily detectable at high redshift. MS acknowledges with gratitude support of NASA grant NAGW-3147 from the Long Term Space Astrophysics Research Program. The research of JS has been supported in part by grants from NASA and NSF. JS also acknowledges with gratitude the hospitality of the Physics Department at the Johns Hopkins University as a Bearden Visiting Professor, and the Institute of Astronomy at Cambridge as a Sackler Visiting Astronomer. \\newpage" + }, + "9708/astro-ph9708135_arXiv.txt": { + "abstract": "We present the first results of a sub-millimeter survey of distant clusters using the new Sub-mm Common-User Bolometer Array (SCUBA) on the James Clerk Maxwell Telescope. We have mapped fields in two massive, concentrated clusters, A370 at $z=0.37$ and Cl\\,2244$-$02 at $z=0.33$, at wavelengths of 450 and 850\\,$\\mu$m. The resulting continuum maps cover a total area of about 10\\,arcmin$^2$ to 1$\\sigma$ noise levels less than 14 and 2\\,mJy\\,beam$^{-1}$ at the two wavelengths, 2--3 orders of magnitude deeper than was previously possible. We have concentrated on lensing clusters to exploit the amplification of {\\it all} background sources by the cluster, improving the sensitivity by a factor of 1.3--2 as compared with a blank-field survey. A cumulative source surface density of $(2.4\\pm 1.0) \\times 10^3$ degree$^{-2}$ is found to a 50\\% completeness limit of $\\sim 4$\\,mJy at 850\\,$\\mu$m. The sub-mm spectral properties of these sources indicate that the majority lie at high redshift, $z>1$. Without correcting for lens amplification, our observations limit the blank-field counts at this depth. The surface density is 3 orders of magnitude greater than the expectation of a non-evolving model using the local {\\it IRAS} 60\\,$\\mu$m luminosity function. The observed source counts thus require a substantial increase in the number density of strongly star-forming galaxies in the high-redshift Universe and suggest that optical surveys may have substantial underestimated the star formation density in the distant Universe. Deeper sub-mm surveys with SCUBA should detect large numbers of star-forming galaxies at high redshift, and so provide strong constraints on the formation of normal galaxies. ", + "introduction": "Surveys of the local Universe have shown that a third of the total bolometric luminosity is emitted at sub-mm and far-infrared wavelengths as a result of reprocessing of star-light by dust (Soifer \\& Neugebauer 1991). Moreover, some of the most vigorous star-forming galaxies in the local Universe are also those in which the effects of dust obscuration are most significant. While there have been striking advances in the identification of `normal' galaxies at high redshift ($z\\sim2$--4) using Lyman-dropout techniques (Steidel \\et 1996), such approaches are insensitive to highly obscured star-forming galaxies at these epochs. The presence of at least modest amounts of dust in distant proto-galaxies, especially forming spheroids, is expected given the highly metal-enriched ISM which must be present during their formation (e.g.\\ Mazzei \\& de Zotti 1996). Thus techniques sensititive to obscured distant galaxies should be used to investigate the formation of these populations in the early Universe. Sensitive sub-mm observations present the first opportunity to detect dusty star-forming galaxies at high redshift. At wavelengths around 100\\,$\\mu$m, the bulk of the luminosity of normal, star-forming galaxies is reprocessed star-light from dust and so observations in the sub-mm band can provide robust estimates of both the dust mass and total star formation rate in a galaxy. Furthermore, the negative {\\it K}-correction at wavelengths $\\lambda \\gs 400$\\,$\\mu$m means that sub-mm observations select star-forming galaxies at $z\\gs 1$ in an almost distance-independent manner, providing an efficient method for the detection of star-forming galaxies at very large redshifts, $z \\ls 10$, and hence studying galaxy evolution (Blain \\& Longair 1993, 1996 --- BL96; Blain 1997; Eales \\& Edmunds 1997; Franceschini \\et 1997; Guiderdoni \\et 1997). \\begin{table*} {\\scriptsize \\begin{center} \\centerline{Table 1} \\vspace{0.1cm} \\centerline{SCUBA observations of A370 and Cl\\,2244$-$02} \\vspace{0.3cm} \\begin{tabular}{lccccccc} \\hline\\hline \\noalign{\\smallskip} {Target} & {R.A.} & {Dec.} & {$\\lambda$} & {Exposure time} & {Area} & {Flux density} \\cr ~ & {(J2000)} & {(J2000)} & {($\\mu$m)} & {(ks)} & {(arcmin$^2$)} & {1$\\sigma$ (mJy)} \\cr \\hline \\noalign{\\smallskip} Cl\\,2244$-$02 & 22~47~11.9 & $-$02~05~38 & 450 & 23.0 & 4.00 & 14.0 \\cr ~ & ~ & ~ & 850 & 23.0 & 5.40 & 1.9 \\cr \\noalign{\\smallskip} A370 & 02~39~53.0 & $-$01~35~06 & 450 & 25.7 & 4.35 & 13.3 \\cr ~ & ~ & ~ & 850 & 25.7 & 5.40 & 1.8 \\cr \\noalign{\\smallskip} \\noalign{\\hrule} \\noalign{\\smallskip} \\end{tabular} \\end{center} } \\vspace*{-0.8cm} \\end{table*} Most published sub-mm studies of distant galaxies have targeted atypical galaxies (e.g.\\ radio-loud galaxies, Ivison \\et 1998). We report here the first deep sub-mm survey to probe the nature of normal galaxies at moderate and high redshift, $z\\gs 0.5$. In this study we have attempted to maximise the available sample of distant galaxies by concentrating on fields in moderate-redshift clusters. While the dominant spheroidal populations of these clusters are expected to be quiescent in the sub-mm band, the in-fall of field galaxies associated with the growth of the clusters (Smail \\et 1997) means that these fields may contain over-densities of moderate-redshift star-forming galaxies, as compared with `blank' field surveys. The main attraction of the clusters observed here, however, is that they are strong gravitational lenses, magnifying any source lying behind them (Blain 1997). Given the expected steep rise in the sub-mm counts (BL96), amplification bias could increase the source counts in these fields by a substantial factor, with a maximum predicted surface density of about 10 sources per SCUBA field down to 1\\,mJy at 850\\,$\\mu$m (Blain 1997). Moreover, by targeting those clusters that contain giant arcs, images of distant field galaxies magnified by factors of 10--20, we can also obtain otherwise unachievable sensitivity ($\\ls 0.1$\\,mJy at 850\\,$\\mu$m) on the dust properties of a few serendipitously-positioned normal galaxies at high redshift. The angular scales of the region where these highly magnified high-redshift galaxies are found is also well-matched to the SCUBA field-of-view. In the following sections we give details of the observations and their reduction, and discuss the results within the framework of current theoretical models of galaxy formation and evolution. We adopt $H_\\circ=50 $\\,km s$^{-1}$ Mpc$^{-1}$ and $q_\\circ = 0.5$. ", + "conclusions": "\\noindent{$\\bullet$} We have presented the first sub-mm survey of the distant Universe, deep enough that we should detect the evolving galaxies predicted by current theoretical models, while at the same time covering a sufficiently large area to be statistically reliable. We derive cumulative source counts of $(2.4\\pm 1.0) \\times 10^3$ degree$^{-2}$ down to 4\\,mJy at 850\\,$\\mu$m. \\noindent{$\\bullet$} The surface density of faint sources in the sub-mm far exceeds a simple non-evolving model using the locally observed 60\\,$\\mu$m galaxy luminosity function. Thus our observations require a substantial increase in the number density of strongly star-forming galaxies at $z\\gs 1$. \\noindent{$\\bullet$} Comparison of our observations with the predictions of simple parametric models implies that the luminosity density of the brightest sub-mm sources continues to increase out to $z> 1$. Models based upon the claimed properties of star-forming galaxies from optically-selected samples of distant galaxies (Madau \\et 1996) significantly underestimate the observed surface density of sub-mm sources. We suggest that such samples may be missing a considerable proportion of the star formation in dust-obscured galaxies at high redshift. We conclude that question of the evolution of the star formation density in the distant Universe and hence the epoch of galaxy formation is still very much open." + }, + "9708/astro-ph9708245_arXiv.txt": { + "abstract": "Core and conal emission beams of radio pulsars have been identified observationally (Rankin 1983,1993; Lyne and Manchester 1988), In the inner gap model (Ruderman and Sutherland 1975, hereafter RS75), the gap continually breaks down (sparking) by forming electron-positron pairs on a time scale of a few microseconds (RS75). This makes a large amplitude low frequency wave to be produced and which would be scattered by relativistic particles moving out from the gap. Under this assumption, Qiao (1988a,1992) presented an Inverse Compton Scattering (ICS) model for both core and conal emissions. This paper presented a development of the model. Retardation and aberration effects due to different radiation components emitted from different heights are considered. The luminosity of pulsar radio emission via the ICS process is discussed. Coherent emission by bunches of particles is adopted and which is adequate to explain pulsar radiation. The theoretical results are in agreement with observations very well. ", + "introduction": "It's convincing that the emission beams of a radio pulsar can be divided into two (core, inner conal) or three (plus an outer conal) emission components through careful studies of the observed pulse profiles and polarization characteristics (Rankin 1983a,1983b,1986,1990,1993; Lyne \\& Manchester 1988). Many pulsar profiles at meter wavelength are dominated by core components. In usual polar cap models of pulsars, it is difficult to get a central or ``core'' emission beam. Many current theoretical models can only get a hollow cone emission beams. Thus, it is needed to make an effort to get the core emission theoretically. Several authors such as Beskin et al. (1988), Qiao (1988a,b), Wang et al. (1989) presented some models for the core emission beam. If the binding energy per ion in the neutron star surface is as large as 10 kev then ions will not be released, a pole magnetospheric vacuum gap (inner gap) is formed (RS75). More accurate variational calculations (e.g. Hillebrandt \\& M\\\"uller 1976; Flowers et al. 1977; K\\\"ossl et al. 1988) have revised downward the binding energy to a few $kev$. According to RS75, primary particles are accelerated to sufficient energies in the gap and emit high energy $\\gamma$-quanta via curvature radiation (hereafter CR), which in turn, ignite pair production cascade to short out the acceleration potential (sparking). This is to say in the calculations for the cascade of the inner gap only CR process is taking into account. If the Inverse Compton Scattering (ICS) process in strong magnetic fields is taking into account the potential drops across the gap will be down, the ``binding energy difficulty'' will be released and the inner gap model (RS model) can still be sound, even if the binding energy downward to a few $kev$ (Zhang and Qiao 1996, here after ZQ96; Qiao and Zhang 1996, here after QZ96; Zhang 1996; Zhang et al.1997). The gap continually breaks down (sparking) by forming electron-positron pairs on a time scale of a few microseconds (RS75), this is inconsistent with the observed short-time scale structure (Hankins,1992). A very large amplitude low frequency wave would be associated with the sparking, which would be scattered by relativistic particles moving out (if the low frequency wave can propagate in the magnetosphere of the neutron star, we will discuss this a little below). Under the assumption that the observed radio emission are produced by the ICS process of the high energy particles (secondary) off the low frequency waves, we can get both core and two conal components (Qiao 1988a,1992). The different emission components are emitted from different heights (Qiao et al. 1992, Lin and Qiao 94, hereafter LQ96, and this paper, see bellow). In this paper we present a development of the work of Qiao (1988a,b;1992) which presented basic ideas and calculations for both core and two conal emission beams. Both observations (Rankin 1993) and the theory (Qiao, et al., 1992) show that the ``core'' emission is emitted at a place relatively close to the stellar surface, the ``inner cone'' is emitted at a lower height, and the ``outer cone'' is emitted at some greater height. If the different emission components are emitted at different heights, two effects, aberration and retardation effects, should be considered in the calculations. Our result shows that these effects move the apparent positions of the emission beams and change their shapes to be asymmetric to the magnetic axis. McCulloch (1992) examined about 20 triple profiles and found no example with the central component close to the leading component but many close to the trailing component, our calculations fit the result well. The basic idea, assumptions and the results of the calculations are presented in section 2; some theoretical respects and observational facts are discussed in sections 3. This is one of a series papers about radio pulsar emission, on the polarizations both linear and circular, the behavior of pulse profiles in different frequencies, $\\gamma$-ray emission and so on will be done later. ", + "conclusions": "\\subsection{Shape of emission beams and pulse profiles} The result in this paper shows that slow-rotation pulsars with long period such as $0.5s-1.0s$ or longer, are ``double-conal pulsars'' with both inner and outer cone beams. This is in agreement with the conclusion given by Rankin (1993). Rankin's table (1993) for {\\bf M} pulsars with five components shows, most of these 19 pulsars have long pulse periods. In the Fig.~\\ref{FigFreq}a, the opening angle (angular radius) becomes larger when the frequency increases (decreases) for the inner cone (outer cone). Fast pulsars, i.e. $P<0.3s$, may only have a core and one inner conal emission component (Fig.~\\ref{FigBeam3} and Fig.~\\ref{FigFreq}b). For this class of pulsars, we suppose that observations tend to get triple ({\\bf T}) profiles. In Fig.~\\ref{FigFreq}b, as the frequency increases, the opening angle for the inner cone always increases. In this case, at higher frequency (from line D to line C), we will get a wider pulse profile (Qiao 1992) in contrast to slow pulsars. This can be seen for some pulsars, such as PSR B1642-03 and PSR B1933+16 (Sieber et al. 1975). We must mention that the theoretical angular radius of the outer conal beam is larger than in previous calculations and the height of the outer conal beam emission region is larger than that of Rankin's. This may be related to the parameter $\\xi $ in Eq.(5). A detailed calculation will show that the controller $\\xi $ is determined by the energy loss of particles, which depends on the strength of the magnetic field and the thermal temperature at the surface of neutron stars (Zhang te al. 1997). \\subsection{Does the ``inner'' cone radius increase as the observed frequency increases?} One result of this paper is that the inner cone radius increases as the observed frequency increases. This result is supported by analysis of observations. Wu et al.(1992) present a method to deal with the structure of the mean pulse profiles of pulsars. With that method and multi-frequency observational data, a diagram of $f'-\\theta_\\mu$ was given, which is very similar to the result of our calculation (see Fig.~\\ref{FigFreq}) and Fig.2 of Qiao (1992). Further analysis of Rankin (1993) did not emphasize that the angular radius of the inner cone $\\rho_{inner}$ increases when the frequency increases. More analysis of this is needed. A very direct method to check the result of this paper is that: for pulsars with ``inner'' and ``outer'' cone (five components), the pulse profiles would become a triple (``inner'' cone and ``outer'' cone get together) with a smaller central component at very high frequency (Fig.~\\ref{FigFreq}, line A); And also become a triple (the inner cone and core get together) with stronger central component at low frequency (Fig.~\\ref{FigFreq}, line B). This is in agreement with the observations, {\\em e.g.} Izvekova et al. (1989). \\subsection{conclusion} Our result shows, ICS process is a possible radiation mechanism for radio pulsars since it can produce the emission beams naturally and well consistent with observations. Rankin (1993) showed that the angular radii of core, ``inner'' cone and ``outer'' cone at a given frequency is a function of P (and only P!). This is just the result of the calculations in this paper. In agreement with the results given by Rankin (1993), we conclude that those pulsars only with ``inner'' cone (core single and triple in Rankin's classification) are generally faster, those with ``outer'' cone (conal single and double) much slower, and the group of five-component ({\\bf % M}) pulsars falls in between the other two. This paper also supports that the ``inner'' cone is emitted at a lower height along a same group of field lines that produce the ``outer'' cone. The shapes of pulse profiles change with frequencies in agreement with some kind of pulsars (Qiao 1992). The retardation and aberration effects induce asymmetry and can be observable, our result fits with observations (McCulloch 1992). These two effects may also change the linear polarization position angle (Xu, et al. 1996). The coherent emission of the ICS process suggested in this paper is an efficient mechanism to produce observed luminosity, and which is also a mechanism to produce observed polarization characters (Qiao et al. 1997)." + }, + "9708/astro-ph9708073_arXiv.txt": { + "abstract": "We present results on the X-ray spectra of the radio-loud, high-polarization quasar, \\pks, based on new data obtained using $ASCA$, and from archival $ROSAT$ data. We find the X-ray spectrum obtained by $ASCA$ to be unusually hard, with the photon index, $\\Gamma$=1.30$\\pm$0.06, while the (non-simultaneous) $ROSAT$ data indicate a steeper spectrum with $\\Gamma$=1.9$\\pm$0.3. However, we find the X-ray flux at 1 keV to be within $\\sim$10\\% during both observations. Thus we suggest the most likely explanation is that there is a break (with $\\Delta \\Gamma \\geq 0.6$) in the underlying continuum at $\\sim$~0.7~keV. Although the sample of high-polarization quasars for which high quality X-ray spectra are available is small, flat X-ray spectra seem to be the characteristic of these objects, and they also appear to be harder than that of the other radio-loud but low-polarization quasars. The multiwavelength spectrum of \\pks~is similar to many other $\\gamma$-ray blazars, suggesting the emission is dominated by that from a relativistic jet. A big blue-bump is also seen in its multiwavelength spectrum, suggesting the presence of a strong thermal component as well. ", + "introduction": " ", + "conclusions": "" + }, + "9708/hep-ph9708383_arXiv.txt": { + "abstract": "s#1{ \\begin{center} {\\begin{minipage}{4.2truein} \\footnotesize \\parindent=0pt #1\\par \\end{minipage}}\\end{center} \\vskip 2em \\par} \\fussy \\flushbottom \\parindent 0.25in \\oddsidemargin 0.75in \\evensidemargin 0.75in \\topmargin=1in \\headheight=0.1in \\headsep= 0.2in \\footskip=0.3in \\footheight=0.3in \\textheight = 6.7in \\textwidth 4.7in \\defs{ The dynamical growth rate of bubbles nucleating in relativistic plasma in thermal first-order phase transitions is analyzed. The framework is a hydrodynamical model which consists of relativistic fluid and an order parameter field. The results of analytical approximations and numerical simulations coincide well. } \\clearpage \\pagestyle{plain} In thermal systems first-order phase transitions normally proceed via nucleation of bubbles of the new phase. Bubbles larger than a certain critical size begin to grow, whereas smaller bubbles will shrink. Langer's formula~\\cite{Langer69}$^{\\!,\\,}$\\cite{Langer73} for the nucleation rate of bubbles of the new phase is given by \\begin{equation} \\Gamma = \\frac{\\kappa}{2\\pi} \\Omega_0 \\: e^{-\\Delta F/T}, \\label{Gamma} \\end{equation} where $\\Gamma$ is the probability of nucleation per volume and time, $\\kappa$ a dynamical and $\\Omega_0$ a statistical prefactor, $\\Delta F$ the free energy difference of the system with and without the nucleating bubble, and $T$ the temperature. The purpose of this study is to investigate the dynamical prefactor $\\kappa$. Let the phase transition be driven by an order parameter field $\\phi(x)$. Initial growth rate of perturbations around an extremum configuration is given by the coefficient $\\kappa$: \\begin{equation} \\delta \\phi \\; \\propto \\; e^{\\kappa t}. \\label{deltaphi} \\end{equation} Similar relations hold for perturbations of energy density and fluid velocity. A parallel definition for the growth coefficient $\\kappa$ can be expressed in terms of the radius $R(t)$ of an expanding spherical bubble of the new phase, \\begin{equation} \\frac{{\\rm d}R(t)}{{\\rm d}t} \\approx \\kappa [R(t) - R_{\\rm cr}], \\label{rkdef} \\end{equation} where $R_{\\rm cr}$ is the radius of the critical or extremum bubble. This equation is valid when $R(t) \\approx R_{\\rm cr}$. In order for the bubble to grow, the initial radius $R(0)$ must be slightly larger than that of a critical bubble. In literature there are at least two calculations of the growth rate $\\kappa$ for relativistic plasma. Csernai and Kapusta write the free energy density as a functional of their order parameter which is just the usual internal energy density.\\cite{Csernai92} Ruggeri and Friedman make the approximation that the interface is infinitely thin and do not need any order parameter.\\cite{Ruggeri96} These two methods produce results for $\\kappa$ which disagree even qualitatively with each other. Let us now introduce a hydrodynamical model which enables both an analytical and a numerical determination of the initial growth rate $\\kappa$. The model~\\cite{Ignatius94} consists of an order parameter field $\\phi$, and perfect fluid which describes the other degrees of freedom. There are three basic locally varying quantities, namely $\\phi(x)$, fluid four-velocity $u^{\\mu}(x)$, and temperature $T(x)$. Due to the small value of the baryon asymmetry all the conserved particle numbers can for the present purposes well be approximated to be zero in the early Universe. Furthermore, the expansion of the Universe can be neglected, since the whole period of nucleation, yet alone the initial growth of bubbles, is an extremely rapid process. Thus the equations of motion can be written in the form \\begin{equation} \\left\\{ \\begin{array}{l} \\partial_{\\mu} T^{\\mu\\nu} = 0 \\\\ \\partial^2 \\phi + \\frac{\\partial}{\\partial\\phi} V(\\phi,T) = - \\eta \\, u^{\\mu} \\partial_{\\mu} \\phi, \\\\ \\end{array} \\right. \\label{EOM} \\end{equation} where $T^{\\mu\\nu} = T^{\\mu\\nu}\\left\\{u^{\\alpha}(x), \\partial_{\\beta}\\phi(x),\\phi(x),T(x)\\right\\}$ is the energy-momentum tensor, $V$ the potential energy density for the order parameter field, $T$ the temperature, and $\\eta$ a phenomenological friction parameter (not to be confused with shear viscosity). The upper equation is the conservation law of total energy-momentum. The lower equation tells how energy is transported between the order parameter and the fluid through the dissipative term, proportional to $\\eta$. The same term is also responsible for the creation of entropy. In the limit of vanishing fluid velocities the lower equation gives the simple dissipative equation \\begin{equation} \\frac{{\\rm d}\\phi}{{\\rm d}t} = - \\frac{1}{\\eta} \\frac{\\delta S_3[\\phi]}{\\delta \\phi}, \\label{dphidt} \\end{equation} where $S_3[\\phi]$ is the usual high-temperature three-dimensional action (equalling free energy at extrema). It is clear that a hydrodynamical description of the system cannot be complete. It is only valid at scales which are longer than particle mean free paths or interaction times. The solutions to Eqs.~(\\ref{EOM}) are smoothly behaving fields, whereas in reality fields have strong thermal fluctuations on short scales. To incorporate thermal fluctuations a Langevin-type equation with a noise term would be needed. Within a purely hydrodynamical model questions for example about damping effects in plasma cannot be answered. The model in Eqs.~(\\ref{EOM}) can be applied to describe both electroweak and QCD phase transition. In electroweak theory the order parameter is obviously identified as the Higgs field. The true coupling term is more complicated than the frictional $\\eta$-term of this model. But that effect should not be significant, as long as the detailed internal structure of the interface is not being discussed (something which a hydrodynamical model cannot accurately determine anyway). The value of the friction parameter $\\eta$ can be fixed by comparing with microscopic calculations.\\cite{Moore95} In the case of QCD the order parameter cannot be identified with a physical particle. However, one can still employ the model as a purely phenomenological construction. Interaction length or time of QCD sets the scale for the friction coefficient $\\eta$. A further complication is the macroscopic mean free path of neutrinos, but luckily the hydrodynamic energy flux is in normal cases clearly superior compared with that carried by the neutrinos. For solving Eqs.~(\\ref{EOM}) the potential $V(\\phi,T)$ must be known. Here the usual quartic potential has been employed. By fixing the parameters of it in a suitable manner, the desired values for the surface tension and latent heat of the transition will be reproduced. The coordinate system is spherically symmetric 1+3 dimensional space-time.\\cite{KurkiSuonio96} In order to create the initial configuration the critical bubble solution must be known with high accuracy. Going closer to the thin-wall limit, that is, using larger critical bubbles, has the advantage that numerical errors in the determination of growth coefficient $\\kappa$ decrease. But in this limit the field equation cannot be directly integrated numerically to produce the critical bubble. Instead, the following Ansatz is used: \\begin{equation} \\phi_A (r) = \\frac{\\phi_{\\rm min }}{2} \\left\\{ 1 - \\tanh \\left( \\frac{r-R_{\\rm cr}}{2\\xi_{\\rm cr}} \\right) \\right\\}. \\end{equation} Here $R_{\\rm cr}$, $\\xi_{\\rm cr}$ are {\\em unknown} parameters, and $\\phi_{\\rm min}$ is the position of the new minimum of the potential. In this two-dimensional subspace set by the Ansatz the extremum of the action becomes saddle point of an ordinary function, $S_3[\\phi_A] = S_{3A} (R_{\\rm cr}, \\xi_{\\rm cr})$. This saddle point can then be located numerically. For larger bubbles this method produces quite accurate results, which was actually unexpected. Fixing $R_{\\rm cr}$ naively by Laplace's relation, $R_{\\rm cr} = 2 \\sigma / \\Delta p$, where $\\sigma$ is surface tension and $\\Delta p$ pressure difference between the two phases, would lead to huge inaccuracies. However, the correct value of $R_{\\rm cr}$ can be found directly by employing curvature-dependent surface tension~\\cite{Ignatius93}, $\\sigma(R)$. Analytically the initial growth rate $\\kappa$ can be determined as follows. Expand the low-velocity dissipation equation (\\ref{dphidt}) around the critical bubble $\\bar{\\phi}$ by making the substitution $ \\phi(t,{\\bf x}) = \\bar{\\phi}(r) + \\varphi(t,{\\bf x})$. The resulting equation for fluctuations is \\begin{equation} \\frac{{\\rm d}\\varphi}{{\\rm d}t} = - \\frac{1}{\\eta} \\frac{ \\delta^2 S_3[\\bar{\\phi}] }{\\delta^2 \\phi} \\varphi. \\end{equation} Next insert the unstable growth mode, Eq.~(\\ref{deltaphi}). The thin-wall result for the negative eigenvalue of the fluctuation operator, $\\lambda_- = 2/R_{\\rm cr}^2$, has been observed to hold surprisingly well in the general case, too.\\cite{Brihaye93} This gives the approximate solution for the initial growth rate in the hydrodynamical model: \\begin{equation} \\kappa \\approx \\frac{2}{\\eta R_{\\rm cr}^2}. \\label{kresult} \\end{equation} Radius of the critical bubble, $R_{\\rm cr}$, depends on one hand on the cooling rate of the system---in cosmology on the strength of gravitational interaction---and on the other hand on the thermodynamical properties of the phase transition, especially on the values of latent heat and surface tension. \\setlength{\\unitlength}{0.7mm} \\begin{figure}[t] \\vspace*{-17mm} \\epsfig{file=figure1.ps,bbllx=20pt,bblly=275pt,bburx=612pt,bbury=792pt,% width=11cm,angle=0} \\caption{Numerical determination of the initial growth rate $\\kappa$. Horizontal axis is time in units of inverse thermodynamical phase transition temperature, $T_c^{-1}$, and vertical axis is $\\ln \\{ [R(t)-R_{\\rm cr} ] / [R(0)-R_{\\rm cr} ] \\} \\equiv y$. Initial growth rate is given by $\\kappa = \\lim_{t \\rightarrow 0} \\frac{{\\rm d}y}{{\\rm d}t}$ (when $R(0) \\rightarrow R_{\\rm cr}$). In the figure $\\eta=1 T_c$, $\\xi_{\\rm cr}=1.07243 T_c^{-1}$, $R_{\\rm cr}=19.695 T_c^{-1}$, and $R(0)/R_{\\rm cr}=1.01$. } \\label{fig:kt} \\end{figure} A more straightforward method is to integrate the dissipation equation~(\\ref{dphidt}) directly in the thin-wall limit.\\cite{Kajantie92} The result for the initial acceleration of the bubble radius is \\begin{equation} \\frac{{\\rm d}R}{{\\rm d}t} \\approx \\frac{2}{\\eta R_{\\rm cr}^2} \\left\\{ \\frac{d-1}{2}R - (d-2)R_{\\rm cr} \\right\\}, \\label{kd} \\end{equation} where the dimensionality of space, $d$, is explicitly visible. In the case $d=3$ Eq.~(\\ref{kresult}) follows from this by comparing with Eq.~(\\ref{rkdef}). Eq.~(\\ref{kd}) states clearly how in the real world the initial growth of bubbles is qualitatively different, much slower, than in one spatial dimension. The opposite approach is to let the bubble to expand in a hydrodynamical computer simulation, and to measure the initial growth rate numerically. Bubble radius is defined to be the distance where the tension or gradient energy has the maximum. The value of $\\kappa$ can be read from Fig.~\\ref{fig:kt} as the slope of the curve at origin. In the example case this gives $\\kappa = (0.0052 \\pm 0.0001)T_c$, which coincides well with the analytical estimate from Eq.~(\\ref{kresult}), $\\kappa \\approx 0.00516 T_c$. These results can be applied in the analysis of thermal cosmological phase transitions. In the case of relativistic heavy-ion collisions there is severe doubt on the validity of the general framework, nucleation in thermal systems. ", + "introduction": " ", + "conclusions": "" + }, + "9708/astro-ph9708163_arXiv.txt": { + "abstract": "Estimates of the total metal production rate ($\\dot{\\rho}_Z$) and star formation rate at $z > 2$ are based on Lyman break systems observed in the rest-frame ultraviolet (UV). These observations are very sensitive to dust obscuration. Here I elucidate and refine the Meurer \\etal\\ (1997) method for calculating UV obscuration, presenting new relationships which accurately model the dust reprocessing of radiation in local starbursts. The median $\\lambda = 1600$ \\AA\\ obscuration factor is $\\sim 10$ at $z \\approx 3$ which is shown to be consistent with other constraints on the high-$z$ $\\dot{\\rho}_Z$. Two tests are proposed to further constrain $\\dot{\\rho}_Z$ at these redshifts. \\endabstract ", + "introduction": "In their pioneering work, Madau \\etal\\ (1996; hereafter M96) use Lyman break galaxies in the Hubble deep field (HDF), in conjunction with other surveys, to evaluate the metal production rate $\\dot{\\rho}_Z$ history of the universe. They find that $\\dot{\\rho}_Z$ peaked at $z \\approx 1$, while at $z \\approx 3$ the universe was fairly quiescent much like the present universe. This result assumes dust absorption is negligible. However the high-$z$ observations, in the rest-frame UV, are highly susceptible to the obscuring effects of dust. Although it is now recognized that absorption corrections are significant, the amount of absorption is still under debate. Following from UV observations of local starbursts (Meurer \\etal\\ 1995, hereafter paper-1) we have argued that the the correction factor to $\\dot{\\rho}_Z$ is considerable (roughly 15; Meurer \\etal, 1997; hereafter paper-2), while others (e.g.\\ Madau 1997, hereafter M97; Pettini \\etal\\ 1997, hereafter P97) prefer fairly modest, $\\sim$ factor of three, corrections. As pointed out by M97, what is at stake is more than just the amount of dust in the early universe. If the absorption corrections to $\\dot{\\rho}_Z$ are low, then hierarchical galaxy formation models are favored, while large corrections at high-$z$ can push back the peak $\\dot{\\rho}_Z$, favoring monolithic collapse models. Here I take the opportunity to elucidate and refine our technique for estimating the UV attenuation due to dust, reapplying it to high-$z$ galaxies. Further constraints and refinements to the high-$z$ star formation rate are discussed. Finally, I present tests to further constrain the high-$z$ $\\dot{\\rho}_Z$. Throughout this paper I adopt $H_0 = 50\\, {\\rm km\\, s^{-1}\\, Mpc^{-1}}$ and $q_0 = 0.5$. ", + "conclusions": "" + }, + "9708/astro-ph9708213_arXiv.txt": { + "abstract": "We describe cosmological simulation techniques and their application to studies of cosmic structure formation, with particular attention to recent hydrodynamic simulations of structure in the high redshift universe. Collisionless N-body simulations with Gaussian initial conditions produce a pattern of sheets, filaments, tunnels, and voids that resembles the observed large scale galaxy distribution. Simulations that incorporate gas dynamics and dissipation form dense clumps of cold gas with sizes and masses similar to the luminous parts of galaxies. Models based on inflation and cold dark matter predict a healthy population of high redshift galaxies, including systems with star formation rates of $20 M_\\odot\\;{\\rm yr}^{-1}$ at $z=6$. At $z \\sim 3$, most of the baryons in these models reside in the low density intergalactic medium, which produces fluctuating \\lya absorption in the spectra of background quasars. The physical description of this ``\\lya forest'' is particularly simple if the absorption spectrum is viewed as a 1-dimensional map of a continuous medium instead of a collection of lines. The combination of superb observational data and robust numerical predictions makes the \\lya forest a promising tool for testing cosmological models. ", + "introduction": "The smoothness of the cosmic microwave background (CMB) tells us that the early universe was remarkably homogeneous, a cosmos without galaxies, stars, planets, or astronomers to admire them. Today's universe, on the other hand, exhibits structure over a vast range of scales. How did this transition take place? The leading hypothesis, strongly supported by COBE's discovery of anisotropies in the CMB (see Ned Wright's contribution to these proceedings), is that gravitational instability amplified tiny fluctuations present in the early universe into the rich structure that we observe today. This broad hypothesis leaves many more specific questions unanswered. What were the properties of the primordial fluctuations, and what was their physical origin? What is the dark matter? What are the values of the mass density parameter, $\\Omega$, and the cosmological constant (a.k.a.\\ vacuum energy density), $\\Lambda$? In principle, a set of answers to these questions, together with values of parameters like $H_0$ and the baryon density parameter $\\Omega_b$, constitutes a theory of structure formation. However, there is a large gap between a theory specified at this level and a set of observationally testable predictions. Numerical simulations that start from the theoretically specified initial conditions and evolve them forward in time play an essential role in bridging this gap. They show whether and how a theoretical model can produce objects like galaxies, quasars, and quasar absorption systems. They often provide greater understanding of observational phenomena, since in the simulations one knows the physical conditions, geometry, and history of the objects that form and the relation between the observable tracers of structure and the underlying distribution of dark matter. This understanding is specific to a theoretical model or class of theoretical models, but to the extent that the simulations reproduce the observations, one may be tempted to believe at least the general features of the physical picture that they provide, even if it is not correct in all of its details. By systematically exploring a range of models, one can see how changes in initial conditions, cosmological parameters, or physical assumptions reveal themselves in properties of observable structure. While cosmological simulations are often complex and computationally intensive, some aspects of the results may be simple to understand, and in such cases the simulations often suggest new analytic approximations, whose accuracy and range of validity can be tested against the numerical results. Simulations can also guide the development of new quantitative tools for characterizing observational data, providing physical motivation for the methods and showing what properties of the underlying structure are constrained by different measures. Finally, simulations yield quantitative predictions for these statistical measures. By comparing them to observational data, one can test the general physical picture of structure formation that emerges from the simulations, and one can distinguish between cosmological models that make different assumptions about primordial fluctuations, dark matter, and the values of $\\Omega$ and $\\Lambda$. In this paper we will describe several applications of cosmological simulations, focusing on our recent work using simulations with gas dynamics to investigate galaxy formation and the \\lya forest at $z \\sim 2-6$. ", + "conclusions": "" + }, + "9708/astro-ph9708025_arXiv.txt": { + "abstract": "The Fornax cluster galaxies NGC 1399 and NGC 1404 are ideal for studying the effects of a cluster environment on globular cluster systems. Here we present new optical imaging of these two galaxies from both the {\\it Hubble Space Telescope's} Wide Field and Planetary Camera 2 and the Cerro Tololo Inter--American Observatory's 1.5m telescope. The combination of both data sets provides unique insight on the spatial and colour distribution of globular clusters. From B--I colours, we find that both galaxies have a broad globular cluster metallicity distribution that is inconsistent with a single population. Two Gaussians provide a reasonable representation of the metallicity distribution in each galaxy. The metal--rich subpopulation is more centrally concentrated than the metal--poor one. We show that the radial metallicity gradient can be explained by the changing relative mix of the two globular cluster subpopulations. We derive globular cluster surface density profiles, and find that they are flatter (i.e. more extended) than the underlying starlight. The total number of globular clusters and specific frequency are calculated to be N = 5700 $\\pm$ 500, S$_N$ = 11.5 $\\pm$ 1.0 for NGC 1399 and N = 725 $\\pm$ 145, S$_N$ = 2.0 $\\pm$ 0.5 for NGC 1404. Our results are compared to the expectations of globular cluster formation scenarios. ", + "introduction": "Globular cluster (GC) systems are in some respects (e.g. luminosity function), more similar to one another than are their host galaxies. This suggests an underlying uniformity in the physical conditions under which GC systems and galaxies formed. On the other hand, properties such as the total number of clusters per unit galaxy luminosity show significant variations, and fundamental questions concerning the origins of these variations remain unanswered. In particular, the roles of tidal encounters, mergers, cooling flows, and initial conditions in the formation of GC systems have yet to be understood. Two galaxies of particular interest in this regard are NGC 1399 at the center of the Fornax cluster and a nearby cluster member NGC 1404. By studying the GC systems of these two galaxies we hope to better understand the environmental influences on the formation of GCs and their host galaxies. Various estimates of the distance modulus to the Fornax cluster have been made in recent years. Here we adopt a typical value of m--M = 31.2, which places the cluster 0.2 mag more distant than Virgo (Jacoby \\etal 1992). With this distance modulus and no Galactic extinction correction, the optical luminosity of NGC 1399 is M$_V$ = --21.74 (Faber \\etal 1989). This cD galaxy is surrounded by hot X--ray emitting gas which extends out to at least 38$^{'}$ or 190 kpc (Mason \\& Rosen 1985; Ikebe \\etal 1996). The mass-to-light ratio increases with galactocentric distance reaching a value of M/L $\\sim$ 70 M$_{\\odot}$/L$_{\\odot}$ at $\\sim$ 5$^{'}$ (Grillmair \\etal 1994a). Well within the X--ray envelope, even allowing for projection effects, at about 10$^{'}$ (50 kpc) to the south--east of NGC 1399 lies NGC 1404. It is classified as an E1 galaxy and has an optical luminosity of M$_V$ = --21.37 (Faber \\etal 1989). The GC systems of NGC 1399 and NGC 1404 have been studied using ground--based imaging by a number of workers (NGC 1399: Hanes \\& Harris 1986; Giesler \\& Forte 1990; Wagner \\etal 1991; Bridges \\etal 1991; Ostrov \\etal 1993; Kissler--Patig \\etal 1997a and NGC 1404: Hanes \\& Harris 1986; Richtler \\etal 1992). For NGC 1399, Kissler--Patig \\etal (1997a) gives the total number of GCs to be N$_{GC}$ = 5940 $\\pm$ 570. This gives a high specific frequency S$_N$ = 11 $\\pm$ 4. A kinematic study of the GCs around NGC 1399 showed that the outer GCs appear to be dynamically related more to the cluster of galaxies than they are to NGC 1399 itself (Grillmair \\etal 1994a). This suggests that accretion of intracluster GCs (West \\etal 1995), tidal stripping (Forte \\etal 1982; Forbes, Brodie \\& Grillmair 1997) or mergers (Ashman \\& Zepf 1992) may help to explain the high specific frequency of NGC 1399. For NGC 1404 the number of GCs is less certain. Hanes \\& Harris (1986) from a photographic study estimated N$_{GC}$ = 190 $\\pm$ 80, while more recently Richtler \\etal (1992) found N$_{GC}$ = 880 $\\pm$ 120. The resulting S$_N$ values range from 0.5 $\\pm$ 0.3 to 2.5 $\\pm$ 0.3 respectively. The latter value is close to the average for ellipticals in the Fornax cluster. Here we present imaging data from {\\it HST's} Wide Field and Planetary Camera 2 (WFPC2) and the Cerro Tololo Inter--American Observatory's (CTIO) 1.5m telescope of NGC 1399 and NGC 1404. The WFPC2 data provide accurate magnitudes and colours with virtually no contamination from foreground stars or background galaxies. These data are complemented by wide field-of-view CCD imaging from CTIO. In this paper we will focus on the colour (metallicity) and spatial distribution of GCs. The GC luminosity functions in these two galaxies and others in the Fornax cluster are addressed by Grillmair \\etal (1997), and a detailed study of NGC 1379, using WFPC2 and CTIO data, is given by Elson \\etal (1997). ", + "conclusions": "\\subsection{Spatial Distribution} Given the magnitude depth and the availability of a background field, our HST data are ideal to examine the surface density distribution of GCs. We first calculate the GC density in 5 annuli around NGC 1399. We make use of the fact that out to 100$^{''}$ radius the WFPC2 gives almost complete 180$^{\\circ}$ coverage for one hemisphere of the galaxy (see Forbes, Franx \\& Illingworth 1995 for details). For the NGC 1399 outer and background fields we simply add up the total number of GCs, i.e. 59 and 14 respectively. We have also made a small correction ($\\sim$ 15\\%) for the expected number of undetected GCs at the faint end of the luminosity function as determined by Grillmair \\etal (1997). Next, we divide the number of GCs by the appropriate spatial coverage to give a surface density. Finally, we subtract off the background density of 2.9 objects per square arcmin giving the background--corrected surface density of GCs. This is shown in Fig. 3. Excluding the innermost 2 data points for NGC 1399 we have fit the data with a function of the form: $\\rho = \\rho_{\\circ} r^{\\alpha}$ \\noindent This fit is shown by a solid line for which $\\rho_{\\circ}$ = 126 arcmin$^{-2}$ and $\\alpha$ = --1.2 $\\pm$ 0.2. It is clear from the NGC 1399 outer pointing that there are GCs at a projected radius of 8.2 arcmin from NGC 1399. Extrapolating the profile to 9.7 arcmin (the projected distance of NGC 1404) suggests that there is about one GC per square arcmin associated with NGC 1399. This corresponds to half a dozen GCs in the WFPC2 field-of-view centered on NGC 1404. The procedure for estimating the surface density profile of NGC 1404 is similar to that of NGC 1399, except that we subtract both the background density and the expected density of GCs in each annulus due to NGC 1399. Again the HST data is fit using a powerlaw profile (excluding the innermost data point) with $\\rho_{\\circ}$ = 36 arcmin$^{-2}$ and $\\alpha$ = --1.3 $\\pm$ 0.2. For both galaxies the GC surface density rises towards the galaxy center but flattens off (in log space) in the very inner regions giving a definable `core' to the GC system where the profile changes slope. This does not appear to be a selection effect and has been observed in other large galaxies (e.g. Grillmair \\etal 1986, 1994b). Forbes \\etal (1996) found that the `core radius' of the GC system is loosely correlated with the galaxy luminosity. We estimate that the GC systems of NGC 1399 and NGC 1404 have core radii of 40$^{''}$ (3.4 kpc) and 30$^{''}$ (2.5 kpc) respectively. These values are within the scatter of the relation found by Forbes \\etal (1996) and are similar to the galaxy effective radii. We have measured a GC surface density slope of --1.2 $\\pm$ 0.2 for NGC 1399. There have been three previous CCD studies which measured the GC slope. Wagner \\etal (1991) found that interior to $\\sim$ 120$^{''}$ the GC slope was --1.4 $\\pm$ 0.11, and this steepened to --1.54 $\\pm$ 0.15 for larger radii. Kissler--Patig \\etal (1997a) derived --1.55 $\\pm$ 0.25 and --1.75 $\\pm$ 0.3 from two pointings for radii beyond $\\sim$ 50$^{''}$. Bridges \\etal (1991) measured --1.4 $\\pm$ 0.2 from their B band data and --1.5 $\\pm$ 0.2 from their V band data. Our value is somewhat flatter than those measured by others. One possible reason for this is that our outer pointing is located in the cD envelope of NGC 1399 where the GC surface density `flattens out' (e.g Kissler--Patig \\etal 1997a). If we exclude the outer point, our fitted slope steepens to --1.25. Another possibility is that, even though we have not included the inner two data points in the fit, we are still fitting into the core--flattened region. The stellar profile (log intensity) for NGC 1399 is also shown in Fig. 3. We measure a slope of --1.6 $\\pm$ 0.1 for the galaxy light. This compares with the measurements of --1.67 $\\pm$ 0.12 (Wagner \\etal 1991) and --1.75 $\\pm$ 0.1 (Kissler--Patig \\etal 1997a). There are been some debate as to whether the GC profile is flatter than the galaxy profile in NGC 1399. In particular, Bridges \\etal (1991) claimed that the GC system had the same slope as the galaxy. Whereas both Wagner \\etal (1991) and Kissler--Patig \\etal (1997a) said it was somewhat flatter. Excluding our data, the weighted mean value from previous studies is --1.47 $\\pm$ 0.09 for the GC slope and --1.72 $\\pm$ 0.06 for the galaxy slope. This indicates that the GC profile is 0.25 $\\pm$ 0.11 flatter than the galaxy starlight. Our data suggest that the difference is even greater, i.e. 0.4 $\\pm$ 0.22. Thus, we conclude that the GC surface density profile is flatter than the galaxy starlight by about 0.3 in the log at about the 2 $\\sigma$ significance level. For NGC 1404, we measure a GC slope of --1.3 $\\pm$ 0.2 and a galaxy starlight slope of --1.9 $\\pm$ 0.1. Thus the GC systems in both galaxies have a flatter, more extended distribution that the underlying starlight. Using the surface density profile we can estimate the total number of GCs and the specific frequency (S$_N$; Harris \\& van den Bergh 1981) for NGC 1399 and NGC 1404. For NGC 1399 GCs are clearly extended beyond 8 arcmin in our data. We have decided to use 10 arcmin as the limiting radius, which is the same limit used by Hanes \\& Harris (1986) in their photographic study. Integrating the powerlaw profile between log r = --0.2 (38$^{''}$) and 1 (10 arcmin), and adding the number of GCs interior to 38$^{''}$ gives a total of 5700 GCs. The uncertainty in deriving the total number of GCs is dominated by the choice of limiting radius. We have used $\\pm$1 arcminute about the limiting radius as our error estimate. Thus limiting radii of 9 and 11 arcminutes give N = 5700 $\\pm$ 500. This is in good agreement with the determination of Kissler--Patig \\etal (1997a) of N = 5940 $\\pm$ 570. For an absolute magnitude of M$_V$ = --21.74, the specific frequency S$_N$ = 11.5 $\\pm$ 1.0. For NGC 1404 the limiting radius is difficult to estimate. Richtler \\etal (1992) determined a background level at $\\sim$ 200$^{''}$. We find that in the direction of NGC 1399, the density of GCs is clearly dominated by NGC 1399 beyond about 250$^{''}$. We have chosen to integrate out to 240$^{''}$ (4 arcmin) with a reasonable range of 3--5 arcminutes. Thus integrating the profile from log r = --0.3 (30$^{''}$) to 0.6 (4 arcmin) and adding the number of inner GCs gives N = 725 $\\pm$ 145 GCs associated with NGC 1404. This lies between the Hanes \\& Harris (1986) value (190 $\\pm$ 80) and that of Richtler \\etal (1992), i.e. 880 $\\pm$ 120. Compared to these ground--based studies, our HST data has the advantages of accurate background subtraction (star, galaxies and NGC 1399 GCs) and nearly 100\\% complete magnitude coverage. For an absolute magnitude of M$_V$ = --21.37, our results indicate a relatively low specific frequency for NGC 1404 of S$_N$ = 2.0 $\\pm$ 0.5. Again using the HST data, we calculate the `local' S$_N$ values (i.e. using the number of GCs and the integrated magnitude at the de Vaucouleurs effective radius) to be $\\sim$ 1.5 for NGC 1399 and $\\sim$ 0.5 for NGC 1404. In Figures 4 and 5 we show the azimuthal distribution of GCs, as determined from the HST data, over the range for which we have uniform radial and azimuthal coverage. For NGC 1399 (Fig. 4) slight enhancements in the GC counts can be seen close to the galaxy major axes (position angles $\\sim$ 100$^{\\circ}$ and --80$^{\\circ}$). There is also a slight deficit around the minor axis P.A. $\\sim$ 10$^{\\circ}$. Thus the GC system is broadly aligned with the stellar isophotes. The direction towards NGC 1404 (P.A. $\\sim$ 150$^{\\circ}$) is not covered by our HST image. For NGC 1404, shown in Fig. 5, there is slight deficit in GC counts along the galaxy major axis lies at P.A. $\\sim$ --20$^{\\circ}$. There are also two enhancements at position angles $\\sim$ --50$^{\\circ}$ and 10$^{\\circ}$ which do not correspond to either galaxy major or minor axes. \\subsection{Colour and Metallicity Distributions} The GC colour distributions, after colour and magnitude selection, for the four HST pointings are shown in Fig. 6. Both NGC 1399 and NGC 1404 are dominated by red (B--I $\\sim$ 2.1) GCs but with a significant tail to the blue (e.g. B--I $\\sim$ 1.6). The GCs in the outer pointing of NGC 1399 are bluer on average than those in the central pointing. The median colours are B--I = 1.97 for NGC 1399 in the central pointing, 1.70 for the outer NGC 1399 pointing and 1.96 for NGC 1404. The CTIO data, after colour and magnitude selection, are shown in Fig. 7. We do not have a background field for the CTIO data. Although we have attempted to remove galaxies using the DAOPHOT sharpness and roundness parameters, our seeing conditions of $\\sim$ 1.5$^{''}$ imply that some galaxies will be included in our candidate list. The issue of background galaxy contamination in GC colour/metallicity distributions is discussed in some detail by Elson \\etal (1997). Background galaxies in our magnitude range peak at B--I $\\sim$ 1.0 and get bluer at fainter magnitudes. The galaxy counts are falling fast at our blue cutoff of B--I = 1.2. Given the richness of the NGC 1399 GC system, background galaxies are likely to be a small affect. However it may be more of a concern for NGC 1404. We can estimate the effects of background contamination using the CFRS sample (Lilly \\etal 1995) within the appropriate B magnitude range and scaled in area to match that of the CTIO data. The background--corrected distributions are shown by a dashed line in Fig. 7. As we have already removed some galaxies in our initial detection process, the true distribution may lie somewhere between the two histograms. In either case (uncorrected or corrected) the distributions reveal a somewhat different situation to the HST data, with both galaxy's GC systems peaking in the blue (B--I $\\sim$ 1.6) with a significant red tail (around B--I $\\sim$ 2.1). The median colours are B--I = 1.84 for NGC 1399 and 1.71 for NGC 1404. Thus the CTIO colours are on average systematically $\\sim$ 0.1 mag bluer than the HST data. This result is not formally significant as the error is also on the order of 0.1 mag, but differences between the shape of the two distributions are clearly seen. As the HST data cover only the central regions of each galaxy, this suggests a radial colour gradient from red dominated objects at small galactocentric radii to blue ones further out. In principle, such a gradient could be caused by blue background galaxies in the CTIO data. However we give further evidence below that this is not the case. In Fig. 8 we show the GC colour distributions for NGC 1399 and NGC 1404 from the HST data. Here the data for NGC 1399 includes both the central and outer pointing. The background--corrected histograms are shown by dashed lines. Some galaxies, such as M87 (Whitmore \\etal 1995; Elson \\& Santiago 1996) reveal a clear bimodal GC colour distribution. In other cases (i.e. NGC 1374, 1379, 1387, 1427) the distribution is essentially consistent with a single unimodal colour for the GC system (Kissler--Patig \\etal 1997a,b). The situation for NGC 1399 and NGC 1404 is not as clear. Nevertheless there is some evidence for multiple GC populations. Firstly, both galaxies have distributions that are significantly broader than that expected from the photometric errors (typically $\\sigma$ $\\sim$ 0.15), i.e. there is a real spread in GC colours. Secondly, we have tested the unbinned colour data using the KMM statistical test (Ashman, Bird \\& Zepf 1994). This test rejects a single Gaussian fit to the NGC 1399 and NGC 1404 data with a confidence of over 99\\%. If we represent the data with two Gaussians then KMM determines means of B--I = 1.7 and 2.1 for NGC 1399, and B--I = 1.6 and 2.1 for NGC 1404. Thirdly, as an exercise, we generated a colour histogram that is the sum of two Gaussians. These Gaussians have mean colours of B--I = 1.6 and 2.1 with dispersions similar to the photometric error, and with total numbers in the ratio of 1:4. They are shown in the lower left panel of Fig. 8. In the lower right panel, we show the same two Gaussians but with alternating Poisson noise included. To the eye, the two noisy Gaussians are qualitatively similar to the HST data, suggesting they could consist of two Gaussian--like populations. Fourth, the two peaks indicated in the HST data are also present in the CTIO data (Fig. 7). {\\it We conclude that both galaxies do not have a single, uniform GC population but rather show evidence for a multimodal GC colour distribution. } Two other large data sets exist for GCs in NGC 1399. They are the CTIO 4m observations by Ostrov \\etal (1993) using Washington photometry and the Las Campanas 2.5m observations by Kissler--Patig \\etal (1997a) in V and I. In order to compare all of the different data sets we have converted each into metallicity. Broad--band colours can be transformed into metallicity using the Galactic relation of Couture \\etal (1990) given the usual caveat that age effects and abundance anomalies may be present. The Washington photometry of Ostrov \\etal is transformed using the relation derived by Geisler \\& Forte (1990). Washington photometry is the most sensitive to metallicity (rms $\\sim$ 0.1), then B--I colours (rms $\\sim$ 0.15) with V--I colours being the least sensitive (rms $\\sim$ 0.25). In Fig. 9 we show the GC metallicity distributions from our HST and CTIO data sets, along with those of Ostrov \\etal and Kissler--Patig \\etal None of the data sets have been background corrected in this figure (although for the HST sample the correction is negligible). In general, the samples reveal a broad distribution with several minor enhancements, many of which appear common to several data sets. From their data, Ostrov \\etal claimed that NGC 1399 has a trimodal GC metallicity distribution, with peaks at [Fe/H] $\\sim$ --1.5, --0.8, --0.2. The metal--rich peak corresponds to that seen clearly in our HST data. The two metal--poor peaks may be real or simply a single population with [Fe/H] $\\sim$ --1.0. The mean B--I colours found by the KMM test on the HST data correspond to [Fe/H] $\\sim$ --1.1 and --0.1. The GC metallicity distribution for NGC 1404 from our HST data is shown in Fig. 10. As indicated by the KMM test, a bimodal distribution is a better representation of the data than a single Gaussian. The two GC populations have means of [Fe/H] $\\sim$ --1.5 and --0.1. We now return to the CTIO data and the issue of background contamination. Earlier we argued that contamination, although present, did not change the basic appearance of the colour distributions. Further evidence that the two peaks in the CTIO data are dominated by {\\it bona fide} GCs comes from examining the spatial distribution of the two subpopulations. First, we define the metal--rich and metal--poor GC subpopulations in each galaxy as being $\\pm$ 0.3 dex about the mean metallicity of each subpopulation. We then calculate the raw surface density (i.e. not corrected for background contamination or for the missing faint end of the luminosity function) for the two subpopulations. These surface density profiles are shown in Fig. 11. A background correction would tend to lower the metal--poor profile, particularly at galactocentric radii beyond 4 arcmin in NGC 1404 (which is excluded from the figure). Although there is considerable scatter, the surface density of each subpopulation declines with distance from its parent galaxy. This indicates that each subpopulation is associated with the galaxy and is therefore dominated by GCs. Secondly, the fact that the HST colours (metallicities) agree with those in the inner annuli of the CTIO data (see figures 12 and 13) indicates that the subpopulations are dominated by GCs. Next we investigate the radial variation of GC metallicity. We have seen that for both NGC 1399 and NGC 1404, the metal--rich (red) subpopulation dominates in the HST data and in the CTIO data the metal--poor (blue) subpopulation is dominant. As the HST data probe the inner $\\sim$ 100$^{''}$ region and the CTIO data cover the regions beyond $\\sim$ 30$^{''}$ this would suggest that the metal--rich GCs are more centrally concentrated than the metal--poor ones. A similar situation is inferred for NGC 4472 (Geisler \\etal 1996) and NGC 5846 (Forbes, Brodie \\& Huchra 1997). An unweighted fit to the HST data indeed indicates a metallicity gradient in both galaxies. For NGC 1399 we find [Fe/H] = --0.31$\\pm$0.13 log R (arcsec) -- 0.046$\\pm$0.22. This slope is similar to that found by Ostrov \\etal (1993) i.e. --0.34$\\pm$0.18. The fit for NGC 1404 gives [Fe/H] = --0.36$\\pm$0.13 log R (arcsec) -- 0.034$\\pm$0.22. In both cases the HST fits are consistent with the CTIO data. In NGC 4472, Geisler \\etal (1996) found that the radial GC metallicity gradient was actually due to the changing relative mix of the two GC subpopulations. Assuming that NGC 1399 and NGC 1404 can be represented by metal--rich and metal--poor GC subpopulations we investigate the radial gradients further. In Figures 12 and 13 we show the metallicity in various radial bins for the different GC subpopulations and the mean metallicity for the combined population. The range in metallicity included in each data point is $\\pm$ 0.3 dex. For NGC 1399 (Fig. 12) the metal--poor and metal--rich GC subpopulations show no obvious radial metallicity gradient. However the mean values for the whole GC system {\\it do} show a global radial gradient. Although not quite as convincing, the same trend appears to be true for NGC 1404 (Fig. 13). For both galaxies, we conclude that the radial metallicity gradient for the overall GC system is consistent with the changing relative proportions of the GC subpopulations; the metal--rich GCs are centrally concentrated while the metal--poor ones are preferentially located in the outer regions. The lack of a real abundance gradient places constraints on the role of gas dissipation in the formation of GCs. In NGC 4472, the metallicity of the galaxy field stars was found to be of the same or of slightly higher metallicity than the metal--rich GCs (Geisler \\etal 1996). For NGC 1399, Brodie \\& Huchra (1991) quote a spectroscopic metallicity of [Fe/H] = +0.2 $\\pm$ 0.9. This is similar to the metallicity derived from B--I colours of the galaxy (from our CTIO data and from Goudfrooij \\etal 1994). For NGC 1404, the galaxy B--I colours indicate [Fe/H] $\\sim$ --0.2. Thus for both galaxies, the field stars appear to have similar metallicities to those of metal--rich GCs (which have [Fe/H] $\\sim$ --0.1). \\subsection{Comparison with Globular Cluster Formation Scenarios} \\subsubsection{The Merger Model} The merger model of Ashman \\& Zepf (1992) and Zepf \\& Ashman (1993) make some specific predictions concerning the properties of GCs in ellipticals that are the result of a gaseous merger. One of their key predictions is that ellipticals will have multimodal GC metallicity distributions. Furthermore the metal--rich GCs will be centrally concentrated and the more metal--poor ones will be preferentially located at large galactocentric radii. This indeed appears to be the case for both NGC 1399 and NGC 1404, and is therefore in qualitative agreement with their model. For NGC 1404 with a low S$_N$ value (i.e. 2.0 $\\pm$ 0.5), Ashman \\& Zepf (1992) expect slightly fewer metal--rich (new) than metal--poor (old) GCs. Our data are generally consistent with this expectation. For NGC 1399 with S$_N$ = 11.5 $\\pm$ 1.0, they would expect N$_R$/N$_P$ $\\sim$ 3--4. In the central regions of NGC 1399 such a ratio may hold, but globally N$_R$/N$_P$ $<$ 1. This is of course complicated by the difficulty in defining clear metal--rich and metal--poor GC subpopulations. However there seems to be a shortfall in the number of newly created GCs if gaseous mergers are responsible for the metal--rich GCs and the high S$_N$ value. Another expectation from their merger model is that the galaxy starlight and GC system should have a similar spatial distribution in galaxies with high S$_N$ values (i.e. those galaxies in which large numbers of new GCs should have been created by the merger). Thus we might expect the surface density of GCs to match that of the galaxy surface brightness profile for NGC 1399 but not for NGC 1404. Our data indicate that for both galaxies the GC distribution is notably flatter than the underlying starlight. We might also expect the metal--rich GCs, which should have formed from the gas of the progenitor galaxies, to reveal a radial metallicity gradient as the formation process should be dissipative. Figures 12 and 13 suggest that there is little or no metallicity gradient for the metal--rich GCs in either galaxy. To summarize, the GC systems in NGC 1399 and NGC 1404 reveal some properties that are consistent with the Ashman \\& Zepf (1992) merger model, but there are also notable conflicts. On an individual basis, some of the disagreements with the model may be accommodated by slight modifications to the initial model (see Zepf, Ashman \\& Geisler 1995). Forbes, Brodie \\& Grillmair (1997) have recently investigated whether a sample of elliptical galaxies meets the general predictions of the Ashman \\& Zepf (1992) model, and concluded that mergers were {\\it unlikely} to account for the GC systems in large elliptical galaxies. They favored a multiphase collapse scenario. Next we compare the data for NGC 1399 and NGC 1404 with this scenario. \\subsubsection{The Multiphase Collapse Model} Forbes, Brodie \\& Grillmair (1997) have proposed that GCs in ellipticals form {\\it in situ} during a multiphase collapse. In the first pre--galactic phase, metal--poor GCs are formed and later on, in the second galactic phase, the metal--rich GCs form. This leads to a bimodal metallicity distribution with the metal--rich GCs more centrally concentrated than the metal--poor ones. In massive ellipticals the collapse could be largely dissipationless and so no strong radial abundance gradients are expected. The radial distribution of (the metal--poor) GCs will be more extended than the galaxy itself, if as expected, they form before the galactic stars. The overall metallicity and spatial properties of GCs derived from our data are consistent with this picture. For cD galaxies, such as NGC 1399, Forbes, Brodie \\& Grillmair (1997) suggested an additional source of GCs from tidal stripping of nearby galaxies. They suggested that some of GC subpopulation identified by Ostrov \\etal (1993) at around [Fe/H] $\\sim$ --0.8 were acquired from NGC 1404 and possibly other Fornax galaxies. Recently, Jone \\etal (1997) have interpreted the asymmetric X--ray emission around NGC 1404 as evidence for tidal interaction with NGC 1399. The idea of tidal stripping of GCs is given some support by our new data. The outer regions of NGC 1404 are dominated by metal--poor GCs, and it is these GCs that may be preferentially stripped and captured by NGC 1399. Thus some of the GCs in NGC 1399 with this intermediate metallicity may have originated from NGC 1404. We confirmed that NGC 1404 has a remarkably low specific frequency value (S$_N$ = 2.0 $\\pm$ 0.5) for a cluster elliptical, suggesting that some GCs may have been lost. If NGC 1404 originally had S$_N$ = 5, i.e. typical of normal ellipticals outside of the Fornax cluster then it has lost $\\sim$ 1000 GCs, which may have been tidally captured by NGC 1399. If we assume that the GC system and galaxy starlight originally (i.e. before tidal stripping) continued beyond 4 arcminutes with the same slope as found in section 3.1, then we can calculate the S$_N$ value of the stripped GCs. We find that the stripped GCs are unlikely reach values of S$_N$ $\\sim$ 11.5 in sufficient numbers to have any appreciable effect on the specific frequency of NGC 1399. We expected the GC surface density profile to match that of the underlying starlight in the inner regions of NGC 1404, but it appears to be somewhat flatter than the stellar profile (at least for galactocentric radii of less than $\\sim$ 150$^{''}$). This may just be indicating that tidal stripping has not yet affected the galaxy inner regions. The above circumstantial evidence suggests that NGC 1399 has acquired GCs from NGC 1404, although this is unlikely to explain the high specific frequency of NGC 1399. Large samples of GC metallicity and kinematic information from spectra are probably required before the multiphase collapse and tidal stripping ideas can be fully tested." + }, + "9708/astro-ph9708096_arXiv.txt": { + "abstract": "The outer-halo globular cluster NGC 6229 has a peculiar horizontal-branch (HB) morphology, with clear indications of a bimodal HB and a ``gap\" on the blue HB. In this paper, we present extensive synthetic HB simulations to determine whether peculiar distributions in the underlying physical parameters are needed to explain the observed HB morphology. We find that a unimodal mass distribution along the HB can satisfactorily account for the observed HB bimodality, {\\em provided} the mass dispersion is substantially larger than usually inferred for the Galactic globular clusters. In this case, NGC 6229 should have a well-populated, extended blue tail. A truly bimodal distribution in HB masses can also satisfactorily account for the observed HB morphology, although in this case the existence of an extended blue tail is not necessarily implied. The other two well-known bimodal-HB clusters, NGC 1851 and NGC 2808, are briefly analyzed. While the HB morphology of NGC 1851 can also be reproduced with a unimodal mass distribution assuming a large mass dispersion, the same is not true of NGC 2808, for which a bimodal, and possibly multimodal, mass distribution seems definitely required. The problem of gaps on the blue HB is also discussed. Applying the standard Hawarden (1971) and Newell (1973) $\\chi^2$ test, we find that the NGC 6229 gap is significant at the $99.7 \\%$ level. However, in a set of 1,000 simulations, blue-HB gaps comparable to the observed one are present in $\\sim 6 \\% - 9 \\%$ of all cases. We employ a new and simple formalism, based on the binomial distribution, to explain the origin of this discrepancy, and conclude that Hawarden's method, in general, substantially overestimates the statistical significance of gaps. ", + "introduction": "The outer-halo globular cluster (GC) NGC 6229 (C1645+476) has recently been studied photometrically by Borissova et al. (1997, hereafter Paper I). These authors showed that the horizontal-branch (HB) morphology of this cluster is characterized by two peculiar properties: i) A distinctly bimodal distribution in $\\bv$ colors, consistent with the relatively small number of RR Lyrae variables, compared to the blue HB and red HB populations; and ii) At least one ``gap\" on the blue HB, as previously noted by Carney, Fullton, \\& Trammell (1991). Bimodal-HB clusters are relatively rare, as are clusters with gaps on the blue HB. NGC 6229 is one of only two known cases---the other being NGC 2808 (Sosin et al. 1997a)---where these two anomalies are simultaneously present. As argued by many authors, an understanding of the nature of the detected peculiarities in ``bimodal\" and ``gap\" clusters would be of paramount importance for understanding the nature of the second parameter phenomenon (e.g., Buonanno, Corsi, \\& Fusi Pecci 1985; Rood et al. 1993; Stetson, VandenBerg, \\& Bolte 1996; Paper I). These considerations have prompted us to perform a theoretical investigation of the problem of HB bimodality and gaps. In particular, we have considered the following questions: i) Can HB bimodality be produced by canonical evolutionary models? ii) If so, what are the requirements for producing HB bimodality in the specific cases of NGC 6229, NGC 1851, and NGC 2808? iii) Are gaps on the blue HB a natural consequence of stellar evolution, as claimed by some authors? iv) What is the statistical significance of gaps? In this paper, we present extensive synthetic HB simulations for NGC 6229. We begin in the next section by laying out a proposed nomenclature for both defining and differentiating ``bimodal\" and ``gap\" clusters. In Sec. 3, we describe the theoretical procedure used in Sec. 4 to compute synthetic HBs for NGC 6229. The statistical significance of gaps on the blue HB is next discussed in Sec. 5. In Sec. 6, we present additional computations for the bimodal-HB clusters NGC 1851 and NGC 2808. In Sec. 7, we discuss some previously proposed scenarios for the origin of HB bimodality and gaps, in the light of the present results. Finally, Sec. 8 summarizes our results. In an Appendix, a compilation is presented and discussed of GCs which may present signs of HB bimodality and gaps, according to the definitions laid out in Sect. 2. Several of these clusters may require more detailed observational and theoretical investigations in the future. ", + "conclusions": "With the insight provided by our HB simulations, we shall now analyze some previously proposed scenarios for the origin of HB bimodality and gaps. For a discussion of non-standard scenarios, we refer the reader to the recent papers by van den Bergh (1996), Catelan (1997), Sosin et al. (1997a), and Sweigart (1997b). \\subsection{HB Bimodality} It has sometimes been suggested that bimodal HBs may be accounted for by unimodal mass distributions along the HB (Norris 1981; Lee et al. 1988; Walker 1992). In particular, Walker has argued that, given a sufficiently wide dispersion in mass along the HB, bimodality is naturally implied. But does this apply equally well to both $\\bv$ colors and temperatures? In Figs. 15a and 15b we show the ($\\log\\,L$, $\\log\\,T_{\\rm eff}$) and ($M_V$, $\\bv$) planes, respectively, for a ``reference\" synthetic HB populated by a large number (100,000) of HB stars randomly sampled from a uniform distribution in mass along the ZAHB over the range $0.495 \\leq M/M_{\\sun} \\leq 0.820$. This large number of HB stars was used to minimize any statistical fluctuations. Fig. 15b shows that, on the observational diagram, the number of stars at intermediate $\\bv$ colors (where the RR Lyrae variables are found) seems to be significantly depleted in comparison with the number of stars both on the blue HB and on the red HB. In Fig. 16, we show the color-temperature relation from the Kurucz (1992) model atmospheres for a representative gravity and metallicity. These diagrams show very clearly the well-known fact (e.g., Rood \\& Crocker 1989) that bimodal color distributions may be produced due to the piling up of HB evolutionary tracks at the low-temperature (i.e., high-mass) end, and to the insensitivity of the $\\bv$ color to temperature changes for $\\log\\,T_{\\rm eff} \\gtrsim 3.9$. Thus, HB bimodality---at least in the observational sense previously defined---is naturally implied for sufficiently wide mass distributions, as we had indeed found on the basis of detailed simulations. However, inspection of Fig. 15a reveals that, for a uniform mass distribution, a rather uniform distribution in temperatures results. Interestingly, Walker (1992) has proposed that the morphology of the HB evolutionary tracks---in particular, the presence of long ``blue loops\" in the post-ZAHB evolution---might produce a bimodal distribution in temperature (and thus also in the number counts and $\\bv$ colors) along the HB. The only requirement, according to Walker, is that the mass dispersion be substantially larger than typically assumed in HB simulations. Can track morphology really lead to two main statistical modes in the HB temperature distribution? To address this question, we show in Fig. 17 a synthetic HB employing 20,000 stars randomly sampled from a uniform mass distribution over a mass range chosen to cover the temperature range in Walker's Fig. 8 for NGC 1851. Our assumption of a flat mass distribution favors Walker's suggestion by adding more blue HB stars than in the detailed simulations for this cluster shown in Fig. 13. The same chemical composition as in Walker's study was also adopted. Two histograms for $T_{\\rm eff}$ are given in Fig. 17: one which includes the canonical evolution of the HB stars away from the ZAHB, and another where all evolution away from the ZAHB is suppressed. As one can clearly see, both histograms are comprised of a single statistical mode on the red clump region, accompanied by a quite uniformly populated distribution extending to higher temperatures. There is no sign of a peak in the temperature distribution blueward of the instability strip. These histograms clearly show that evolution off the ZAHB does not help in producing two main modes in the temperature distribution. We have carried out a similar simulation using the HB evolutionary tracks of Lee \\& Demarque (1990), as shown in Fig. 18. As remarked by Catelan \\& de Freitas Pacheco (1993), the Lee \\& Demarque HB evolutionary tracks, for some still unclear reason, present blueward loops which are much longer in $\\Delta\\log\\,T_{\\rm eff}$ (for the same chemical composition) than in other independent HB evolutionary calculations. This difference in track morphology dramatically affects the synthetic HBs, particularly the luminosity width, as can clearly be seen by comparing Figs. 17 and 18. Interestingly, the more recent Koopmann et al. (1994) and Yi, Lee, \\& Demarque (1995) HB tracks show blueward loops which are in much better agreement with those presented by, e.g., Sweigart (1987), as well as with those employed in the present work. The differences in track morphology notwithstanding, this simulation also shows a remarkably flat temperature distribution blueward of the instability strip. (The slight irregularities in this HB simulation are due to the coarseness of the Lee \\& Demarque grid.) The histograms once again demonstrate that the HB evolution does not contribute to the production of two main statistical modes in the temperature distribution. Our results do not therefore support the argument (Lee et al. 1988) that bimodality can arise naturally from ``evolution away from the ZAHB.\" There is, however, one circumstance under which the blueward loops may be {\\em directly} responsible for producing HB bimodality. For some combinations of evolutionary parameters---for instance, for large values of the envelope helium abundance, or for low values of the helium-core mass $M_{\\rm c}$---blueward loops may become {\\em very} long indeed (see, e.g., Sweigart \\& Gross 1976). Primordial (Shi 1995) or evolutionary (Sweigart 1997a, 1997b) processes might be responsible for a large helium abundance in HB stars. A substantial decrease in $M_{\\rm c}$ would, however, be harder to justify (Sweigart 1994; Catelan, de Freitas Pacheco, \\& Horvath 1996). In either case, for sufficiently large $Y$ or low $M_{\\rm c}$ values ZAHB stars on the red HB clump will evolve along extremely long blueward loops and thereby produce a substantial number of blue HB stars, and possibly even a relatively long blue tail. Moreover, provided the ZAHB mass distribution does not populate the region between the red clump and the blue HB, few stars would be present near the instability strip, because of the rapid evolutionary pace at these intermediate temperatures. If such a scenario were to be responsible for a bimodal HB, however, we would expect the HB distribution to be {\\em sloped} on the CMD (Catelan \\& de Freitas Pacheco 1996). Interestingly, bimodal {\\em and sloped} HBs have recently been reported for the metal-rich GCs NGC 6388 and NGC 6441 (Piotto et al. 1997; Rich et al. 1997). We have carried out extensive simulations to explore this scenario in detail and have confimed that high $Y$ values can lead to both bimodal and sloped HBs. Preliminary results for these simulations have been presented by Sweigart \\& Catelan (1997a) and will be discussed more fully in a future paper (Sweigart \\& Catelan 1997b). \\subsection{Gaps on the Blue HB} The change in the morphology of canonical post-ZAHB evolutionary tracks as a function of mass has sometimes been invoked to explain the presence of blue-HB gaps (e.g., Norris et al. 1981; Lee et al. 1994). In particular, it has been argued that the NGC 288 gap, and possibly the NGC 6752 one, may be easily explained in this way (Lee et al. 1994). However, numerical experiments by Castellani et al. (1995) failed to provide conclusive evidence for or against this suggestion. Inspection of Figs. 15a and 15b shows no significantly underpopulated regions on the blue HB in either diagram. Similar plots using the Lee \\& Demarque (1990) HB tracks do not show any significantly underpopulated regions either. This indicates that, contrary to previous claims, {\\em gaps on the blue HB are not caused by the morphology of canonical HB tracks}. One noteworthy feature in Fig. 15b is the non-monotonic behavior of the number counts as a function of \\bv at $\\bv \\approx 0.2 - 0.3$. No corresponding feature seems to be present in the theoretical plane (cf. Figs. 15a and 17). Inspection of the Kurucz (1992) color transformations (cf. Fig. 16) discloses that this anomaly is related to an abrupt change in slope of the $(\\bv ) - \\log\\,T_{\\rm eff}$ relation at this region. We refer the reader to the papers by Gratton, Carretta, \\& Castelli (1996, esp. their Sec. 5) and Lejeune, Cuisinier, \\& Buser (1997) for discussions of this feature, which is not present in their more recent color-temperature transformations. Some experiments with the non-standard evolutionary tracks of Sweigart (1997a, 1997b) show that track morphology may play a r\\^ole in generating gaps, but only if some physical parameter(s) (such as $Y$ or $M_{\\rm c}$) changes {\\em abruptly} as a function of temperature along the ZAHB, thus effectively generating a {\\em break} in track morphology at a given temperature, as opposed to the continuous and smooth change that is found in the canonical case (Fig. 15). A gap on the extreme HB is also predicted by the scenario outlined by D'Cruz et al. (1996). In their framework, very hot extreme HB stars ($\\log\\,T_{\\rm eff} \\approx 4.5$; cf. their Fig. 2) are produced by the stars which undergo the He flash on the white dwarf cooling curve as a consequence of high mass loss rates on the RGB (Castellani \\& Castellani 1993). In this case, however, the predicted gap location would be even hotter than the gap in NGC 6752 (Sweigart 1997b)." + }, + "9708/astro-ph9708119_arXiv.txt": { + "abstract": "Properties of high redshift clusters are a fundamental source of information for cosmology. It has been shown by Oukbir and Blanchard (1997) that the combined knowledge of the redshift distribution of X-ray clusters of galaxies and the luminosity-temperature correlation, $L_X-T_X$, provides a powerful test of the mean density of the Universe. In this paper, we address the question of the possible evolution of this relation from an observational point of view and its cosmological significance. We introduce a new indicator in order to measure the evolution of the X-ray luminosity-temperature relation with redshift and take advantage of the recent availability of temperature information for a significant number of high and intermediate redshift X-ray clusters of galaxies. From our analysis, we find a slightly positive evolution in the $L_X-T_X$ relation. This implies a high value of the density parameter of $0.85\\pm0.2 $. However, because the selection of clusters included in our sample is unknown, this can be considered only as a tentative result. A well-controlled X-ray selected survey would provide a more robust answer. XMM will be ideal for such a program. ", + "introduction": "The Press-Schechter (1974) formalism seems to give an accurate determination of the mass function. The reason for this has been a matter of debate, but it has allowed investigation of the non-linear evolution of structure formation in much detail. This has been widely used to put constraints on cosmological parameters, such as the amplitude and the shape of the power spectrum on galaxy clusters scales by comparison with observations. However, cluster masses are not directly measured, therefore it is necessary to establish relations between the mass and ``observables\" such as the X-ray luminosity or the X-ray temperature. As the luminosity of a cluster is difficult to relate to its virial mass from theoretical arguments, it has been argued by some authors that the temperature distribution function is more adequate for a fruitful comparison, although Balland \\& Blanchard (1997) showed that hydrostatic equilibrium does not provide a one-to-one correspondence between mass and temperature. This relation should therefore be taken from the numerical simulations where gas dynamics are taken into account (Evrard, Metzler \\& Navarro 1996). The consequence for cosmology of the observed X-ray luminosity and X-ray temperature distribution functions has been investigated in recent years. OBB97 have shown that a comprehensive description can be constructed, in a consistent way, provided that the relation between luminosity and temperature is specified (from the observations). Such a scheme has been used by OB97 to investigate the cosmological implication of X-ray clusters in an open cosmology and to establish a self-consistent modeling in such a context. They have shown that although the properties of X-ray clusters at redshift zero can be well reproduced in an open model, the redshift evolution is significantly different: this is due to the fact that in open model universes the growth rate of fluctuations is lower than in a Einstein--de Sitter universe. Therefore, at high redshift, a higher number density is expected than in an $\\Omega_0 = 1$ universe. Specifically, they prove (see their section 4) that the redshift evolution of the number of clusters of a given mass, or equivalently of a given apparent temperature, is almost independent of the spectrum of the primordial fluctuations, but that it depends on the mean density of the universe (and on others cosmological parameters of the universe). The evolution of the mass function therefore allows one to measure the mean density of the universe (and the cosmological constant), providing a new cosmological test, based on the dynamics of the universe as a whole. \\begin{table}[ht] \\begin{center} \\begin{tabular}{lllllcllclllcr} \\hline Cluster & z &$T_{X}$ & $L_{bol}$ \\\\ & & ${\\sf keV}$ & $10 ^{44} {\\sf erg/s}$& Ref.\\\\ \\\\ \\hline \\\\ Virgo\t & \t0.0038 & 2.34$^{+0.02}_{-0.02}$ &0.68 &{\\sl 28} \\\\ Centaurus &\t0.01 & 3.9$^{+0.2}_{-0.20}$ &1.22 &{\\sl 1} \\\\ A1060\t &\t0.0114 & 3.1$^{+0.3}_{-0.5}$ & 0.658&{\\sl 22} \\\\ A262\t &\t0.0164 & 2.4$^{+0.3}_{-2.20}$ & 0.85 &{\\sl 1} \\\\ AWM7\t &\t0.0176 & 4.0$^{+0.3}_{-0.20}$ & 2.76 &{\\sl 1}\\\\ A426 & 0.0183 & 6.3$^{+0.3}_{-0.3}$ & 23.1 &{\\sl 1}\\\\ A539 \t &\t0.0205 & 3.0$^{+0.8}_{-0.6}$ & 0.64 &{\\sl 1} \\\\ A1367\t &\t0.0215 & 3.5$^{+0.18}_{-0.18}$ & 2.2 &{\\sl 1} \\\\ 3C 129\t &\t0.0218 & 6.2$^{+0.8}_{-0.6}$ & 3.7 &{\\sl 1} \\\\ A1656\t &\t0.0232 & 8.11$^{+0.07}_{-0.07}$ & 17.2 &{\\sl 1} \\\\ Ophicius &\t0.028 & 9.8$^{+0.7}_{-0.3}$ & 27.4 &{\\sl 22}\\\\ A2199\t &\t0.0299 & 4.5$^{+0.3}_{-0.2}$ & 6.4 &{\\sl 1} \\\\ A496\t &\t0.032 & 3.91$^{+0.06}_{-0.06}$ & 7.16 &{\\sl 1} \\\\ A576\t &\t0.0381 & 4.3$^{+0.5}_{-0.4}$ & 2.94 &{\\sl 1} \\\\ A3558\t &\t0.048 & 5.5$^{+0.3}_{-0.2}$ & 10. &{\\sl 1,8}\\\\ Triangulum &\t0.051 &10.3$^{+0.8}_{-0.8}$ & 30. &{\\sl 12}\\\\ A85\t &\t0.052 & 6.2$^{+0.4}_{-0.5}$ & 15.8&{\\sl 1} \\\\ A3667\t &\t0.053 & 6.5$^{+0.8}_{-0.99}$& 21.4&{\\sl 15}\\\\ A754\t &\t0.0534 & 8.5$^{+0.82}_{-0.82}$ & 29.4 &{\\sl 26} \\\\ A2319\t &\t0.0564 & 10.0$^{+0.7}_{-0.7}$ & 37. &{\\sl 11} \\\\ A2256\t &\t0.0601 & 7.51$^{+0.19}_{-0.19}$ & 24.3&{\\sl 1} \\\\ A1795\t &\t0.0616 & 6.7$^{+0.66}_{-0.66}$ & $23.75^{*}$ &{\\sl16}\\\\ A399 \t &\t0.07 & 7.0$^{+1.0}_{-1.0}$ & 14.2&{\\sl 5}\\\\ A644\t &\t0.0704 & 6.6$^{+0.17}_{-0.17}$ & 23.5&{\\sl 1}\\\\ A401\t &\t0.0748 & 8.0$^{+1.0}_{-1.0}$ & 30.8&{\\sl 5}\\\\ A2142\t &\t0.0899 & 9.0$^{+0.2}_{-0.2}$ & $56^{*}$ &{\\sl 30} \\\\ \\hline \\end{tabular} \\caption{\\small Temperatures and bolometric luminosities for a sample of low-redshift clusters. Asterisks indicate clusters considered as having a strong cooling flow. The numbers in column 6 indicate the references from which the luminosity and the temperature, respectively, are taken ($\\Omega_0 = 1$ and $H_{0}=50$ {{\\rm km/s}/Mpc}). The quoted error bars are given at the 90\\% confidence level.} \\label{Table 1a} \\end{center} \\end{table} \\subsection{Measuring $\\Omega_0$ with clusters} OB97 and OBB97 have used the observed correlation $L_{X}-T_{X}$ to construct a self-consistent modeling of X-ray clusters. The relation between luminosity and temperature they used is: \\begin{equation} L_{bol }= L_1 T_{\\rm keV}^{\\alpha} \\end{equation} with $ L_1 = 0.049 \\; {10^{44} {\\rm erg/s/cm}^2}$ and $\\alpha = 3$. A more recent analysis (Arnaud \\& Evrard, 1997) showed that the $L_{X}-T_{X}$ does have a moderate intrinsic dispersion when cooling flow clusters are removed, with $ L_1 = 0.067 \\; 10^{44} {\\rm erg/s}$ and $\\alpha = 2.89$. By comparing to the EMSS cluster redshift distribution, OB97 have shown that in the absence of evolution of the $L_{X} - T_{X}$ relation, open models with ${\\Omega_{0}} {\\sim} 0.2 $ predict much more clusters than observed while $\\Omega_{0}= 1$ model fits the data reasonably well, although the abundance of high redshift clusters was not very well reproduced. They have investigated the possibility of evolution by allowing the relation $L_{X} - T_{X}$ to change with redshift according to the following form : \\begin{equation} L_{bol} = L_1 (1 + z)^{\\beta} T_{\\rm keV}^{\\alpha} \\label{LTevol} \\end{equation} where ${\\beta}$ is a free parameter which can be derived by fitting the redshift distribution of the cluster EMSS survey. They found that ${\\beta}=1$ for a flat universe (${\\Omega_0}=1$), corresponding to positive evolution, while a significant negative evolution, corresponding to ${\\beta} = - 2.3$ was required for an open universe (${\\Omega_0} \\sim 0.2$). Following a similar approach and by fitting the ROSAT number counts, Kitayama \\& Suto (1997) provided constraints on the parameters of the model and Mathiesen \\& Evrard (1997) reached essentially identical conclusions. Kitayama et al. (1997) has extended the predictions to the SZ counts. The self-similar model predicts $\\alpha = 2.0$ and $\\beta = 1.5$, but $\\alpha = 2.$ and is clearly ruled out by observations, and physical processes specific to the baryonic content have to be advocated. For instance, HA91 assumed an isentropic model. Bower (1997) has recently re-examined this question in more detail. \\section {The high-redshift cluster sample} \\begin{table*}[ht] \\begin{tabular}{lllll} \\hline Cluster & z &$T_{X}$ & $L_{bol}$ \\\\ & & ${\\sf keV}$ & $10 ^{44} {\\sl erg/s}$& Ref.\\\\ \\hline A2204\t &\t0.153 & 8.5$^{+0.4}_{-0.45}$ & 76. &{\\sl 2}\\\\ A3888\t &\t0.168 & 7.9$^{+0.3}_{-1.0}$ & 33. &{\\sl 17}\\\\ A1204 &\t0.17 & 3.6$^{0.13}_{-0.13}$& 14. &{\\sl 18}\\\\ A586\t &\t0.171 & 6.8$^{+0.7}_{-0.67}$ & 17.96&{\\sl 3}\\\\ RXJ 1340+4018& 0.171 & 0.92$^{+0.13}_{-0.13}$ & 0.45&{\\sl 27}\\\\ A2218\t &\t0.175 & 6.72$^{+0.83}_{-0.83}$ & 19.3&{\\sl 1}\\\\ A1689\t &\t0.181 & 8.7$^{+0.51}_{-0.49}$ & 47.9 &{\\sl 3}\\\\ A665\t &\t0.182 & 8.9$^{+0.62}_{-0.61}$ & 28. &{\\sl 3}\\\\ A1763\t &\t0.187 & 9.0$^{+1.02}_{-0.84}$ & 35.5&{\\sl 6,2}\\\\ A1246\t &\t0.187 & 6.3$^{+0.54}_{-0.51}$ & 19.5&{\\sl 2}\\\\ SC 2059-25 & 0.188 & 7.0$^{+6.9}_{-2.2}$ &23.15 &{\\sl 9}\\\\ MS0839.8 &\t0.194 & 3.8$^{+0.4}_{-0.31}$ & 7.4&{\\sl 3} \\\\ A2507 & 0.196 & 9.4$^{+2.7}_{-1.9}$ & 40. &{\\sl 1}\\\\ MS0440\t &\t0.1965 & 5.6$^{+0.80}_{-0.60}$ & 6.1&{\\sl 31,19} \\\\ A2163\t &\t0.201 & 14.6$^{+0.9}_{-0.80}$ & 143. &{\\sl 13} \\\\ A520(MS0451+02)&0.201 & 8.6$^{+0.93}_{-0.90}$ & 17. &{\\sl 4,2} \\\\ A963\t &\t0.206 & 6.76$^{+0.44}_{-0.49}$ & 21.&{\\sl 2}\\\\ A1851\t &\t0.2143 & 5.04$^{+0.8}_{-0.68}$ & 6.23&{\\sl 3} \\\\ A773\t &\t0.217 & 9.6$^{+1.03}_{-0.90}$ & 30. &{\\sl 6,2} \\\\ A1704\t &\t0.219 & 4.5$^{+0.56}_{-0.34}$ & 13. &{\\sl 2}\\\\ A1895\t &\t0.225 & 6.7$^{+1.38}_{-1.05}$ & 10.5&{\\sl 3}\\\\ A2390\t &\t0.228 & 8.9$^{+0.97}_{-0.77}$ & 54. &{\\sl 2} \\\\ A2219\t &\t0.23 &11.8$^{+1.26}_{-0.74}$ & 57. &{\\sl 6,2}\\\\ MS1305.4 &\t0.241 & 2.98$^{+0.52}_{-0.41}$ & 2.86 &{\\sl 2} \\\\ A1835\t &\t0.252 & 9.1$^{+2.10}_{-1.30}$ & $80.^{*}$ &{\\sl 7}\\\\ Zw 7160 &\t0.258 & 5.2$^{+2.2}_{-0.70}$ & $30.^{*}$ &{\\sl 7}\\\\ A348 &\t0.274 & 4.85$^{+0.6}_{-0.6}$ & 3.6 &{\\sl 36}\\\\ A33 &\t0.28 & 4.06$^{+0.5}_{-0.5}$ & 2.4 &{\\sl 36}\\\\ A483\t &\t0.28 & 8.7$^{+3.3}_{-2.2}$ & 48.9 &{\\sl 1}\\\\ \\hline \\end{tabular} \\hspace*{2mm} \\begin{tabular}{lllll} \\hline Cluster & z &$T_{X}$ & $L_{bol}$ \\\\ & & ${\\sf keV}$ & $10 ^{44} {\\sl erg/s}$& Ref.\\\\ \\hline A1758N\t &\t0.28 &10.2$^{+2.3}_{-1.7}$ & 29.&{\\sl 6,2} \\\\ Zw3146\t &\t0.291 & 6.2$^{+1.2.}_{-0.7}$ & $64.^{*}$&{\\sl 7}\\\\ A1722\t &\t0.301 & 5.87$^{+0.51}_{-0.41}$ & 21.&{\\sl 2}\\\\ MS1147.3 & 0.303 & 5.5$^{+1.32}_{-1.00}$ & 6.6& {\\sl 32}\\\\ MS1008\t &\t0.306 & 7.9$^{+2.00}_{-1.65}$ & 16.32&{\\sl 32}\\\\ AC118\t &\t0.308 &9.33$^{+1.09}_{-0.83}$ & 48. &{\\sl 3}\\\\ MS1241.5 & 0.312 & 6.2$^{+3.00}_{-2.15}$ & 6.4& {\\sl 32}\\\\ MS0811.6 & 0.312 & 4.6$^{+1.50}_{-1.00}$ & 4.76& {\\sl 32}\\\\ MS2137 & 0.313 &4.7$^{+0.5}_{-0.3}$ & 34.19&{\\sl 32}\\\\ A1995\t &\t0.318 & 9.44$^{+2.18}_{-1.54}$ & 22.5&{\\sl 3}\\\\ MS0353\t &\t0.32 & 6.2$^{+1.65}_{-1.32}$ & 12.85&{\\sl 32} \\\\ MS1426.4 & 0.32 & 5.5$^{+1.82}_{-1.15}$ & 8.85& {\\sl 32}\\\\ MS1224.7 & 0.326 & 4.3$^{+1.15}_{-1.00}$ & 7.46 &{\\sl 32} \\\\ MS1358\t &\t0.329 & 6.6$^{+0.82}_{-0.82}$ & 18.94&{\\sl 32}\\\\ A959\t &\t0.353 & 6.47$^{+1.15}_{-1.01}$ & 16. &{\\sl 3}\\\\ MS1512 & 0.3726 & 3.8$^{+0.66}_{-0.50}$ & 8.212 &{\\sl 32}\\\\ A370\t &\t0.373 & 6.39$^{+1.02}_{-0.81}$ & 21.7&{\\sl 3} \\\\ Cl 0939+47 & \t0.41 & 2.9$^{+1.3}_{-0.8}$ & 11.&{\\sl 34}\\\\ Cl 09104+4109\t&0.442 & 11.4$^{+3.2}_{-3.2}$ & 75.&{\\sl 20}\\\\ RXJ1347.5-1145& 0.451 & 9.3$^{+1.1}_{-1.0}$ & $210.^{*}$&{\\sl 23}\\\\ A851\t &\t0.451 & 6.7$^{+2.7}_{-1.70}$ & 15.&{\\sl 2}\\\\ 3C295\t & 0.46 & 7.13$^{+2.06}_{-1.35}$ & 24.&{\\sl 2}\\\\ MS0451-03 &\t0.5392 &10.4$^{+1.60}_{-1.30}$ & 46.8&{\\sl 24}\\\\ RXJ 0018.8+160&0.544 & 1.6$^{+0.7}_{-0.4}$ 0.6 & 3.2& {\\sl 29}\\\\ CL0016+16 &\t0.5466 & 7.55$^{1.188}_{-0.957}$ & 55.&{\\sl 10}\\\\ RXJ1716+67 & 0.813 & 6.7$^{+2.0}_{-2.0}$ & 17.72&{\\sl 33}\\\\ MS1054.5-0321&\t0.826 &14.7$^{+4.6}_{-3.5}$ & 42.&{\\sl 25}\\\\ AXJ2019 &\t1.0 &8.6$^{+6.9}_{-4.9}$ & 19 &{\\sl 35}\\\\ & & & & \\\\ \\hline \\end{tabular} \\caption{\\small X-ray temperatures and bolometric luminosities for high redshift clusters. \\label{Table 1b} } ${\\bf Ref.}$ {\\sl (1)} David et al., 1993; {\\sl (2)} Mushotzky \\& Scharf, 1997; {\\sl (3)} Tsuru et al., 1996; {\\sl (4)} Nichol et al., 1997; {\\sl (5)} Fujita et al. 1997; {\\sl (6)} Ebeling et al. 1995; {\\sl (7)} Allen et al. 1995; {\\sl (8)} Markevitch \\& Vikhlinin, 1997; {\\sl (9)} Arnaud et al., 1991; {\\sl (10)} Hughes \\& Birkinshaw, 1997; {\\sl (11)} Markevitch, 1996; {\\sl (12)} Markevitch et al., 1996; {\\sl (13)} Elbaz et al., 1996; {\\sl (14)} Tamura et al., 1996; {\\sl (15)} Knopp et al., 1996; {\\sl (16)} Briel \\& Henry, 1996; {\\sl (17)} White \\& Fabian, 1996; {\\sl (18)} Matsuura et al., 1996 ; {\\sl (19)} Gioia private communication ; {\\sl (20)} Hall et al., 1996; {\\sl (21)} Hamana et al., 1997; {\\sl (22)} Matsuzawa,et al., 1996 ; {\\sl (23)} Schindler et al., 1996; {\\sl (24)} Donahue, 1996 ; {\\sl (25)} Donahue, et al., 1997 ; {\\sl (26)} Henriksen \\& Markevitch, 1996; {\\sl (27)} T.J, Ponman et al., 1994; {\\sl (28)} Arnaud \\& Evrard, 1997; {\\sl (29)} Connolly et al., 1996; {\\sl (30)} White et al., 1994; {\\sl (31)} Gioia \\& Luppino, 1994; {\\sl (32)} Henry 1997; {\\sl (33)} Henry \\& Gioia 1997; {\\sl (34)} Schindler, 1997; {\\sl (35)} Hattori et al., 1997; {\\sl (36)} Colafrancesco, 1997; \\end{table*} In order to perform this new cosmological test, we have collected from the literature all the information on temperature measurements of X-ray clusters. Although the number of high redshift clusters ($z > 0.4$) with an estimation of the X-ray temperature remains small, ASCA observations of clusters of relatively high redshift ($z \\approx 0.3$) has begin to appear in the literature. We have tried to compile all the existing X-ray clusters with redshift greater than 0.15 for which the temperature has been measured with various X-ray satellites, mostly ASCA. This represents 57 clusters whose properties are summarized in Table 2.\\\\ In order to address the question of the possible evolution of clusters, we have decided to include some clusters for which a reliable mass estimate is available and derived the corresponding temperature. This has led us to include the CNOC survey of clusters by C97, as it contains accurate velocity dispersion measurements, and the sample of Smail et al. (1997) (hereafter S97), who have used deep HST images to study weak shear of background galaxies by distant clusters of galaxies to estimate their masses. From these two samples, only clusters for which temperature information was not available were retained in the analysis. The final compilation contains a total number of 57 clusters at high redshift ($z > 0.15$), 26 being at redshift greater than $ 0.3$. The most distant redshift clusters of the sample are MS1054.5-0321 at z = 0.83 which has been measured by Donahue et al. (1997) and AXJ2019 at z = 1 by Hattori et al. (1997). An additional list of clusters at low redshift, for which accurate temperature information was available, has been added to the sample for the purpose of our analysis, but we do not attempt to be complete. This set is used as a template (Table 1). \\\\ Ideally, in order to perform the cosmological test, it would be better to estimate the evolution of $L_X-T_X$ from the EMSS sample itself. However, too few measurements exist up to now. Still, our sample contains a number of high redshift clusters which is comparable to that of the EMSS sample in the various redshift ranges. For comparison, the EMSS survey (Gioia \\& Luppino 1984) has the same number of clusters at $z > 0.3$ but a total number of $z > 0.15$ of only 49 (Figure 1). This is important as it does mean that we are testing the evolution of the $L_X-T_X$ relation over the same redshift range as OB97 have investigated, with a similar number of clusters involved. \\begin{figure}[h] \\epsfxsize=9cm \\centerline{\\epsfbox{hist3.ps}} \\caption{\\small \\sl The redshift distribution in our sample (full line) compared with the one in the EMSS cluster survey (hashed line). } \\label{Fig. 1} \\end{figure} \\subsection{Fluxes} As the available data come from different satellites, the published luminosities are given in various energy bands. Therefore, when the bolometric luminosity was not available, it has been estimated by using a Raymond--Smith code, taking an average metallicity of 0.33 when its value was not available (see Tables 1, 2). Flux calibration is a serious worry when data from different satellites are used. Arnaud \\& Evrard (1997) have shown that GINGA and ROSAT fluxes agree very well, while an offset in the calibration of EXOSAT is suspected. In some cases, several flux estimates are given, which disagree from time to time. Our fluxes were taken from different authors: David et al. (1991), Mushotzky \\& Scharf (1997) (hereafter {M\\&S97}), Ebeling (1996) from ROSAT PSPC observations and Nichol {\\it et al.} (1997) from HRI observations. In general, we have preferred ROSAT measurements when available. \\subsection{Cooling flows} Cooling flow clusters present a central enhancement in their luminosity, in a region were the cooling time of the gas is shorter than the Hubble time. The inclusion of such clusters in our analysis is problematic : as they are more luminous than ``normal clusters\", they might introduce a bias. However, as the EMSS is flux selected, there is no reason to assume that cooling flow clusters are not present in the EMSS sample as well. Their point-like nature could even represent a bias favoring their presence in the EMSS, as the high background makes the detection of extended source more dificult. Still, it is also conceivable that cooling flow clusters are overabundant in the general compilation compared to the EMSS sample. We have therefore applied a correction to clusters for which the presence of a cooling flow was known: only the flux outside the cooling flow was taken into account. Such cooling flow clusters are flagged by an asterisk. In practice, this correction is never larger than 50\\%. This has been applied only to a few number of clusters. \\\\ \\begin{table}[t] \\begin{center} \\begin{tabular}{llrrr} \\hline Cluster & z & ${\\sigma}$~~ &$T_{est}^{v}$&$T_{X}$ \\\\ & & ${\\sf km/s}$& ${\\sf keV}$ &${\\sf keV}$\\\\ \\hline A2390 &0.228 & 1104 & 9.95 & $8.9^{a}$ \\\\ MS0440 &0.1965 & 606 & 3.00 & $5.6^{a}$ \\\\ MS1008 &0.306 & 1054 & 9.06 & $7.7^{a}$ \\\\ MS1358 &0.329 & 934 & 7.12 & $6.5^{a}$ \\\\ MS1512 &0.373 & 690 & 3.88 & $4.2^{a}$ \\\\ MS0451+02 &0.2011&1031 & 8.70 & $8.6^{a}$ \\\\ MS0451-03 &0.5392&1371 & 15.30 & $10.4^{b}$ \\\\ MS0839 &0.1928 & 756 & 4.66 & $3.8^{c}$ \\\\ MS1455 &0.2568 &1133 & 10.50 & $5.0^{a}$ \\\\ MS1224 &0.3255 & 802 & 5.25 & $4.3^{d}$ \\\\ \\hline \\end{tabular} \\caption {\\small Comparison between virial temperatures estimated from the velocity dispersion ${\\sigma}$ (CNOC Survey C97) with measured temperatures : $^{a}$M\\&S97; $^{b}$Donahue 1997; $^{c}$Tsuru et al. 1997; $^{d}$Henry 1997 .} \\end{center} \\end{table} \\subsection{ Deriving the temperature } As explained in section 3, our sample is not based on clusters with measured temperature only: it also contains distant clusters for which the temperature information has been estimated in an indirect way. From the lensing analysis, S97 have used deep WFPC-2 imaging of 12 distant clusters at redshifts between $z = 0.17$ and 0.56. Using the distortion of faint galaxies detected in these fields, they measured the mean shear and inferred a mass estimate within a radius of $200 h^{-1}$ kpc from the cluster lens center, assuming a singular isothermal profile $\\propto r^{-2}$. We have used these masses to derive the X-ray temperature of these 11 distant clusters, applying the following scaling relation:\\\\ \\begin{equation} T_{X} = 5.38 M_{lens} {\\rm keV} \\end{equation} where $ M_{lens}$ is in units of $10^{14}h^{-1}{\\rm M}_{\\odot}$. For the C97 clusters, we have converted the virial masses derived inside the total radial extent of the sample from the measured velocity dispersions to an X-ray temperature. For this, we used the scaling relations derived from Evrard's numerical simulations (Evrard 1989, Evrard 1997), we infer a relation between velocity dispersion and temperature: \\begin{equation} T_{X}=({\\sigma}/350{\\rm km/s})^{2} {\\rm keV} \\end{equation} In some cases, it has been possible to compare the temperature as estimated from (3) and (4) with satellite temperature measurements (see Tables 3 and 4). Although the number of clusters for which the information is available is small, the agreement is rather good (Figure 2). This suggests that lensing and the virial mass estimates are essentially in agreement with the X-ray mass at a level better than a factor of two, although one can notice a slight tendency towards overestimation, but the samples are too small to draw any firm conclusion. \\begin{table}[t] \\begin{center} \\begin{tabular}{lllrr} \\hline Cluster & z & $M_{lens}$& $T_{X}^{lens}$& $T_{X}$\\\\ & & $10^{14}{\\sf M_{\\odot}}$& {\\sf keV} & {\\sf keV}\\\\ \\hline A2218 & 0.17 & 1.05 & 5.65 & $6.35^{a}$ \\\\ AC118 & 0.31 & 1.85 & 9.92 & $9.95^{b}$ \\\\ CL0016 & 0.55 & 1.87 & 10.06 & $7.55^{c}$\\\\ CL0939 & 0.41 & 0.73 & 3.93 & $2.90^{d}$ \\\\ 3C295 & 0.46 & 2.35 & 12.64 & $7.50^{a}$\\\\ \\hline \\end{tabular} \\caption {\\small Comparison between estimated temperatures from $ M_{lens}$ (weak lensing, S97) with measured temperatures: $^{a}$M\\&S97; $^{b}$Tsuru et al. 1997; $^{c}$Hughes \\& Birkinshaw, 1997; $^{d}$Schindler 1997.} \\end{center} \\end{table} \\begin{figure}[h] \\epsfxsize=9cm \\centerline{\\epsfbox{comp_Tx.ps}} \\caption{\\small \\sl Comparison between estimated temperatures from $ M_{lens}$ as inferred from weak lensing S97 (triangles) and from CNOC survey velocity dispersions (crosses) with measured temperatures $T_{X}$.} \\label{Fig. 2b} \\end{figure} \\noindent \\section {Method and Results} \\subsection{Investigating the possible evolution} Ideally, one would like to estimate directly the $L_{X}-T_{X}$ in a redshift bin centered on a value of the redshift as high as possible. However, since the number of clusters decreases rapidly for redshifts greater than 0.25, it becomes difficult to get any information from this method at high redshift: dividing the sample in several redshift bins and trying to fit the luminosity--temperature relation in the different redshift bins becomes unpractical for redshifts greater than 0.35. Another method has been used which consists in plotting the mean temperature of clusters above some threshold luminosity in order to minimize any possible systematic effect (Arnaud et al, 1991, M\\&S97). It remains possible, however, that higher redshift clusters are brighter in the mean, introducing a bias in the sample. Furthermore, this method results in a rude elimination of some of the data. We have tried to find an efficient estimator of the evolution which is adequate for the kind of evolution introduced by OB97. To get round this problem we have introduced a new evolution estimator. For each cluster $i$,with measured $L_i, T_i,$ we have estimated the following quantity: \\begin{equation} C_i = \\frac{L_i}{L_1T_i^{\\alpha}} \\label{eq:cz} \\end{equation} where $ L_1 = 0.049 {10^{44} {\\rm erg/s/cm}^2}$ and $\\alpha = 3$. The dependance of $C_i$ on redshift probes the evolution: clearly, if the cluster population is not evolving, the mean value of this quantity should remain constant with redshift. Note that it is possible that the $L_{X}-T_{X}$ evolves and that the measured $C$ will not probe this evolution, this would need, however, some kind of conspiracy. Because $L_X$ is estimated from the apparent flux, $C(z)$ also contains a term coming from the cosmological parameters of the universe in the luminosity distance. In practice, $L_X$ is estimated in a Einstein-de Sitter universe; therefore, the theoretical value $C(z)$ has to be corrected when one is comparing data with low-density universe predictions (however, this term is small, as can be seen in fig. \\ref{fig:cz}). It is also clear that this method can be applied without removing any data, and that it takes fully taking into account the information in redshift. We have applied this test to our sample using the OB97 parameters for the $L_{X}-T_{X}$ relation. For each cluster, we have computed the quantity given by (5) and estimated an uncertainty range on this quantity, neglecting the uncertainty in the luminosity. The result is plotted in Fig \\ref{fig:cz}. \\begin{figure*} \\centerline{\\psfig {figure=C-fit.ps,height=12cm}}% \\caption{\\small In this figure we present the coefficient $C_i$ for all the clusters in our sample with redshift smaller than 0.6. The filled circles are used for actual X-ray measurement of the temperature, filled squares are for weak lensing measurements of S97 and filled triangles are for the CNOC clusters. The error bars are $1 \\sigma$. The thick line represents the best fitting power--law, the shaded area represents an estimate of the 90\\% confidence range. The prediction of a low density universe ($\\Omega_0 = 0.2$) is represented by the thin line. The dashed, thin line is the curve when the correction in the luminosity distance is not taken into account in the expression of $C(z)$.} \\label{fig:cz} \\end{figure*} The measure of evolution can now be directly obtained by fitting a power law to the data: \\begin{equation} C(z) = \\alpha(1+z)^{\\beta} \\end{equation} in which $\\alpha$ and $\\beta$ are determined by a likelihood analysis. However, the intrinsic dispersion of the values of $C$ is often higher than the uncertainty on the coefficient itself, due to the errors in the temperature. We have therefore estimated the intrinsic dispersion of $C$ from our template sample and added it in quadrature with the uncertainty on the individual values of the coefficient. We have also checked that the results are insensitive to the assumed dispersion. In order to investigate the robustness of our results, we have performed various analyses, whose results are summarized in Table 5: for various sub-samples, we have reported the best estimated parameters from the likelihood analysis and the 68\\% and 90\\% confidence limits. The expected value of $\\alpha$ is 1. If only clusters with redshift greater than 0.15 are used, there is a degeneracy between $\\alpha$ and $\\beta$: in the two-parameter plane, the confidence domain looks like an elongated ellipse, allowing a wide range of $\\beta$, but for unrealistic values of $\\alpha$. This is the main motivation to include a set of low redshift clusters in our sample, before performing the likelihood analysis. With this approach, $\\alpha = 1$ value always falls in the 90\\% confidence range when $\\alpha$ and $\\beta$ are determined by a likelihood analysis. In our analysis, we have generally assumed a gaussian distribution of the errors, which is equivalent to a $\\chi^2$ minimization. When the full sample is used, the $\\chi^2$ of the best fitting model is rather poor : $\\chi^2$ is equal to 219 for 97 clusters. This is due to a few outliers, which are far away from the general trend. This can be seen by using the $l_1$ norm instead of the standard $\\chi^2$: five clusters were found to have high $C$ value (A1204, RXJ 1340+4018, MS2137, Cl 0939+47, RXJ 0018+16) and were removed in most sub-samples (in this case the sub-sample is flagged by -o in Table 5).\\\\ \\begin{table*} \\begin{center} \\begin{tabular}{lrrllc} \\hline \\vspace*{-2mm}\\\\ Sub-sample & $N _c$ & $\\delta \\chi^2$ & $\\alpha\\pm 68\\% \\pm 90\\%$ & $\\beta\\pm 68\\% \\pm 90\\%$ & $\\Omega_0\\pm 90\\%$ \\vspace*{1mm}\\\\ \\hline \\vspace*{-2mm}\\\\ Complete sample & 71 & 193. & $1.15^{+0.10+0.20}_{-0.10-0.15}$ & \\hspace*{2.5mm}$ 0.00^{+0.30+0.50}_{-0.30-0.60}$ & $0.75^{+0.12}_{-0.15}$\\vspace*{1mm}\\\\ ($L_1$ norm) & 71 & 63. & $1.05^{+0.15+0.25}_{-0.15-0.25}$ & $ -0.10^{+0.60+1.00}_{-0.60-1.10}$ & $0.72^{+0.25}_{-0.27}$\\vspace*{1mm}\\\\ X-ray sample & 57 & 182. & $1.14^{+0.10+0.15}_{-0.10-0.15}$ & $ +0.28^{+0.30+0.50}_{-0.35-0.60}$ & $0.82^{+0.15}_{-0.17}$\\vspace*{1mm}\\\\ & 57 & 184. & {\\bf 1.00 } & $ +0.65^{+0.25+0.35}_{-0.20-0.35}$ & $0.91^{+0.09}_{-0.09}$\\vspace*{1mm}\\\\ X-ray sample -o & 52 & 68. & $1.07^{+0.10+0.17}_{-0.05-0.13}$ & $+0.20^{+0.30+0.60}_{-0.35-0.70}$ & $0.80^{+0.15}_{-0.17}$\\vspace*{1mm}\\\\ & 52 & 69. & {\\bf 1.00 } & $+0.40^{+0.20+0.40}_{-0.30-0.50}$ & $0.85^{+0.10}_{-0.12}$\\vspace*{1mm}\\\\ X-ray sample $0.150.15$), the $\\chi^2$ value in col. 3 ( the contribution of the low redshift cluster subset to the $\\chi^2$ is nearly constant and equal to 27.), the best fit estimates of $\\alpha$ and $\\beta$ with the uncertainties at 68\\% and 90\\% confidence level are given in col. 4 and col. 5 respectively and the corresponding value of $\\Omega_0$ as inferred from eq. 7 and the 90\\% uncertainties are given in col. 6. } \\label{tabres} \\end{center} \\end{table*} The X-ray data, even when they are split into various sub-samples (like the EMSS X-ray clusters alone or the Henry set alone) always lead to a high value of the best $\\beta$, in the range $[0.-1.0]$. It is important to note that no systematic trend is found: the various X-ray selected samples always give a $\\beta$ in the same range, with an uncertainty at worst of the order of 0.5. Therefore the range $[0.-1.0]$ can be taken as the 90\\% confidence interval. The only two samples which do not lead to a high value of $\\beta$ are the CNOC sample and the S97 sample with their estimated temperatures (from virial and lensing mass estimates). However, when these samples are restricted to clusters for which the temperature has been measured (the X-ray CNOC and X-ray S97 sample in Table 3), the resulting $\\beta$ is much higher and in agreement with other analysis based on X-ray data. This may be due to the fact that lensing and virial mass estimates are higher than X-ray masses, which introduces a bias toward lower $C$ and, consequently, to artificially lower $\\beta$ at high $z$. \\\\ Our result confirms previous investigations: available data on X-ray clusters are consistent with the lack of significant evolution in the $L_X-T_X$ relation. Finally, it is somewhat troublesome that the highest redshift clusters ($Z > 0.8$) lead to a small value of $C$. Better temperature and flux measurements will be of great interest, and we obviously need more (very) high redshift clusters temperature measurements.\\\\ \\begin{figure} \\vspace*{-1.5cm} \\hspace*{1cm} \\centerline{\\psfig {figure=omega-C.ps,height=11cm}}% \\vspace*{-2cm} \\caption{\\small \\sl Comparison of the predicted values of the coefficient to the values we have estimated from the observations. The grey area corresponds to our 90\\% confidence range for $\\beta$. The solid thick line corresponds to the best evolution law when the EMSS redshift distribution is fitted, the dashed (resp. dotted-dashed) correspond to an estimation of the 68\\% contour (95\\% resp.). } \\label{fig.om} \\end{figure} \\subsection{What do the data tell us?} As discussed previously, the knowledge of the evolution of the number density of clusters with redshift is a powerful cosmological test. By fitting the EMSS cluster redshift distribution, OB97 showed that the knowledge of evolution of the temperature-luminosity relation provides an alternative way to measure the mean density of the universe. We have iterated OB97 analysis for various values of the density parameter: the temperature distribution $N(T_x)$ was fitted in order to derive the power spectrum index of the fluctuations as well as its normalization; the best $\\beta$ in eq. (\\ref{LTevol}) was then determined by fitting the EMSS redshift distribution, as well as the 1- and 2- sigma interval. The best fitting parameter $\\beta$ is tighly related to $\\Omega_0$ accordingly to the following relation: \\begin{equation} \\beta = 4.\\times \\Omega_0 -3. \\end{equation} This relation, as well as the 1- and 2- sigma interval contours, are presented on figure \\ref{fig.om}. The high value of the slope illustrates the strong dependence on ${\\Omega_0}$, highlighting the power of this test.\\\\ From the range of $\\beta$ we estimated from the data, we can obtain an estimation of $\\Omega_0$. This is presented in figure \\ref{fig.om}, where we have plotted the value of $\\beta$ which is necessary to fit the EMSS cluster redshift distribution as well as the confidence range. The grey area represents the range $[0.-1.]$, which is used as our 90\\% range obtained from the likelihood analysis. As one can see, the range of values we obtained favors a high density universe with a formal determination of the mean density parameter: $\\Omega_0 = 0.85^{+0.2}_{-0.2}$. The main potential problem in our analysis is that the correspondence between $\\beta$ and $\\Omega_0$ has been obtained by OB97 from the EMSS survey, while the sample we used consists of clusters for which selection rules cannot be defined in a simple way. There are some limitations in the present work which imply that the conclusion should be considered as only preliminary: firstly, because the selection function of X-ray clusters in EMSS is not well understood, it conceivable that a substantial number of clusters have been missed (for instance because their surface brightness was too low). This is not supported by the fact that the modeling of X-ray clusters as done by OB97 predicted correctly the abundance of faint clusters as detected by ROSAT. Secondly, it is possible that our sample contains clusters which have preferentially high $\\beta$ as we have already mentioned; it is possible that these clusters suffer from a systematic bias favoring high $\\beta$. Still, we do not find any evidence of such a bias. For instance, the typical luminosity in the sample does not seem to increase with redshift. In general, no systematic tendency was found by eye inspection. Thirdly, OB97 normalized the models at $z = 0$ by use of the HA91 temperature distribution function, which may suffer from systematic uncertainties (Eke at al., 1996, Henry, 1997). A more detailed investigation of these various questions would need a Monte-Carlo simulation as well as a systematic investigation of possible biases. This will be addressed in a future paper. ", + "conclusions": "In this paper, we have addressed the question of the cosmological evolution of the $L_{X}-T_{X}$ relation and we have investigated the possible cosmological implications of this evolution (or lack of it). The new indicator $C(z)$ is well adapted to measure the possible evolution of the $L_{X}-T_{X}$ relation when the number of available clusters is small. We have applied this measure to a sample of high and intermediate redshift clusters for which the temperature information is available. We have found no strong evidence of evolution of the $L_{X}-T_{X}$ relation, in agreement with other works. Despite the fact that our sample is not complete and that the number of high redshift clusters is not large, we have shown that our method is robust and leads to statistically significant result. We have furthermore used our results on $L_{X}-T_{X}$ evolution to constrain the value of ${\\Omega_o}$ accordingly to the test of OB97. This is indeed the first time that this test is applied to observations. This test is extremely powerful because it results from a fundamental difference between high-- and low--density universes: the rate of structure formation. Therefore, it provides a global test of the mean density of the universe, rather than a local dynamical one, as are classical $M/L$ estimates. One can therefore expect to obtain in this way a definitive answer on the value of the mean density of the universe. The absence of negative evolution in the $L_{X}-T_{X}$ relation, as we have found, provides an indication of a high-density universe: accordingly to our analysis, the range of evolution we find is consistent with $\\Omega_0 = 0.85^{+0.2}_{-0.2}$ (at 90\\% confidence level). The fact that our sample is not drawn from an X-ray selected sample implies a possible bias and therefore our conclusion can be considered as only preliminary. A more robust answer on the mean density of the universe can be obtained only from a well-controlled X-ray selected sample of clusters. Our work shows that even if no definitive conclusion can be drawn, a reasonable number of X-ray temperature measurements can provide a very interesting answer on the mean density of the universe and that open model universes seem to be facing a serious problem. It is interesting to note that Barbosa et al. (1996) have also pointed out an other piece of evidence coming from clusters which disfavors open model universes. It is realistic to envisage that a definitive answer could be obtained with XMM by a follow-up of a sample of high redshift clusters selected from an X-ray flux limited survey." + }, + "9708/astro-ph9708082_arXiv.txt": { + "abstract": "We re-examine the interaction of ultra high energy nuclei with the microwave background radiation. We find that the giant dipole resonance leaves a new signature in the differential energy spectrum of iron sources located around 3 Mpc: A depression before the bump which is followed by the expected cutoff. \\noindent {\\it PACS number(s):} 96.40, 13.85.T, 98.70.S, 98.70.V ", + "introduction": "In 1966 Greisen, Zatsepin and Kuz'min \\cite{G,Z} noted that the microwave background radiation (MBR) makes the universe opaque to cosmic rays of sufficiently high energy, yielding a steep drop in the energy cosmic ray spectrum at approximately $ 5 \\times 10^{19}$ eV (GZK cutoff). More recently, a fresh interest in the topic has been rekindled since several extensive air showers have been observed which imply the arrival of cosmic rays with energies above $10^{20}$ eV. In particular, the Akeno Giant Air Shower Array (AGASA) experiment recorded an event with energy 1.7 - 2.6 $\\times 10^{20}$ eV \\cite{Yoshi,Hasha}, the Fly's Eye experiment reported the highest energy cosmic ray event ever detected on Earth, with an energy 2.3 - 4.1 $\\times 10^{20}$ eV \\cite{Bird1,Bird2}, both events being well above the GZK cutoff. Deepening the mystery, the identification of the primary particle in these showers is still uncertain. On the one hand, the Fly's Eye group claims that there is evidence of a transition from a spectrum dominated by heavy nuclei to one of a predominantly light composition \\cite{Bird1}, while on the other hand, it has also been suggested that a medium mass nucleus also fits the shower profile of the highest energy Fly's Eye event \\cite{H}. In addition, there is an unexpected energy gap before these events. Although heavy nuclei can be accelerated to high terminal energies by ``bottom up'' mechanisms, one should note that, for energies above 100 EeV the range of the corresponding sources is limited to a few Mpc \\cite{JCronin}. Sigl and co-workers \\cite{Sigl} have analysed the structure of the high energy end of the cosmic ray spectrum. They found that most ``bottom up'' models can be ruled out except for those involving a nearby source, which is consistent with data at the 1$\\sigma$ level. Their argument for this is that a nearby source can account for the ultrahigh energy events but would also produce events in the apparent gap in data obtained to date. In this direction, Elbert and Sommers have suggested that the highest energy event recorded by Fly's Eye, could have been accelerated in the neighborhood of M82, which is around 3 Mpc away \\cite{ES,W}. In relation to the aforementioned possibilities, we have re-examined the interaction of ultrahigh energy nuclei with the microwave background radiation and we have found a new feature in the ultrahigh energy cosmic ray spectrum from iron sources located around 3 Mpc which forms the motivation for the present article. ", + "conclusions": "We have studied the interaction of ultra high energy nuclei with the MBR. We have presented a parametrization of the fractional energy loss for Lorentz factors up to $10^{11}$ that allows us to analyse the evolution of the energy spectrum for different nuclei sources. When considering an iron source located around 3 Mpc, the spectrum exhibits a depression before a bump not previously reported. In the light of this finding it is tempting to speculate whether the apparent gap in the existing data is due to the relative weight of the depression and the bump if a source of iron nuclei is responsible for the end of the cosmic ray spectrum. This speculation, if true, reclaims \"botton up\" models as a possible scenario for the origin of the highest energy cosmic rays. The limited statistics in the observed data make it impossible to resolve the question definitively at this time, and we are obliged to present this idea as a hypothesis to be tested by experiment. The existence of a cutoff or a gap which might be present in the observed spectrum is of fundamental interest in cosmic ray physics, allowing stringent tests of existing models. The future Pierre Auger Project \\cite{Desrep} should provide enough statistics for a final veredict on these open questions, and in particular on the ideas discussed in this paper." + }, + "9708/astro-ph9708031_arXiv.txt": { + "abstract": "\\noindent We investigate the effects on cosmological clustering statistics of empirical biasing, where the galaxy distribution is a local transformation of the present-day Eulerian density field. The effects of the suppression of galaxy numbers in voids, and their enhancement in regions of high density, are considered, independently and in combination. We compare results from numerical simulations with the predictions of simple analytic models. We find that the bias is generally scale-dependent, so that the shape of the galaxy power spectrum differs from that of the underlying mass distribution. The degree of bias is always a monotonic function of scale, tending to an asymptotic value on scales where the density fluctuations are linear. The scale dependence is often rather weak, with many reasonable prescriptions giving a bias which is nearly independent of scale. We have investigated whether such an Eulerian bias can reconcile a range of theoretical power spectra with the twin requirements of fitting the galaxy power spectrum and reproducing the observed mass-to-light ratios in clusters. It is not possible to satisfy these constraints for any member of the family of CDM-like power spectra in an Einstein - de Sitter universe when normalised to match {\\em COBE\\/} on large scales and galaxy cluster abundances on intermediate scales. We discuss what modifications of the mass power spectrum might produce agreement with the observational data. ", + "introduction": "The great advances in observational cosmology during the past decade have provided a wealth of new data on galaxy clustering. These results facilitate a phenomenological approach to the study of large-scale structure, in which the mass power spectrum is reconstructed directly from the observed clustering statistics (e.g. Peacock \\& Dodds 1994, PD94). This method is complementary to the hitherto conventional technique of judging a pet cosmogony on the basis of its predictions for a series of observationally-determined quantities, and has some advantages so long as no specific model appears capable of accounting for the full set of observational data. There is, however, a fundamental problem with both the empirical and the a priori approaches to large-scale structure. The linear mass power spectrum that one deduces from the observed clustering depends on, and may be very sensitive to, an assumed relationship between the distributions of mass and galaxies in the Universe, and ignorance of the detailed processes through which galaxies form leaves this relationship poorly constrained. Different classes of galaxy are observed to cluster differently (see, e.g., the compilation by PD94), implying that at least some galaxies are biased and do not directly trace the mass distribution. Bias is also required by advocates of an Einstein -- de Sitter universe, to reconcile observed cluster $M/L$ values that imply $\\Omega \\simeq 0.2$ with the assumed critical density. In the absence of a complete understanding of galaxy formation and evolution, large-scale structure theorists are left to model bias in as plausible a manner as they can. Many such attempts have used the high-peak bias method (Davis et al. 1985; Bardeen et al. 1986, BBKS). We will term this a {\\it Lagrangian\\/} bias scheme, since it identifies the sites of nascent objects (galaxies or clusters) with peaks in the {\\it initial\\/} density field, smoothed on a scale appropriate to the mass of the objects under consideration. If the initial density field is Gaussian, then these peaks are more strongly clustered that the mass distribution as a whole, producing the desired bias. {\\em N\\/}--body simulations (e.g. Katz, Quinn \\& Gelb 1993) have revealed a poor correspondence between the particles found in galaxy halos selected at late times and those located at appropriate peaks in the initial density field, thus casting doubt on Lagrangian galaxy bias prescriptions (although the work of Mo \\& White 1996 and Mo, Jing \\& White 1996 suggests Lagrangian bias may work statistically): this is not surprising, since the density field smoothed on the scale of galaxies is highly non-linear. For cluster-sized halos, however, Lagrangian bias may still be appropriate (Kaiser 1984; Cole \\& Kaiser 1989; Mann, Heavens \\& Peacock 1993), since their spatial distribution has undergone far less dynamical evolution. In reality, galaxy formation must be more complex than this, involving a range of feedback mechanisms (Dekel \\& Rees 1987; Babul \\& White 1991). One might hope that all these issues would eventually be clarified by large numerical simulations which include a hydrodynamical treatment of baryons, as well as dissipationless dark matter, but such codes are in their infancy. An intriguing early result from one such code was the observation by Cen \\& Ostriker (1992; CO) of a tight correlation between the present-day (Eulerian) density field, $\\rho_{\\rm m}$, and the local number density, $\\rho_{\\rm g}$, of galaxies, formed using what those authors termed a `heuristic but plausible' prescription, which creates a dissipationless proto-galactic particle wherever the local baryonic component is sufficiently dense, rapidly cooling, contracting and Jeans unstable. Together with the observed uniformity of $M/L$ values as a function of scale, these results motivate the general idea of {\\it Eulerian\\/} bias models, in which the galaxy number density at redshift zero is a function of the local value of the evolved mass density field. Some general aspects of such prescriptions were considered by Coles (1993), who deduced various important inequalities relating the mass and galaxy correlation functions in these models. Our approach here is as follows. We make no assumptions concerning the physical processes that lead to bias, simply taking their overall effect to be a local transformation of the Eulerian density field, the principal features of which are the suppression of galaxy numbers in low-density regions and their enhancement in regions of high density. In Section 2 we outline the set of bias prescriptions combining these components which are to be studied here, using the combination of numerical and analytical methods described in Section 3, while Section 4 shows the effect of biasing CDM-like mass models by these methods. We focus on the question of the scale-dependence of the Fourier space bias parameter, $b(k)$, defined as the square root of the ratio of the galaxy and mass power spectra: \\beq b(k)\\equiv [\\Delta_{\\rm g}^2(k)/\\Delta_{\\rm m}^2(k)]^{1/2}, \\eeq where $\\Delta_{\\rm g}^2(k)$ and $\\Delta_{\\rm m}^2(k)$ are, respectively, the galaxy and mass power spectra. We employ the dimensionless form of the power spectrum, $\\Delta^2(k)$, which is defined to be the variance per $\\ln k$ [i.e. $\\Delta^2(k)=\\rmd \\sigma^2/ \\rmd \\ln k \\propto k^3 P(k)$]. Our results indicate that $b(k)$ tends to a constant value on scales where the density field is linear, with only a weak tendency to change with scale for most reasonable bias prescriptions: we shall denote this asymptotic large-scale bias value by $b_\\infty$. Finally, we attempt to find a combination of a CDM-like mass model with a reasonable normalisation and a biasing prescription that can reproduce the observed power spectrum of galaxy clustering, while reconciling observed cluster $M/L$ ratios with the critical density assumed in Einstein -- de Sitter models. Details of this exercise and its results are given in Section 5, and are discussed in Section 6, where we also present the conclusions we draw from them. ", + "conclusions": "We have investigated the effect of local biasing on the clustering statistics of galaxies. In general, we find that local biasing gives rise to a scale-dependent bias, so the shape of galaxy power spectrum no longer directly reflects the underlying mass distribution. However, the change in $b(k)$ with scale is often rather modest: if $b_\\infty \\simeq 1.5$, as suggested by observations and {\\em COBE} normalisation, then the models we have investigated would only predict at most $b(k) \\simeq 3$ at $k \\simeq 1 \\hompc$. The scale dependence of the bias appears to be monotonic, but $b(k)$ does not decrease with scale in all cases, showing how the inequalities of Coles (1993) break down once the density field is sufficiently non-linear that it is no longer well described by a Gaussian field, or a local transformation thereof. The $b(k)$ curves produced by our biasing prescriptions are also quite smooth, so it would appear that any model capable of producing sharp changes in galaxy power with respect to that of the mass would have to involve non-local bias. The scale dependence of bias clearly presents difficulties in the interpretation of observed galaxy clustering, by allowing a wider range of mass power spectra to map onto the galaxy power spectrum data through the use of a suitably curved $b(k)$. Interpretation of even these models is simplified on large scales, however, where the bias tends to a constant value. We have looked at a range of galaxy clustering models in the context of an Einstein - de Sitter universe, to see if any of these models can simultaneously fit the APM power spectrum and the observed mass-to-light ratios in clusters of galaxies. The strong conclusion we reach is that {\\em COBE\\/}-normalised CDM-like models with $\\sigma_8$ fixed as required by the observed abundances of clusters do not succeed, given the biasing functions we have explored. In addition, the models fail to reproduce the required mass-to-light ratios by large factors, so if any successful bias function were found, it would probably be rather pathological. There are several ways to improve the agreement of the models with observation. One is to lower the density below $\\Omega=1$, thus reducing or removing the constraints on our models in the high-density regions. The lower $\\Omega$ has the effect of increasing the required $\\sigma_8$ from cluster abundances, but the non-linear effects on the mass spectrum depend on $\\Omega$ only on scales smaller than those we have considered here (Peacock \\& Dodds 1996). The second is to change the shape of the underlying power spectrum to one with a sharper break than CDM, such as mixed dark matter (e.g. Taylor \\& Rowan-Robinson 1992; van Dalen \\& Schaeffer 1992; Klypin et al. 1993). This latter possibility is the only existing model which can explain the full inflection in the APM power spectrum around $k\\simeq 0.1 \\hompc$, since CDM-like models produce a much smoother variation at this point. The conclusions we draw from our results are thus as follows: \\begin{itemize} \\begin{enumerate} \\item No linearly-biased CDM-like model can reproduce the APM power spectrum: a non-linear $b(k)$ curve is required if the cosmological fluctuations spectrum is one of the $\\Gamma^*$ models given by equations (12) and (13). Analyses (such as that of PD94) which do allow a linearly-biased $\\Gamma^*$ model can do so only because they ignore clustering data on non-linear scales, which greatly weakens their power to reject models. \\item Non-linear bias prescriptions generically produce a scale-dependent bias, but it is difficult to produce a strongly curved $b(k)$ relation with Eulerian bias, so a confirmation of the APM inflection would be strong evidence for some feature in the primordial linear spectrum at that scale or for non-local bias in galaxy clustering. \\item Analytical models based on a lognormal mass distribution are useful heuristics for gauging the rough level of the large-scale bias produced by a given prescription, but they are not sufficiently accurate to replace numerical simulations. \\item No Einstein -- de Sitter CDM-like mass model can be biased, using the range of bias prescriptions studied here, to match the APM power spectrum if it is required to be {\\em COBE\\/}-normalised and account for the abundances and $M/L$ values of clusters of galaxies: if these conditions are to be satisfied using the bias models considered here, then the mass power spectrum must be flatter than CDM-like models on small scales. \\end{enumerate} \\end{itemize}" + }, + "9708/astro-ph9708177_arXiv.txt": { + "abstract": "Dwarf, irregular and infrared-luminous starburst galaxies are all known to have ``steep'' luminosity functions with faint-end behavior roughly $\\phi(L)\\propto L^{-1.8}$. This form is exactly what is expected if the luminosities of these objects fade with time as $L\\propto t^{-1.3}$, because the objects spend more time at low luminosities than high, even if they form with a wide range of initial masses. Models of young stellar populations show this fading behavior when the star formation has occured in a single, short, recent burst. Steep luminosity functions therefore do not require steep mass functions if the galaxies are powered by fading bursts. The local-galaxy H$\\alpha$ luminosity function---which is less steep than $L^{-1.8}$---is also well-fit by this mechanism, because ionizing photon flux fades much more quickly than broad-band optical luminosity. An age-luminosity relation and a wavelength-dependence of the luminosity function are both predicted. In the context of this mechanism, the slope of the luminosity function provides a constraint on the stellar initial mass function in the bursts. ", + "introduction": "While the normal field galaxy luminosity function (GLF) $\\phi(L)$ (number density per unit luminosity) or $\\phi(\\log L)$ (number density per unit log luminosity) is ``flat'' in optical bandpasses at the faint end, i.e., $\\phi(L)\\propto L^{-1}$ or $\\phi(\\log L)={\\rm constant}$ (Efstathiou, Ellis \\& Peterson 1988; Loveday et al 1992; Mobasher, Sharples \\& Ellis, 1993; Marzke et al 1994a; Lin et al 1996, 1997; Gardner et al 1997; Ratcliffe et al 1997), many studies have found that objects in which the luminosity is thought to be dominated by young stars show a ``steep'' GLF, with roughly $\\phi(L)\\propto L^{-1.8}$ at the faint end. Parameterizing $\\phi(L)\\propto L^{\\alpha}$, the $60~\\mu{\\rm m}$ GLF from the IRAS Bright Galaxy Sample appears to show $\\alpha=-1.8$ at the faint end (Soifer et al 1987; although see Saunders et al 1990) despite the fact that these same objects lie in the flat part of the optical GLF. The $60~\\mu{\\rm m}$ luminosity is thought to originate in dust heated by the radiation from young stars at ultraviolet wavelengths where dusty galaxies are optically thick. Although the faint end of the optical GLF is flat, there may also be a small ``upturn'' at the very faintest end, around absolute magnitude $M_B=-16$~mag (Marzke et al 1994a; Driver \\& Phillips 1996; Loveday 1997), which is explained by a luminosity function with $\\alpha=-1.8$, among dwarf (i.e., low-luminosity) galaxies. In the case of the CfA survey, the upturn can be explained entirely by the luminosity function of the Sm-Im galaxies (identified on the basis of morphology) which show $\\alpha=-1.87\\pm0.2$ (Marzke et al 1994b). Local galaxies spectrally classified as strongly or recently star-forming also show a steep luminosity function (Heyl et al 1997). A steep upturn at the faint end of the GLF is observed for dwarf galaxies in rich clusters with $\\alpha$ ranging from $-1.4$ to $-2.2$ (Sandage, Binggeli \\& Tammann 1985; Driver et al 1994b; Bernstein et al 1995; De~Propris et al 1995; Lobo et al 1996; Wilson et al 1997). A recent measurement of the luminosity function of dwarf galaxies or ``knots'' formed in the tidal tails of merging galaxies finds $\\alpha=-1.75\\pm0.27$ in the $R$-band for these objects (Hunsberger, Charlton \\& Zaritsky 1996). Compact ``super star clusters'' observed in the vicinity of starburst galaxies or galaxy mergers and interpreted as the progenitors of globular clusters show a luminosity function consistent with $\\alpha=-1.8$ although the numbers are small (Lutz 1991; Holtzman et al 1992; Whitmore et al 1993; Conti \\& Vacca 1994). Finally, a luminosity function with faint end behavior $\\alpha\\approx -1.8$ is often invoked as a natural explanation of faint galaxy counts and redshift distributions (Broadhurst, Ellis \\& Shanks, 1988; Eales 1993; Koo, Gronwall \\& Bruzual 1993; Driver et al 1994a; Treyer \\& Silk 1994; Metcalfe et al 1995; Smail et al 1995; Lilly et al 1995; Ellis et al 1996). In these studies, the steep GLF is largely required to account for the large numbers of faint blue galaxies, which are mainly irregulars (Glazebrook et al 1995a; Driver et al 1995; Abraham et al 1996) and thought to have luminosities dominated by young stars. The steep faint end of the GLF is usually attributed to a steep underlying galaxy mass function. A steep mass function at small halo mass $M_h$ is natural for cold and mixed dark matter models. In the Press \\& Schechter (1974) formalism, on small mass-scales $\\phi(M_h)\\,dM_h\\propto M_h^{-((9-n)/6)}\\,dM_h$, where the post-recombination power spectrum of density fluctuations has $P(k)\\propto k^n$, with $n\\rightarrow -3$ for adiabatic fluctuations on small scales. Thus $\\phi(M_h)\\,dM_h\\propto M_h^{-2}dM_h$. This has been amply verified by numerical simulation for both cold (Brainerd \\&\\ Villumsen 1992) and mixed (Ma \\&\\ Bertschinger 1994) dark matter halos. However, the ejection of gas by early generations of stars in shallow potential wells implies that the mass converted to stars rises faster than linearly with $M_h$, so for identical stellar populations, the galaxy luminosity function should be shallower than the halo mass function (see Silk \\& Wyse 1993 for a review). Furthermore, in the IR-luminous galaxy sample, the large scatter in optical-IR colors (Soifer et al 1987) and the lack of correlation between IR luminosity and galaxy mass inferred from rotation curves (Lehnert \\& Heckman 1996) suggest that the starburst GLF is not strongly tied to the host galaxy mass function. In this {\\sl Letter,} we remark that there is a natural mechanism which ensures a steep GLF among young objects: Even if the galaxy mass function (where now by ``galaxy mass'' is meant ``the mass of that part of galaxy's baryonic mass which is turned into stars'') is flat, a GLF of roughly the form $\\phi(L)\\propto L^{-1.8}$ will be observed among any population of objects whose luminosities are dominated by light from short, recent bursts of star formation with a Salpeter-like initial mass function. This is because their luminosities decrease with time in such a way that they spend more time (and are therefore more numerous) at low luminosities than at high luminosities. This kind of mechanism underlies the theoretical explanation of the $60~{\\rm \\mu m}$ GLF by Scoville \\& Soifer (1991) and a discussion of the super star cluster luminosity function by Meurer (1995). An important feature of this mechanism is that steep mass functions are not required to explain steep luminosity functions. ", + "conclusions": "" + }, + "9708/astro-ph9708226_arXiv.txt": { + "abstract": "We present $K$-band observations of the low-luminosity galaxies in the Coma cluster, which are responsible for the steep upturn in the optical luminosity function at $M_R \\sim -16$, discovered recently. The main results of this study are \\vskip 1pt \\noindent (i) The optical$-$near-infrared colours of these galaxies imply that they are dwarf spheroidal galaxies. The median $B-K$ colour for galaxies with $-19.3 < M_K < -16.3$ is 3.6 mag. \\vskip 1pt \\noindent (ii) The $K$-band luminosity function in the Coma cluster is not well constrained, because of the uncertainties due to the field-to-field variance of the background. However, within the estimated large errors, this is consistent with the $R$-band luminosity function, shifted by $\\sim3$ magnitudes. \\vskip 1pt \\noindent (iii) Many of the cluster dwarfs lie in a region of the $B-K$ vs.~$B-R$ colour-colour diagram where background galaxies are rare ($B-K < 5$; $1.2 < B-R < 1.6$). Local dwarf spheroidal galaxies lie in this region too. This suggests that a better measurement of the $K$-band cluster luminosity function can be made if the field-to-field variance of the background can be measured as a function of colour, even if it is large. \\vskip 1pt \\noindent (iv) If we assume that none of the galaxies in the region of the $B-K$ vs.~$B-R$ plane given in (iii) in our cluster fields are background, and that all the cluster galaxies with $15.5 < K < 18.5$ lie in this region of the plane, then we measure $\\alpha = -1.41_{-0.37}^{+0.34}$ for $-19.3 < M_K < -16.3$, where $\\alpha$ is the logarithmic slope of the luminosity function. The uncertainties in this number come from counting statistics. ", + "introduction": "Recent studies of the optical luminosity function (LF) of galaxies in the Coma cluster (Trentham 1997a; Bernstein et al.~1995; Secker \\& Harris 1996) have revealed a steep rise at the faint-end, with $-1.7 < \\alpha < -1.3$ for magnitudes fainter than about $M_R = -15$ (where $\\alpha$ is the logarithmic slope of the LF: $\\phi(L) \\propto L^{\\alpha}$). The optical colours and scale-lengths of these faint galaxies suggest that they are probably dwarf spheroidal (dSph, alternatively called dwarf elliptical) galaxies. These galaxies form a distinct family of objects, separate from giant ellipticals on the luminosity $-$ surface-brightness $-$ radius parameter correlations (Kormendy 1987, Binggeli 1994). They have lower surface-brightnesses as their luminosity decreases; well known examples in the Local Group are NGC 205 and Draco. The colour distribution of the galaxies which produce a steep rise in the LF is heavily peaked at $B-R = 1.3$. This is towards the blue end of the range of colours exhibited by local dSphs (Trentham 1997b), but is redder than most of the dwarf irregulars (dIrrs). We cannot distinguish between the two types of galaxies based on their morphologies because (i) they have similar scale-lengths as a function of luminoisty (Binggeli 1994), and (ii) these scale-lengths are comparable to the seeing for galaxies in the Coma, so we cannot probe fine details in the structure. Here, we extend these studies to the near-infrared wavelengths. The K-band measurements probe the old stellar populations in galaxies as opposed to the younger population probed by the optical wavebands. The aim of this study is two fold: 1). to measure the near-infared LF of galaxies in Coma and ascertain whether or not the steep rise found at optical wavelengths is seen in the $K$-band, and 2). to study the nature of galaxies which dominate the faint-end of the LF (ie. dwarfs), using their optical-infrared colours. The measurement of the near-infrared LF is not an easy task because of background contamination. At optical wavelengths, the background counts have been well characterized because CCDs cover substantial areas (e.g.~Bernstein et al.~1995). In the near-infrared wavelengths, due to the smaller format of the arrays, such detailed measurements cannot be made. A number of medium-deep and deep surveys (Gardner et al.~1993) have permitted detailed measurements of the mean number counts. However, these measurements do not constrain the distribution of the counts around this mean, from field to field as a function of angular size. We compute this distribution from a combination of optical and $K$-band observations, and describe this calculation in detail. This is an important part of this study because the field-to-field variance of the background is the dominant source of uncertainty in the LF. The positions of the cluster dwarfs on an optical-infrared colour-colour diagram allow us to identify the type of galaxy responsible for the upturn in the optical LF, with somewhat more confidence than using the optical observations alone (see above). The difference in optical-near infrared colours between dSph and dIrr galaxies is much bigger than the difference in optical colours, because the dSphs, unlike the dIrrs, have a substantial fraction of the old stellar populations. The similarity in the optical colours between the bluest dSphs and the reddest dIrrs is normally attributed to recent star formation in the blue dSphs; a plausible physical mechanism for this is given by Silk et al.~(1987). This paper is organized as follows. In Section 2, we describe our observing strategy, the observations and data reduction. Section 3 presents the background counts. The near-ir LF for galaxies in the Coma cluster is studied in Section 4. The nature of the dwarf galaxies in the Coma cluster is explored in section 5, using their optical-infrared colours. Finally, our conclusions are summarised in section 6. Throughout this paper we assume $H_0 = 75$ km s$^{-1}$ Mpc$^{-1}$ and $\\Omega_0 = 1$. ", + "conclusions": "A K-band survey of the Coma cluster has been carried out. This consists of a wide-angle shallow ($K\\sim 19$) survey with the QUIRC (UH 2.2m) and a deeper survey ($K\\sim 21$), covering a much smaller area with the IRCAM3 (UKIRT). These observations were used to construct the near-infrared LF of galaxies in the Coma cluster and to study the nature of the population dominating the faint-end of the LF. The results of this study are summarised as follows: \\begin{enumerate} \\item The K-band LF in the Coma cluster is not well constrained. However, within the estimated (large) errors, this is consistent with the R-band LF, shifted by $\\sim 3$ magnitudes. \\item The optical-infrared colours of these faint galaxies confirm that they are dwarf spheroidals. The median $B-K$ colour for galaxies with $-19.3 < M_K < -16.3$ is 3.6 mag. \\item Using the $B-K$ vs. $B-R$ colour-colour diagram, a region on this plane is identified where the background galaxies are rare ($B-K < 5$; $1.2 < B-R < 1.6$). It is proposed that a more accurate measurement of the K-band LF for clusters can be carried out if the background contamination is estimated as a function of colour. \\item Using the $B-K$ and $B-R$ colours to correct the Coma cluster data for background contamination (as in 3 above), the K-band LF is again constructed. The logarithmic slope of the LF in the range $-19.3 < M_K < -16.3$ mag. is $\\alpha = -1.41_{-0.37}^{+0.34}$. \\end{enumerate}" + }, + "9708/astro-ph9708010_arXiv.txt": { + "abstract": "The abundances of beryllium and boron have been measured in halo stars of metallicities as low as [Fe/H] =-3. The observations show that the ratios Be/Fe and B/Fe are independent of metallicity and approximately equal to their solar values over the entire range of observed metallicity. These observations are in contradiction with the predictions of simple models of beryllium and boron production by spallation in the interstellar medium of a well mixed galaxy. We propose that beryllium and boron are produced by spallation in the ejecta of type II supernovae. In our picture, protons and alpha particles are accelerated early in the supernova event and irradiate the heavy elements in the ejecta long before the ejecta mixes with the interstellar medium. We follow the propagation of the accelerated particles with a Monte-Carlo code and find that the energy per spallation reaction is about 5 GeV for a variety of initial particle spectra and ejecta compositions. Reproducing the observed Be/Fe and B/Fe ratios requires roughly $3 \\times 10^{47}$ ergs of accelerated protons and alphas. This is much less than the $10^{51}$ ergs available in a supernova explosion. ", + "introduction": " ", + "conclusions": "" + }, + "9708/astro-ph9708156_arXiv.txt": { + "abstract": "We present the catalogue of gamma-ray bursts (GRB) observed with the WATCH all-sky monitor on board the GRANAT satellite during the period December 1989 to September 1994. The cosmic origin of 95 bursts comprising the catalogue is confirmed either by their localization with WATCH or by their detection with other GRB experiments. For each burst its time history and information on its intensity in the two energy ranges 8--20 keV and 20--60 keV are presented. Most events show hardening of the energy spectrum near the burst peak. In part of the bursts an X-ray precursor or a tail is seen at 8--20~keV. We have determined the celestial positions of the sources of 47 bursts. Their localization regions (at $3\\sigma$ confidence level) are equivalent in area to circles with radii ranging from 0.2 to 1.6 deg. The burst sources appear isotropically distributed on the sky on large angular scales. ", + "introduction": "From December 1989 to September 1994 the astrophysical observatory GRANAT performed pointed observations of different celestial regions. During that period, the X-ray instrument WATCH, a part of the scientific payload of the observatory, was monitoring the whole of the sky. WATCH is uniquely capable of precisely measuring the celestial positions (the radius of the localization region is generally smaller than 1~deg at the 3$\\sigma$ confidence level) of short-lived hard X-ray sources, which include GRBs. Another feature of the instrument relevant to observations of GRBs is that its detectors are sensitive over an X-ray energy range that reaches down to $\\sim 8$~keV, the domain where the properties of GRBs are known less than at higher energies. In this paper, we present the catalogue of GRBs detected with WATCH in 1989--1994. For nearly half of the events we have been able to determine the location of the burst source on the celestial sphere. Earlier, a preliminary catalogue covering the WATCH observations carried out before October 1992 was compiled by \\cite{castro-tirado94}. The new catalogue has been updated mainly in the two aspects: 1) the bursts detected between October 1992 and September 1994 have been added, 2) more accurate positions have been determined for many of the previously catalogued events due to the use of more precise information on the spacecraft attitude and an improved model of the instrument. ", + "conclusions": "The WATCH GRB catalogue provides a set of 47 burster positions, 39 of which have a total (statistical plus systematic) uncertainty of less than 1~deg. For 13 events error boxes with radii smaller than 30~arcmin are now available. The WATCH sample thus contributes to the currently available list of moderately accurate GRB positions a number of locations comparable to that accumulated by the interplanetary satellite networks and never before obtained with any stand-alone instrument. We therefore hope that the new data presented here will be useful for burster counterpart searches in different energy ranges as well as for studying possible correlations in GRB positions. The WATCH GRBs appear distributed both isotropically on the celestial sphere and homogeneously in space. These two results seem to be consistent with the implications from the third BATSE catalogue, as WATCH is about an order of magnitude less sensitive than the large-area detectors of BATSE, and the brighter bursts in the BATSE catalogue also show both isotropy and homogeneity (\\cite{meegan96}). The light curves of most bursts observed by WATCH show hardening of the energy spectrum near the burst maxima. Several bursts demonstrate a significant 8--20~keV activity in the absence of hard X-ray flux either before or after the GRB. To all appearances, these X-ray events accompanying gamma-ray bursts are higher energy manifestations of the soft X-ray precursors and tails observed at 1.5--10~keV by GINGA." + }, + "9708/gr-qc9708043_arXiv.txt": { + "abstract": "{\\bf Abstract.} This paper and the others in the series \\cite{bi:EBD96,bi:SEB96} challenge the standard model of the effects of gravitational lensing on observations at large distances. We show that due to the cumulative effect of lensing, areas corresponding to an observed solid angle can be quite different than would be estimated from the corresponding Friedmann-Lema\\^{\\i}tre model, even when averaged over large angular scales. This paper concentrates on the specific example of spherically symmetric but spatially inhomogeneous dust universes, the Lema\\^{\\i}tre-Tolman-Bondi models, and shows that radial lensing significantly distorts the area distance-redshift and density-redshift relations in these exact solutions compared with the standard ones for Friedmann-Lema\\^{\\i}tre models. Thus inhomogeneity may introduce significant errors into distance estimates based on the standard {\\sc fl} relations, even after all-sky averaging. In addition a useful new gauge choice is presented for these models, solving the problem of locating the past null cone exactly. ", + "introduction": "Our aim in this series of papers is to show that when analyzing observations at large redshift in the real universe, the assumption of an area distance corresponding to that of a best-fit Friedmann-Lema\\^{\\i}tre ({\\sc fl}) model\\footnote {We follow recent moves to use `Friedmann-Lema\\^{\\i}tre' to describe the dynamics of and the application of the Einstein field equations to the standard model, and `Robertson-Walker' to describe its metric and geometry.} --- that is, one with a matching averaged matter density --- may not be a good approximation, even when averaging over large angular scales. This claim is important because of the ubiquitous use of {\\sc fl} models in studies of number counts versus redshift and area distance versus redshift. The first paper in the series (paper I) \\cite{bi:EBD96} gave general arguments for this thesis. It was explained there that whilst the Dyer-Roeder distance \\cite{bi:DR72} is generally regarded as a good approximation for ray bundles moving between high-density clusters of matter, resulting in a de-focussing relative to the comparable {\\sc fl} model, matter moving near or through higher density regions is more focused than in the {\\sc fl} model, so resulting in a compensating effect. It is commonly believed \\cite{bi:W76,bi:SEF92} that an exact cancellation between the two effects takes place when one averages over large angular scales containing both high and low density regions --- so the {\\sc fl} area distance is the correct one on these scales. However, it was pointed out in \\cite{bi:EBD96} that, after caustics have formed through the focussing of light rays by the high density regions, these light rays too are rapidly diverging so that sufficiently far down the past light cone {\\it all} light rays will be diverging relative to the comparable {\\sc fl} model. Consequently there is good reason to question the general opinion in this matter --- we find that in fact `shrinking' takes place.\\footnote {The term shrinking \\cite{bi:L88} describes the following. If at some redshift the area distance is greater in the inhomogeneous universe than in the best-fit or background {\\sc fl} model, scales at that redshift will be underestimated if the {\\sc fl} area distance is used to calculate the size of objects from their measured angular size, and so will have `shrunk' relative to their actual scales in the more realistic model. Put another way, if a sphere at given $z$ about the observer has a larger total area in the inhomogeneous model, an object of fixed size will subtend a smaller solid angle on the sky, and so be shrunk. (See \\cite{bi:EBD96} for more discussion.)} It is often argued that photon conservation excludes this possibility, but these arguments are based on incorrect assumptions about the geometry. Given this feature, even when there is only weak lensing,\\footnote {We refer to `lensing' when the light rays are different than they would be in the background {\\sc fl} model. Thus weak lensing implies a change of apparent positions and alteration in the usual distance relations, usually combined with image distortions, but not necessarily multiple imaging.} and caustics have not formed, it is not obvious that the {\\sc fl} area distance-redshift relation is an accurate description. The aim of this paper is to construct an exact inhomogeneous model and its {\\sc fl} approximation, and compare the area distance-redshift and density-redshift relations in the two. To do so we examine Lema\\^{\\i}tre-Tolman-Bondi ({\\sc ltb}) spherically symmetric dust solutions \\cite{bi:L33,bi:T34,bi:B47}, where exact integrations of the field equations are available for the past light cone of observers at the central position. Although the lensing that occurs for this central observer is purely radial, we find the inhomogeneity has a tangible effect on observational relations.\\footnote {Radial lensing is a spherically symmetric distortion of the null cone compared with an {\\sc fl} model, resulting in a uniform delay of the wavefront. There is no image distortion, no dependence of magnification or time delay on direction, and no multiple imaging, but our results show its effects are observable.} In the particular case examined, because of the spherical symmetry about the observer, the effect in any specific direction will not be compensated by an opposite effect in another direction --- on the contrary, the effect is uniform because it is the same in all directions, and will occur on large as well as small angular scales. However unlike the previous paper \\cite{bi:EBD96}, this will not be associated with the formation of caustics. The failure of the {\\sc fl}-like area distance-redshift assumption will thus have been shown to occur for observers at a very special position in a family of high-symmetry space-times. Undoubtedly this is not the kind of situation the authors of the papers mentioned above, claiming the effect does not occur, had in mind. However the usual statements of this effect contain no clauses that exclude this situation: the result is supposed to hold in all cases when the `lumpy' universe is reasonably close to an {\\sc fl} model, and observations are averaged over the sky; their arguments do not exclude the situation envisaged here where the observer is at the centre of a spherically symmetric inhomogeneity. Indeed, a spherically symmetric model may be regarded as describing data that has been averaged over the {\\em whole} sky, but not over distance. Our example thus confirms the claims of paper I, in the setting of particular exact inhomogeneous solutions of the Einstein Field Equations. It does not generically establish the magnitude of the effect, precisely because the high-symmetry geometry considered here precludes formation of caustics and the consequent fractal-like structure of the real light cone. Paper III \\cite{bi:SEB96} will provide confirmation of the overall shrinking effect due to caustics, and attempt more realistic estimates of its magnitude than those given in paper I, which used simple analytic formulae for this purpose. In developing the results of this paper, we solve one of the problems that has made analysis of observations in Lema\\^{\\i}tre-Tolman-Bondi solutions difficult, namely the problem of precisely locating the past light cone of the chosen central event P, by use of a special choice of radial coordinate that ensures a very simple form for the past light cone of P in these inhomogeneous space-times. This technical development has other uses in terms of analysing observational relations in these models.\\footnote {For a slightly different analysis of {\\sc ltb} spacetimes, based on null cone coordinates, see \\cite{bi:MHMS95}, and for a consideration of observations away from the centre of symmetry see \\cite{bi:HMM96}.} Our approach is complementary to that of Kurki-Suonio and Liang \\cite{bi:K-SL92}, who did numerical calculations of observational relations in a hyperbolic {\\sc ltb} model derived from an $\\Omega = 0.1$ {\\sc fl} model plus some overdensities. ", + "conclusions": "\\label{sec:conc} The general belief that photon conservation implies that the total area of an incoming wavefront must be the same as in the background, matter-averaged, {\\sc fl} model is contradicted by our results. The functions $\\hat{R}(z)$ and $\\hat{\\rho}(z)$ in an inhomogeneous model differ from their standard {\\sc fl} forms, because the redshift $z$ is path dependent and the null cone is warped. Thus, even if the inhomogeneous model is exactly homogeneous beyond some radius, the standard {\\sc fl} forms are not recovered. This means that a universe which is {\\sc fl} {\\em on average} will not in general present standard {\\sc fl} forms of $\\hat{R}(z)$ and $\\hat{\\rho}(z)$ to an observer. The spherically symmetric model used here is simple but effective, since averaging over direction cannot change the results. In more realistic models of the lumpy universe this effect will still be present, and we expect full gravitational lensing to occur, resulting in more significant deviations from the {\\sc fl} formula. This investigation used a parabolic {\\sc ltb} model, where the areal radius $\\hat{R}$ is also the area distance of the 2-sphere wavefronts of the past null cone. The density in the {\\sc ltb} model is averaged to give a background Einstein-de Sitter ($\\Omega = 1$) model, and it is tested against this model. Although there exists no covariant way to perform this averaging, we use the `natural' one defined by the use of junction conditions, here equivalent to the one used in astrophysical problems: that is, averaging by integrating the rest mass and proper volume on constant time slices. More importantly, the functions $\\hat{R}(z)$ and $\\hat{\\rho}(z)$ in the examples are obviously perturbed away from the standard {\\sc fl} ones, so that no one {\\sc fl} model can give the same $\\hat{R}$ and $\\hat{\\rho}$ values as an inhomogeneous one at a variety of $z$ values. The results show that it is quite easy to have areas in the inhomogeneous models which differ significantly from areas in the background, matter-averaged {\\sc fl} model. The result may either be shrinking (the background {\\sc fl} area distance underestimates the real area distance at that redshift) or magnification (background {\\sc fl} area distance is an overestimate). The presence of loops in the $\\hat{R}$-$z$ and $\\hat{\\rho}$-$z$ graphs is analogous to the well known `finger of God' effect familiar in redshift maps of the galaxy distribution. Our results and conclusions generally agree with those of Kurki-Suonio and Liang \\cite{bi:K-SL92}, who calculated obervational relations numerically in 4 hyperbolic {\\sc ltb} models out to $z = 0.5$. They generated mild and strong deviations from the {\\sc fl} observational relations, with `bang time' inhomogeneities and `areal radius' inhomogeneities having opposite effects on the oberved matter distribution in redshift space, compared with the present day distribution. Their method did not permit the redshift to be disordered with distance. Whilst the major aim of this paper has been the above thesis, the choice of radial coordinate which locates the observer's null cone will be of use in future analyses of observations in these isotropic dust models. An important caveat is that since the {\\sc ltb} model does not allow for formation of caustics in the null cone of the central observer, it cannot be considered a useful model for obtaining quantitative `real world' results. Rather this paper should be viewed as a proof that even purely radial lensing distorts the area distance-redshift relation significantly. If the observer moves away from the central position, then continuity ensures that the radial effects found here will still be present, and the effects of true lensing will be superimposed. As argued in \\cite{bi:EBD96,bi:SEB96}, we expect caustics to skew the area towards larger values, so that most objects in the universe are demagnified. The importance of all this is that it opens up the way for considering the effects of lensing by inhomogeneities on large angular-scale number counts and area distances as opposed to limiting discussion to lensing effects on small scales." + }, + "9708/astro-ph9708232_arXiv.txt": { + "abstract": "We present deep ($V \\simeq 27$) \\V- and \\I-band stellar photometry of G302 and G312, two globular star clusters in the halo of M31. These data were obtained using the {\\sl Hubble Space Telescope\\/}'s Wide Field/Planetary Camera 2. We find iron abundances of $\\FeH = -1.85 \\pm 0.12$ for G302 and $\\FeH = -0.56 \\pm 0.03$ for G312, consistent with spectroscopic measurements. The color--magnitude diagrams for each cluster show no evidence for an intermediate-aged population of stars, or a second parameter effect in the morphology of the horizontal branch. G302 shows no evidence for a color gradient but the inner regions of G312 are bluer than the outer regions. G312 shows no evidence of ellipticity or an extended halo of unbound stars. G302 has a projected ellipticity of $\\epsilon = 0.195 \\pm 0.012$ with the projected major axis oriented towards the center of M31. G302 also shows evidence of an extended asymmetric stellar halo extending to at least twice the fitted Michie--King tidal radius. The amount of mass beyond the tidal radius of G302 is consistent with the stellar escape rates which have been predicted by $N$-body simulations of globular clusters in the Galactic tidal field. ", + "introduction": "} Globular star clusters (GCs) are pressure-supported collections of between $\\sim 10^4$ and $\\sim 10^6$ stars which are usually associated with galaxies, although there is evidence that some clusters of galaxies contain a population of ``free'' GCs which are associated with the cluster potential as a whole and not any individual galaxy (West \\etal \\markcite{WC95}1995). Individual Galactic GCs are made up of stars with a single overall chemical abundance suggesting that they formed in a single star formation event. GCs typically have integrated magnitudes of $-10 \\le V \\le -4$ making GC systems visible out to redshifts of $z \\sim 0.04$, the approximate distance to the Great Wall galaxies. The shape of the GC luminosity function (LF) has been assumed to be universal, so the GC LF has been used as a distance indicator (see Harris \\markcite{H91}1991 for a review). However, recent work has suggested that the shape of the GC LF may depend on the metallicity of the GC system (Ashman \\etal \\markcite{AC95}1995) and the details of stellar evaporation from individual GCs (Okazaki \\& Tosa \\etal \\markcite{OT95}1995). Because of the wide-spread use of GCs to determine distances to external galaxies, it is important to determine if GCs truly are the same from one galaxy to the next. This is best determined by studying the physical structures and stellar populations of GCs in nearby galaxies. The nearest large GC system outside the Milky Way Galaxy is that of the Andromeda Galaxy ($=$ M31 $=$ NGC 224). M31 is located at a distance of 725 kpc ($\\mu_0 = 24.3$, van den Bergh \\markcite{V91}1991) so individual stars in M31's GCs can be easily resolved with the {\\sl Hubble Space Telescope\\/} ({\\sl HST\\/}) and large ground-based telescopes at sites with sub-arcsecond seeing. M31's low inclination ($i = 12\\fdg5$, Hodge \\markcite{H92}1992) means that many of its GCs are not superimposed against the disk of M31, making identification of GCs, and photometry of their stars, relatively straightforward. M31 offers a unique laboratory to study the outer regions of GCs. Because of their small fields-of-view it is not practical to used CCDs to obtain star counts in the outer regions of Galactic GCs. Beyond distances of approximately half the tidal radius, $r_t$, the projected stellar densities of the clusters are overwhelmed by random fluctuations in the background stellar number density (Innanen \\etal \\markcite{IH83}1983). However, M31's GCs have sufficiently small angular sizes ($\\theta \\sim 5\\arcsec$ to $15\\arcsec$) that both the cluster and the background can be imaged on a single large format CCD image. This eliminates the need to match photometric zero-points between the cluster and the background, which makes possible a more precise subtraction of the background light from the cluster light. The first study of the internal structures of M31 GCs was undertaken by Battistini \\etal \\markcite{BB82}(1982), who estimated core radii for several clusters. Pritchet \\& van den Bergh \\markcite{PV84}(1984) found that the surface brightness profile for G1 ($=$ Mayall II; the G-numbers used in this paper are from Sargent \\etal \\markcite{SK77}[1977]) had an excess of light at large radii compared to the best-fitting seeing-convolved analytical King \\markcite{K66}(1966) model. They did find that G1 was well fit by empirical King \\markcite{K62}(1962) models with core radii of $r_c \\lesssim 0.5$ pc. Crampton \\etal \\markcite{CC85}(1985) used seeing-convolved King \\markcite{K62}(1962) models to derive core radii for nearly 500 M31 GCs. Their values, however, are systematically $\\sim$50\\% larger than those of Battistini \\etal \\markcite{BB82}(1982), despite the better seeing conditions of the Crampton \\etal \\markcite{CC85}(1985) data set. Bendinelli \\etal \\markcite{BP90}(1990) used ground-based data to produce seeing-deconvolved radial profiles for six bright GCs in M31, but seeing and pixel scale limitations restricted them to resolutions of $\\sim$0\\farcs3, insufficient to resolve the cores of the clusters. Still, their data suggest that M31 GCs have King-like profiles similar to those of Galactic GCs. The core structures of some of M31's GCs have been studied using the pre-refurbished {\\sl HST}. Bendinelli \\etal \\markcite{BC93}(1993) detected a power-law density cusp in G105 using {\\sl HST\\/}'s Faint Object Camera (FOC) images and a variety of image restoration and seeing deconvolution techniques. In addition, Fusi Pecci \\etal \\markcite{FB94}(1994) used similar methods to obtain half-width at half-maxima (HWHM) and half-light radii ($r_h$) for thirteen M31 GCs from FOC images. They found HWHMs similar to the core radii of Galactic GCs. Their data, however, could not be used to find the tidal radii of these clusters because the point-spread function (PSF) before the refurbishment of the {\\sl HST\\/} overfilled the FOC's field of view. Cohen \\& Freeman \\markcite{CF91}(1991) derived tidal radii for thirty M31 GCs by fitting seeing-convolved King models. Although their fits to individual clusters were quite uncertain they did find a mean tidal radius for the M31 clusters---after adjustment for differences in galactic masses and rotation velocities---which was very similar to that of the Milky Way clusters. GCs do not exist in isolation but sit in the tidal field of a galaxy. Any stars that evaporate from a GC by the galaxy will have velocity vectors similar to the velocity vector of the GC\\@. This can result in the GC being surrounded by an extended halo of unbounded stars which move in approximately the same direction as the GC and have approximately the same velocity. This idea has been explored numerically by Oh \\& Lin \\markcite{OL92}(1992) who predicted that GCs could be surrounded by extended halos of escaped stars which can persist for up to a Hubble time. The size and shape of such a halo can be influenced by the tidal field of the parent galaxy. Evidence for extended halos has been observed in some Galactic GCs by Grillmair \\etal \\markcite{GF95}(1995). In addition, Grillmair \\etal \\markcite{GA96}(1996, hereafter referred to as GAF) have observed an excess of resolved and unresolved stars beyond the formal Michie--King (MK) tidal radii of several GCs in M31, as would be expected if extended halos were present. There has been some interest in determining the ellipticities of M31 GCs. Pritchet \\& van den Bergh \\markcite{PV84}(1984) measured an ellipticity of $\\epsilon = 0.22$ for the region of G1 with $12\\arcsec \\lesssim r \\lesssim 35\\arcsec$. Spassova \\etal \\markcite{S88}(1988) measured ellipticities for approximately two dozen GCs while a study by Lupton \\markcite{L89}(1989) suggests that the mean ellipticity measured in the inner 7 to 14 pc of a M31 GC ($\\overline{\\epsilon} = 0.08$) is indistinguishable from the mean ellipticity of Galactic GCs. In the outer 14 to 21 pc, however, the mean ellipticity of an M31 cluster is $0.11 \\pm 0.08$, greater than that of the Galactic GCs ($\\overline{\\epsilon} = 0.08 \\pm 0.07$) and similar to GCs in the Large Magellanic Cloud ($\\overline{\\epsilon} = 0.11 \\pm 0.07$). Baev \\etal \\markcite{BS97}(1997) found systematic differences between the shapes of M31's disk and halo GCs. They found that the disk GCs are triaxial ellipsoids while the halo GCs are oblate or prolate spheroids, but cautioned that this is a preliminary result since the sample of halo GCs in their study is small compared to the sample of disk GCs. The first color--magnitude diagrams (CMDs) for GCs in M31 were for G1 by Heasley \\etal \\markcite{HC88}(1988) and G219 ($=$ Mayall IV) by Christian \\& Heasley \\markcite{CH91}(1991). Couture \\etal \\markcite{CR95}(1995) undertook a systematic study of five M31 GCs (G11, G319, G323, G327 $=$ Mayall VI, and G352 $=$ Mayall V) with a range of iron abundances. Unfortunately none of these ground-based studies were able to reach the level of the horizontal branch (HB) at $V \\simeq 25$. The first CMDs constructed from {\\sl HST\\/} data were for G1 (Rich \\etal \\markcite{RM96}1996); G58, G105, G108, and G219 (Ajhar \\etal \\markcite{AG96}1996); and G280, G351, and Bo468 (Fusi Pecci \\etal \\markcite{FP96}1996). These CMDs were able to resolve stars one to two magnitudes below the red portion of the HB\\@. In order to study the stellar populations and structures of GCs in M31 we obtained deep {\\sl HST\\/} Wide-Field/Planetary Camera-2 (WFPC2) images of two halo GCs (G302 and G312) located $32\\farcs1$ and $49\\farcs8$ respectively from the center of M31 along the southeast minor axis. Our data is unique in that we have centered the GCs on the WF3 CCD in order to take advantage of the WFC's field of view to search for extended halos around each GC\\@. ", + "conclusions": "} We have used {\\sl HST\\/} WFPC2 photometry to construct deep ($V \\simeq 27$) CMDs for two GCs in the halo of M31. Both GCs appear to have a single old population of stars similar to what is found in Galactic GCs. The shape of the RGB for G302 gives an iron abundance of $\\FeH = -1.85 \\pm 0.12$, in agreement with the published values obtained using spectroscopy. For G312 we obtain $\\FeH = -0.56 \\pm 0.03$, which is somewhat more metal-rich than the spectroscopic value. Neither GC shows any indication that there is a second parameter acting upon their HB morphologies. Both GCs have MK tidal radii of $r_t \\simeq 10\\arcsec$, core radii of $r_c \\simeq 0\\farcs2$, central concentrations of $c \\simeq 1.7$, and half-mass radii of $r_h \\simeq 0\\farcs5$. There is no evidence for velocity anisotropy in either G302 or G312. G302 has a color of $(V-I)_0 = 0.83$ while the color of G312 is $(V-I)_0 = 1.07$. G302 has a projected ellipticity of $\\epsilon = 0.195$ with the major axis oriented approximately towards the center of M31. This GC has an excess of light beyond its formal tidal radius which is not consistent with either an isotropic or an anisotropic MK model. The two-dimensional distribution of stars around G302 is consistent with the presence of an extended halo extending to two to three times the formal tidal radius from the cluster. G312, on the other hand, has an ellipticity of $\\epsilon \\simeq 0$. Neither the integrated light, nor the star counts, show any evidence for an extended halo. It is possible that such a tail does exist for G312, but is oriented along the line of sight. We have estimated the projected mass-loss rate from G302 to be $\\dot {\\cal M} = 4500 \\pm 1800 \\, {\\cal M}_{\\sun}$ per Gyr which corresponds to a projected escape rate of $\\dot r = (2.3 \\pm 0.9) \\times 10^{-3}$ per half-mass relaxation time. The projected escape rate from G312 is $\\dot r = (0.38 \\pm 2.95) \\times 10^{-3}$ per half-mass relaxation time. These are consistent with the escape rates predicted by Oh \\& Lin \\markcite{OL92}(1992) although the large photometric uncertainties near the tidal radius of G312 makes the escape rate for this GC much less reliable than that for G302." + }, + "9708/astro-ph9708138_arXiv.txt": { + "abstract": "We have made spectroscopic identifications of 39 additional quasar candidates from the Deep Multicolor Survey (DMS) of Hall {\\it et al.} (1996, ApJ, 462, 614). We have identified 9 new quasars with $0.3 < z < 2.8$ and $16.8 < B < 21.6$, all from the group of candidates with ultraviolet excess (UVX). No new quasars with $z > 3$ were found among the observed candidates selected due to their red ($B-R$) and ($V-R$) colors. As a result, there are now 55 confirmed quasars in the survey: 42 with $0.3 < z < 2$, nine with $2 < z < 3$, three with $3 < z < 4$, and 1 at $z = 4.3$. One new quasar, DMS 0059$-$0055, is very bright with $B=16.8$ and $z=0.3$, making its detection by our survey very unexpected. Including this new spectroscopy, the results of the DMS are converging with the predicted space densities of other surveys. In particular, we no longer find an excess of quasars with $z<2.3$ and $B<21$ in the survey over predictions based on models by Koo \\& Kron. Also, the excess in the number of quasars seen at $z>3$ over predictions based on models by Warren, Hewett, \\& Osmer is less than previously suggested. We also demonstrate the success of our quasar color modeling which is important in assessing the completeness of our survey. ", + "introduction": "This is the third in a series of papers describing results from the Deep Multicolor CCD Survey (DMS) of \\cite{HALL2} (1996ab; Papers I and II) conducted at the KPNO 4m telescope. The survey covered 0.83 deg$^2$ using six filters (U,B,V,R,I75,I86) covering the range from 0.34 to 0.86$\\mu$m in six fields at high galactic latitude. The imaging data have average 5$\\sigma$ limiting detection magnitudes ranging from 22.1 to 23.8. The motivation of the survey is to search for lower luminosity quasars at $z > 3$ than have previously been studied and to search for lower luminosity quasars with $z \\approx 2$ to constrain the nature of the evolution of the luminosity function. In addition, the survey is valuable for the study of faint field galaxies and stars in the galactic halo ({\\it e.g.}, \\cite{LIU}.) Paper I contains the details of the construction of the stellar catalog of 21375 objects, including a detailed description of the imaging observations and reductions, object classification, photometry, astrometry, and the catalog completeness and contamination. Paper II describes the search for quasars including candidate selection criteria, survey completeness estimates, and initial spectroscopy which resulted in the discovery of 46 quasars, including one at $z=4.3$, and a comparable number of compact narrow emission line galaxies (CNELG's). The initial spectroscopy described in Paper II was used along with determinations of the survey efficiency to estimate the number of quasars contained in the survey for different redshift and magnitude bins. Our expected numbers were compared with predictions based on models from \\cite{KK88} (1988, hereafter KK88) and \\cite{WHO} (1994, hereafter WHO). It was concluded that the DMS contains more quasars with $B < 21$ and $z<2.3$ than expected from the results of Koo \\& Kron and more quasars at $z>3$ than expected from the results of Warren, Hewett, \\& Osmer. However, the initial spectroscopy included less than half of the candidates at the brightest magnitudes and even fewer of the fainter candidates. Here we report additional follow-up spectroscopy obtained at the KPNO 4m telescope in April and October 1995 resulting in the identification of 39 additional candidates, including the discovery of 9 additional quasars. In \\S2 we describe the candidate selection process, in \\S3 we describe the spectroscopic observations and reductions, and in \\S4 we discuss the spectroscopic results. In \\S5, we compare our findings with predictions from the determinations of quasar densities from other surveys for quasars. \\S6 contains further discussion. ", + "conclusions": "Multicolor surveys can suffer from large selection biases that vary with redshift, magnitude, and spectral energy distribution (SED.) We have attempted to model these biases by producing synthetic quasar spectra with a variety of properties (see Paper II, \\S5 for a detailed description of the synthetic spectra) and determining the probability of selecting such quasars given our candidate selection procedures. The question then arises: are the models actually representative of the data? Figure 2 demonstrates that the theoretical quasar colors computed from the synthetic spectra do accurately model the colors of the quasars contained in the DMS. Shown is a ($B-V$) {\\it vs.} ($U-B$) color-color diagram containing the 27 UVX quasars in the range ($16.5 < B < 21.0$) along with their redshifts. Theoretical quasar colors computed from the synthetic spectra are plotted as tracks. Four different quasar SED are shown with varying line strengths and spectral slopes ($\\alpha$) and with redshifts ranging from $z=0.05$ to 3.05 in steps of 0.1. The theoretical tracks do bracket the survey quasar colors and the trend in color with redshift is also matched quite well. In particular, note the two quasars at $z\\approx2.8$ at $(B-V)\\approx0.3$ and $(U-B)\\approx0$. They would have been predicted to be in the redshift range between $z=2.4$ and 2.9 based on their colors alone. Since the synthetic quasar spectra play a major role in computing the survey completeness, the good agreement between their colors and those of the real quasars lends more confidence in our results. The discovery of the one bright quasar at $z=0.296$ and $B=16.8$ and the quasar at $z=4.3$ and $R=20.1$ are unexpected in 0.83 deg$^2$, the area of the DMS survey. We could expect to find 0.04 quasars at $z<2.2$ and $16.5 < B < 17.0$ in the DMS based on the QLF of \\cite{BSP} (1988), or one every 28 deg$^2$. The Large Bright Quasar survey of \\cite{HEW} found 15 quasars with $16.5 < B < 17.0$ in an effective area of 454 deg$^2$, or one every 30 deg$^2$. The Hamburg/ESO Survey (\\cite{HES}) and the Homogeneous Bright QSO Survey (\\cite{LFC}) also predict one such quasar every $\\sim30$ deg$^2$. Given the agreement of these surveys at these magnitudes, it was indeed unlikely that this bright object would be found in the DMS. As for the $z=4.3$ quasar, we could expect to find 0.06 quasars at $4.0 < z < 4.5$ and $19.5 < R < 20.5$ in the DMS based on the results of WHO or 0.28 based on the results of KDC, or one quasar every 13 or 3 deg$^2$, respectively. While these results are unexpected, due to their small numbers, we do not find their discovery statistically significant. The spectroscopy reported here brings the total number of quasars in the DMS to 55, four of these having $z>3$. Given our observational efficiencies including this new spectroscopy, we conclude that the number of quasars estimated to be contained in the DMS is in good agreement with the results of other surveys for quasars. The absence of new quasars with $z > 3$ among the BRX candidates reduces the expected total number of such quasars in the survey from 8.4 to 4.0. This is not a significant excess over the predicted number of 1.2 (allowing for efficiency of detection) based on results by WHO or of 2.0 based on results by KDC, and, therefore, the estimated number of objects in the DMS at $z>3$ are in good agreement with these two surveys. The results for the UVX candidates indicate that the number of quasars in the survey with $z < 2.3$ and $B < 22$ is also less than estimated in Paper II and in good agreement with the predictions of the KK88 and BSP surveys. Unfortunately, we were unable to identify any additional UVX candidates at $B>22.0$ beyond what was reported in Paper II. This spectroscopy along with that of the remaining BRX candidates will require an instrument/telescope combination with greater sensitivity than the configurations used here. The next steps are to complete the spectroscopy of the BRX candidates so that a definitive value of the surface density of faint $z > 3$ quasars in the survey is established; complete the spectroscopy of the UVX candidates with $B > 22$ to constrain the evolution of the faint end of the luminosity function at $z < 2.3$; make the entire catalog available to the community in electronic form; and identify and analyze the field galaxies and faint stars." + }, + "9708/hep-ph9708281_arXiv.txt": { + "abstract": "\\setlength{\\baselineskip}{14pt} A particular class of variant axion models with two higgs doublets and a singlet is studied. In these models the axion couples either to the $u$-quark or $t$-quark or both, but not to $b$, $c$, $s$, or $d$. When the axion couples to only one quark the models possess the desirable feature of having no domain wall problem, which makes them viable candidates for a cosmological axion string scenario. We calculate the axion couplings to leptons, photons and nucleons, and the astrophysical constraints on the axion decay constant $v_a$ are investigated and compared to the DFSZ axion model. We find that the most restrictive lower bound on $v_a$, that from SN1987a, is lowered by up to a factor of about 35, depending on the model and also the ratio of the vacuum expectation values of the higgs doublets. For scenarios with axionic strings, the allowed window for $v_a$ in the $u$ quark model can be more than two orders of magnitude. For inflationary scenarios, the cosmological upper bound on $v_a/N$, where $N$ is the QCD anomaly factor, is unaffected: however, the variant models have $N$ either 3 or 6 times smaller than the DFSZ model. ", + "introduction": "The relevance of instantons \\cite{1} to physics was first shown by 't Hooft \\cite{2} in his resolution of the $U_A(1)$ problem \\cite{3}. It was first pointed out by him that the topological structure of the vacuum of any non-abelian gauge theory, which includes QCD, is non-trivial. There are gauge transformations, characterized by different topological numbers $n$ (the Pontryagin index), which can not be continuously deformed to one another. This gives rise to distinct ground states, labelled by different $n$ and separated by finite energy barriers. Instantons can be physically interpreted as quantum mechanical tunneling events, in Euclidean spacetime, between these different ground states \\cite{4,5}. One then has to construct the true, gauge-invariant, vacuum out of these degenerate $n$-vacua by taking a linear combination \\begin{equation} |\\theta\\rangle=\\sum_n e^{-in\\theta}|n\\rangle. \\end{equation} In the path integral approach, the effect is to add the so-called $\\theta$-term, to the ordinary QCD lagrangian, or \\begin{equation} \\cL_\\theta=\\theta{g_{s}^2\\over 32\\pi^2}G_{b}^{\\mu\\nu}{\\tilde{G}_{b\\mu\\nu}}. \\end{equation} Under the combined action of charge conjugation and parity transformation the $\\theta$-term changes sign and hence it violates CP invariance. There is also CP violation communicated from the quark mass matrix $M$: if we diagonalise it with a bi-unitary transformation, we find that the coupling constant $\\theta$ is modified to \\begin{equation} \\bar\\theta=\\theta+\\arg(\\det M). \\end{equation} Despite being the coupling constant of a total derivative term, the parameter $\\bar\\theta$ is observable, through its effect on the neutron electric dipole moment \\cite{6}. The current experimental upper limit \\cite{7} on the electric dipole moment implies that $\\bar\\theta$ is less than about $10^{-9}$. Such a small value of $\\bar\\theta$ contradicts our expectation that a dimensionless free parameter should be of order one. Many ideas have emerged in trying to resolve the puzzle. One of the most elegant solutions was proposed by Peccei and Quinn in 1977 \\cite{8}. The authors postulated the invariance of the lagrangian under the transformations of a new extra global chiral $U(1)$ symmetry, called PQ-symmetry, thus enlarging the symmetry group of the SM to $SU(3)_C\\times SU(2)_{L}\\times U(1)_Y\\times U(1)_{PQ}$. To accomodate the extra charges of the new symmetry one needs (at least) one extra higgs doublet. When this symmetry is spontaneously broken, a pseudoscalar boson appears \\cite{9} in the theory, called the {\\it axion}. Normally, one would think that the axion is massless, as it is the Nambu-Goldstone boson of the PQ symetry. However, the $U(1)_{PQ}$ symmetry is an anomalous one, spoiled by the effect of instantons in the QCD vacuum, and the axion picks up a small mass via the axion-gluon-gluon triangle anomaly. The assignment of appropriate PQ-charges to the higgs fields and consequently to the quarks is responsible for the presence of the anomaly in the PQ current, and also for the variety of the axion models. In the original model, known as the PQWW (Peccei-Quinn-Weinberg-Wilczek) axion model \\cite{9}, all the quarks of the same chirality were assigned the same PQ charge. Unfortunately the model was ruled out both by particle physics experiments and by astrophysical observations. The former gave constraints which came from $K,\\ J/\\psi$ and $\\Upsilon$ meson decays, reactor and beam dump experiments and nuclear deexcitations \\cite{10}. The latter are in fact more restrictive and imply that $v_a>10^{10}$ GeV for most axion models considered to date \\cite{11}. A way out of this was first proposed by Kim in 1979 and subsequently by Shifman, Vainstein \\& Zakharov \\cite{12}, who added a higgs singlet, thus enabling the axion decay constant to be much higher than the electroweak scale. However, it cannot be too high, as it is also possible to restrict $v_a$ from above through cosmological arguments. Coherent oscillations in axion field can be produced after inflation or via the formation of axionic cosmic strings. The requirement that the energy density in these oscillations is not large enough to overclose the universe puts an upper limit on $v_a$. The current values on these limits for the axion decay constant are $v_a$ less than about $10^{12}$GeV \\cite{13,11} from inflationary scenarios and $v_a\\lesssim 2.6\\times 10^{11}$GeV \\cite{14} from axion strings (for $H_0 = 50$ km s$^{-1}$ Mpc$^{-1}$, and making the conservative assumption that radiation from infinite strings dominates that from loops). As we can see there is only a very small window left for the axion. Cosmology also restricts on the value of $N$, the parameter that characterises the QCD anomaly, which is related to the number of quarks that couple to the axion. If there is no inflation between the PQ symmetry-breaking transition and the present day, a dense network of axion strings is formed. At around a temperature of 1 GeV, each string becomes the junction of (in our normalisation convention) $2N$ domain walls \\cite{15}. In order to avoid the domain walls dominating the energy density, $2N$ must be equal to unity, so that the string-wall system can annihilate. Models with $N>{1\\over 2}$ must have a period of inflation at a low energy scale, or must reheat after inflation to less than the PQ symmetry-breaking temperature, to remain viable. In this paper we examine some variant axion models based on those proposed first by Peccei, Wu \\& Yanagida \\cite{16} and independently by Krauss \\& Wilczek \\cite{17}. Their models were constructed with two higgs doublets and assigned of different PQ charges to different quarks. The original reason for this is that in order to make an axion model which avoided the particle physics constraints at the time it was essential to decrease the axion couplings to $c$ {\\it and} $b$ quarks on one hand, and on the other to sufficiently suppress the $K^{+}\\rightarrow \\pi^+a$ decay rate. One must also have a sufficiently short lived axion so it cannot be detected by the other experiments. To accomplish that, one has to couple both the $c$ and $b$ quarks to the same higgs field. However, as we know the limits on the strangeness changing neutral currents are very tight and for that reason we have to couple the strange quark to the same higgs doublet as $c$ and $b$. Thus the only way to realise the Peccei-Quinn symmetry is to couple either the $u$ or the $t$ or both quarks to a second higgs doublet. So we have three different models in our hands, two of which have the cosmologically desirable property that $N={1\\over 2}$, and hence have no domain wall problem. Although the original models are also ruled out, along with the PQWW axion, by the astrophysical constraints, they also have extensions with a higgs singlet. The astrophysical constraints on these models, presented in the next section, have not previously been considered, and we present an analysis in this paper of the bounds on the axion scale that can be inferred from their couplings to electrons, photons, and nucleons in astrophysical processes. We find, as for the standard DFSZ \\cite{18} and KSVZ \\cite{12} invisible axions, that the tightest constraint comes, via the nucleon coupling, from SN1987a. However, in these variant models, the coupling is generally weaker, and weakens the lower bound on $v_a$ by a factor of between 1.4 and 35, depending on which quarks the axion couples to, and on the ratio of the vacuum expectation values of the higgs doublets. The cosmological upper bound on $v_a$ from the axion density is also reduced, by a factor of either 3 or 6, as a result of the smaller QCD anomaly factor. ", + "conclusions": "Although the axion solution to the strong CP problem is one of the most physically appealing, axions themselves face a great problem. Despite the fact that they interact very weakly with matter and so are very difficult to track, particle physics experiments together with astrophysical considerations and cosmology have managed to constrain axion models significantly. In this paper we have found the constraints on for axion models with non-standard couplings to quarks and leptons, using data from E143 to determine the values of the nucleon couplings, which provide the strongest constraint. We find that the bounds are generally weakened, as the nucleon couplings in our variant axion models are smaller. The most spectacular effect is for Model I near $\\beta \\sim \\pi/4$: the bound dips to about $v_a > 2\\times 10^8$ GeV, about a factor 35 less than the DFSZ value. Models I and II have the desirable feature that the QCD anomaly coefficient $N={1\\over 2}$, which means that they have no domain wall problem, and therefore are viable models for an axion string scenario. For Model I, the lower bound on $v_a$ can dip to $2\\times10^8$ GeV for values of $\\beta$ near $\\pi/4$. Recalling the upper bound on $v_a$ in the axion string scenario, $v_a < 2.6\\times10^{11}$ GeV for $H_0 = 50$ km s$^{-1}$ Mpc$^{-1}$ \\cite{14}, we see that the `window' for this axion string scenario is actually quite large. In an inflationary scenario, the cosmological upper bound is on $v_a/N$, and so the upper bound on $v_a$ itself is reduced by a factor of 6 for Models I and II, and a factor 3 for Model III." + }, + "9708/astro-ph9708127_arXiv.txt": { + "abstract": "\\noindent We report the serendipitous detection of an X-ray source, \\AXc\\, with the ASCA Gas Imaging Spectrometer. \\AX is identified with a LINER/starburst-type spiral galaxy KUG\\thinspace 1750+683A at a redshift $z = 0.05$. It has a hard X-ray spectrum, consistent with that of the X-ray background (XRB) in the 1-10~keV band. Despite the optical classification, the X-ray luminosity cannot be explained by starburst activity. Combined with spatial variations in the optical emission line ratios, this suggests the presence of an obscured Seyfert nucleus embedded within a starforming galaxy. Similar behaviour could explain the ambiguous properties of the faint narrow-line X-ray galaxies (NLXGs) emerging from deep X-ray surveys. ", + "introduction": "Although the origin of much of the soft XRB below $\\sim 2$~keV is now understood in terms of the integrated contribution of quasars and faint NLXGs (Boyle et al 1995; Roche et al 1995; Carballo et al 1995; Griffiths et al 1996; McHardy et al 1996; Hasinger 1997), the situation is less clear at higher energies where the bulk of the energy density occurs. A strong possibility is that the NLXGs have hard X-ray spectra similar to that of the XRB, which is well fitted by a power-law of photon index $\\Gamma=1.4$ in the 1-7~keV band (Gendreau et al 1995; Chen et al 1997). The NLXGs may therefore dominate the background contribution at harder energies. In the softer 0.1--2.4~keV ROSAT band there are already indications that NLXGs show significantly harder spectra than other types of X-ray sources (Almaini et al 1996, Romero-Colmenero et al 1996). The hard spectra of NLXGs may be intrinsic to the continua of those sources, making them unlike any other well-studied objects, or it may be produced by the integrated effect of varied levels of intrinsic absorption in a more typical active galaxy population. The 2--10 keV spectra of sources detected serendipitously are therefore of great interest since they may be the brighter, and probably nearer, members of the population dominating the XRB. Studies of such objects may therefore help to reveal the true nature of this X-ray population. So far most NLXGs have been discovered in deep ROSAT observations, where they are now thought to dominate the X-ray source population at faint fluxes below $S(0.5-2.0$\\,keV$)=4\\times10^{-15}$erg$\\,$s$^{-1}$cm$^{-2}$. A small number of bright NLXGs have been known for many years (see e.g. Piccinotti et al 1982; Lawrence \\& Elvis 1982) but their relevance to the XRB is unclear. Serendipitous sources detected with ASCA provide a new way forward, since good spectra are then obtained over the whole 1-10~keV band. The discovery of possible flat X-ray spectrum from NGC 3628 (Yaqoob et al 1995) has already highlighted the possible importance of previously hidden AGN in explaining the hard XRB. We report here on the serendipitous detection of a previously unidentified X-ray source with a flat X-ray spectrum, discovered during the ASCA observation of Mrk~507. The object is $\\sim 20$ arcmin away from Mrk~507 and seen in the same Gas Imaging Spectrometer (GIS) field of view. We refer to this source as \\AXc. A ROSAT Position Sensitive Proportional Counter (PSPC) image of the same field greatly improved the positional uncertainty of the source and allowed us to identify it with a galaxy which has a UV excess. An optical spectrum of \\AX then allows us to obtain the redshift and deduce the nature of the activity in this galaxy. Our results show that it is a low redshift and hence one of the brightest ($\\sim 10^{-12}$erg cm$^{-2}$s$^{-1}$) members of the NLXG population. A value of the Hubble parameter of $H_0=50$ km s$^{-1}$ Mpc$^{-1}$ and a cosmological deceleration parameter of $q_0={1\\over 2}$ have been assumed throughout. ", + "conclusions": "We have serendipitously detected a hard-spectrum X-ray source, \\AXc, which is identified with a spiral galaxy at a redshift of 0.05. This galaxy shows an integrated LINER/starburst-like optical spectrum but the X-ray luminosity is two orders of magnitude higher than any known starburst galaxy. Spatial variations in the optical emission lines suggest the presence of a more ionized component. In addition, the lack of significant far infra-red emission is inconsistent with a starburst origin for the huge X-ray flux. A plausible explanation is the presence of a moderately obscured active nucleus surrounded by a starforming galaxy. This would simultaneously explain the X-ray emission and the optical spectrum. The flat X-ray spectrum below 5~keV, combined with the optical spectrum, suggest that this object may be a local counterpart to the faint NLXGs thought to be responsible for the origin of the hard XRB. This raises the intriguing possibility that many of these objects could also contain AGN, embedded within starforming host galaxies. The use of integrated optical spectra would then lead to an ambiguous classification. Many of the NLXGs currently classified as starburst galaxies could well be more distant counterparts to \\AXc. Hidden AGN, of which \\AX is a low redshift and low column density example, may therefore provide the origin of the hard XRB, as originally suggested by Setti \\& Woltjer (1989), and later modelled by Madau et al (1994) and Comastri et al (1995)." + }, + "9708/astro-ph9708061_arXiv.txt": { + "abstract": "Parallaxes for 581 bright K giants have been determined using the Hipparcos satellite. We combine the trigonometric parallaxes with ground based photometric data to determine the K giant absolute magnitudes. For all these giants, absolute magnitude estimates can also be made using the intermediate band photometric DDO system (Janes 1975, 1979). We compare the DDO absolute magnitudes with the very accurate Hipparcos absolute magnitudes, finding various systematic offsets in the DDO system. These systematic effects can be corrected, and we provide a new calibration of the DDO system allowing absolute magnitude to be determined with an accuracy of 0.35 mag in the range $2 > M_V > -1$. The new calibration performs well when tested on K giants with DDO photometry in a selection of low reddening open-clusters with well-measured distance moduli. \\begin{keywords} G and K giants -- absolute magnitudes, parallaxes \\end{keywords} ", + "introduction": "K giants are bright and ubiquitious, occuring in a wide range of Galactic populations, and are convenient tracer objects for examining the structure and kinematics of the Galaxy. The chief difficulty with these objects has been the uncertainty in their absolute magnitudes, which arises from the fact that stars on the giant branch form from a wide range of mass and age, as well as the giant branch being rather steep as a function of colour. K giants span a broad range of absolute magnitude $M_V$ from about $2 < M_V < -3$. In some studies accurate distances to the tracer objects are required such as in Bahcall, Flynn and Gould (1991) and Flynn and Fuchs (1994), who used the kinematics of K giants to constrain the amount of dark matter present in the Galactic disc. In these studies the absolute magnitudes and hence distances of the giants were estimated using David Dunlop Observatory (DDO) photometry. This is an intermediate band photometric system of six filters, four of which can be used to estimate physical parameters for late type giants (McClure 1976). The four filters are called 41, 42, 45 and 48 and are at the wavelengths 4166, 4257, 4517 and 4886 \\AA~and have passbands of 83, 73, 76 and 186 \\AA~respectively. Three colours, C4142, C4548 and C4245 are formed from these four filters. C4245 is primarily sensitive to effective temperature, C4548 to luminosity and C4142 can be used in combination with the other two colours to estimate stellar metallicity, [Fe/H]. Janes (1979) describes the measurement of absolute magnitude and abundance using these filters. His absolute magnitude scale was based on distances determined by means of the Wilson-Bappu effect, and was later adjusted slightly when DDO photometry had been obtained of K giants in open clusters. Janes' calibration applies to metal rich disc stars ([Fe/H] $>-1$), although the system has been extended to metal weak populations (e.g. Norris, Bessell and Pickles 1985, Morrison, Flynn and Freeman 1990, and Claria et.al. 1994). In this paper we are concerned with metal rich K giants only, [Fe/H]$>-0.5$. The {\\it European Space Agency's} Hipparcos satellite has observed all bright (apparent $V<8.0$) K giants, which are included in the all sky part of the Hipparcos Input Catalogue (HIC), so that very high accuracy trigonometric parallaxes are now available in the Hipparcos Catalogue (ESA 1997). Before Hipparcos, only a handful of giants had accurate parallax measurements, whereas Hipparcos has now measured accurate parallaxes for around 600 giants. In this paper, we use the parallaxes of local K giants from Hipparcos to check the DDO system's absolute magnitude calibration. Our sample is described in section 2. A number of systematic offsets are found in the DDO absolute magnitudes, particularly for the redder ($B-V>1.2$) giants and around the clump giants. In section 3 we develop a new calibration of the DDO system, tied to the Hipparcos results. In addition to removing the systematic errors in the old system, the new calibration is a good deal simpler to use. Absolute magnitudes of K-giants can thus be determined photometrically over the range $2 < M_V < -1$ with an accuracy of 0.34 magnitudes. In section 4 the new calibration is found to be satisfactory when checked with K giants in old open clusters for which DDO photometry is available in the literature. The calibration developed here can be applied in a wide range of Galactic structure studies, one of which is the measurement of the amount of dark matter in the Galactic disc. We draw our conclusions in section 5. ", + "conclusions": "We have used very accurate Hipparcos data for a sample of 581 local K giants to measure their absolute magnitudes. These stars were used to check the absolute magnitudes derived using intermediate band DDO photometry. A number of systematic offsets in the DDO system emerged from this comparison. The Hipparcos data were then used to derive a new calibration of absolute magnitude in the DDO system. We have checked our new calibration satisfactorily against K giants with DDO photometry in 17 open clusters. Our calibration is appropriate for K giants in the colour range $0.85 < $C4245$ < 1.15$ or approximately $0.95 < B-V < 1.3$, and with [Fe/H]$ > -0.5$, (i.e. metal rich stars) with an accuracy of 0.35 mag. The quality of the Hipparcos data enable a great improvement to be made in the estimation of K giant $M_V$ by photometric methods. With the release of the full data set in mid-1997, it should be possible to extend the calibration presented here to giants with [Fe/H] $< -0.5$ and to subgiants. K giants have a venerable tradition in Galactic structure studies, such as tracing inner and outer disc kinematics, measuring the disc scale length, properties of the Bulge, and the disc dark matter problem. The latter was our primary motivation for this study, since the work of Bahcall, Flynn and Gould (1992), Flynn and Fuchs (1994) used K giants to trace the scale height and kinematics of the disc and place limits on its dark matter content. This new calibration will provide an improved measure of the distances of K giants and we are looking forward to using it to reanalyse the disc dark matter problem." + }, + "9708/astro-ph9708257_arXiv.txt": { + "abstract": " ", + "introduction": "The data were obtained with the Mid--infrared Array eXpandable (MAX) camera constructed by Infrared Labs for the Max--Planck--Institut f\\\"ur Astronomie. The MAX camera is built around a Rockwell 128$\\times$128 Si:As BIB array which provides a field of view 35'' $\\times$ 35'' when mounted on the 3.8m United Kingdom Infrared Telescope (UKIRT). Observations were made at UKIRT on August 26--27, 1996 during photometric conditions with diffraction--limited images (FWHM $\\sim 0.7$'') obtained through an N--band filter ($\\lambda_{eff} = 10.16 \\mu$m; $\\Delta \\lambda = 5.20 \\mu m$). Data were collected while chopping the telescope N--S (12'') at a rate of 2 Hz, and nodding the telescope (12'') every 50 seconds to correct for non--uniform illumination effects introduced by chopping. Data were reduced according to standard image processing techniques except that no flat--field corrections were applied. Images obtained at each end of the ``chop'' were subtracted from each other to remove bias, dark current, and thermal background. Co--added images from both ``nod'' positions were averaged and aperture photometry was performed on the final images with a diameter of 3.12'' using a sky annulus of 5.2--10.4''. Flux calibration was derived by observing standards from the list of Cohen {\\it et al.} (1992). Both DI and the nearby DH Tau ($sep = 15.1$''; $PA = 307^{\\circ}$) were observed simultaneously on the array ($T_{int} = 250.0$s), interspersed with observations of the standard star HR1370 ($T_{int} = 50.0$s) at nearly the same airmass ($\\Delta X < 0.1$). Comparison of photometry from stellar images appearing on different portions of the array indicates residual uncertainties in the calibration less than $\\pm 5$\\%. Derived fluxes and associated errors (dominated by the thermal background) are: DH Tau $F_N = 0.137 \\pm 0.005$ Jy (6.26$^m$) and DI Tau $F_N = 0.030 \\pm 0.005$ Jy (7.90$^m$). Additional observations were obtained in the Q--band ($\\lambda_c = 19.91 \\mu$m; $\\Delta \\lambda = 1.88 \\mu m$). during non--photometric conditions yielding a flux ratio of $> 1.7$ between DH Tau (detected) and DI Tau (undetected). ", + "conclusions": "" + }, + "9708/astro-ph9708075_arXiv.txt": { + "abstract": "Although it is just over a year since the data was made public, the HDF exposure has stimulated considerable progress towards our understanding of the faint galaxy population. I present a brief personal account of the history of faint galaxy studies culminating in the HDF, and describe what I consider to be the main highlights thus far from this remarkable image. The HDF has given considerable impetus to studies of galaxy evolution and this has led to the emergence of a convincing empirical framework. Further exploitation of deep HST images in conjunction with ground-based 2-D spectroscopy will assist in the physical understanding of the evolutionary processes involved. ", + "introduction": "We're here to celebrate and discuss scientific results from the Hubble Deep Field (Williams et al 1997). Most would agree that this exposure represents an observational landmark in the long arduous path of exploring and understanding the Universe of faint galaxies. Indeed, it is difficult to remember a single observation in astronomy that has influenced our subject so quickly. Moreover, its full impact may not yet be realised. In this workshop, we will debate the significance of the conclusions so far derived and learn of new developments that follow directly from the HDF. This remarkable image has acted as an inspiration to many astronomers because, exceptionally, we were granted immediate access to the data. Those working with other facilities, such as ISO (Rowan-Robinson, this volume) and the VLA (Kellermann, this volume) have been quick to follow the example by concentrating their deepest exposures on this same field. As well as explaining what I consider to be the main extragalactic highlights from the HDF at this point, largely to set the scene for the more detailed articles that follow, I will recall some of the earlier work which inserted pieces of the jigsaw that we now recognise more clearly via the HDF. Of course, HDF has also had significant impact in the non-extragalactic area. I'm glad this is well covered in the workshop but won't attempt to review progress in those important areas. ", + "conclusions": "It's an exciting time to be doing cosmology! The last time we changed government in England was when the first deep photographic counts were published and Beatrice Tinsley suggested measuring the redshifts of a sample of galaxies to $B$=21. She predicted a small fraction of the bluest sources might be high redshift primordial galaxies. After 15 years of ground-based work and only 4 years of post-refurbishment HST data, we have clearly come a long way. The acceleration of this subject in the past 2 years owes a great deal to HST and, within that context, to the HDF itself. There is much more data to come and much more physics to do. The much heralded `synthesis' of theory and data is premature in my view. So far we are mainly surveying. The more fundamental task of understanding will take considerably longer. When we finally get there, I believe we will all recall that moment when we first saw the spectacular image of the Hubble Deep Field." + }, + "9708/astro-ph9708133_arXiv.txt": { + "abstract": "We present results from spectral analysis of ASCA data on the strong Fe{\\sc ii} narrow-line Seyfert 1 galaxy Mrk\\thinspace 507. This galaxy was found to have an exceptionally flat ROSAT spectrum among the narrow-line Seyfert 1 galaxies (NLS1s) studied by Boller, Brandt \\& Fink (1996). The ASCA spectrum however shows a clear absorption feature in the energy band below 2 keV, which partly accounts for the flat spectrum observed with the ROSAT PSPC. Such absorption is rarely observed in other NLS1s. The absorption is mainly due to cold (neutral or slightly ionized) gas with a column density of (2--3)$\\times 10^{21}$\\psqcm. A reanalysis of the PSPC data shows that an extrapolation of the best-fit model for the ASCA spectrum underpredicts the X-ray emission observed with the PSPC below 0.4 keV if the absorber is neutral, which indicates that the absorber is slightly ionized, covers only part of the central source, or there is extra soft thermal emission from an extended region. There is also evidence that the X-ray absorption is complex; an additional edge feature marginally detected at 0.84 keV suggests the presence of an additional high ionization absorber which imposes a strong O{\\sc viii} edge on the spectrum. After correction for the absorption, the photon index of the intrinsic continuum, $\\Gamma\\simeq 1.8$, obtained from the ASCA data is quite similar to that of ordinary Seyfert 1 galaxies. Mrk\\thinspace 507 still has one of the flattest continuum slopes among NLS1, but is no longer exceptional. The strong optical Fe{\\sc ii} emission remains unusual in the light of the correlation between Fe{\\sc ii} strengths and steepness of soft X-ray slope. ", + "introduction": "Mrk~507 is a narrow-line Seyfert galaxy at a redshift of $z = 0.0559$ (Halpern \\& Oke 1987). Although this object had been classified as a Seyfert-2 galaxy or a LINER based on the permitted line width (FWHM$\\sim 800$\\kmps; Koski 1978) and line ratios (Heckman 1980), the detection of strong optical Fe{\\sc ii} permitted lines and a small [OIII]$\\lambda 5007$/H$\\beta$ ratio favours classification as a narrow line Seyfert 1 galaxy (NLS1; Halpern \\& Oke 1987). NLS1 were first studied by Osterbrock \\& Pogge (1985); they have optical spectra which are similar to those of normal Seyfert 1 galaxies but the Balmer lines are narrow, typically with FWHM$\\leq 2000$\\kmps. A prototype of this class of galaxy is I\\thinspace Zw\\thinspace 1 (see Goodrich 1989 for the classification criteria in optical emission-line spectra for NLS1). NLS1s are different from ordinary Seyfert galaxies in their soft X-ray properties, as shown by ROSAT observations. Rapid, large amplitude X-ray variability and the lack of evidence for internal X-ray absorption strongly suggest that we are seeing direct radiation from an unobscured central source. This is consistent with the large X-ray luminosities of NLS1s as previously pointed out by Halpern \\& Oke (1987). Another remarkable X-ray property is their steep soft X-ray spectra. As shown in the ROSAT 0.1--2.4 keV photon-index--FWHM(H$\\beta$) diagram of Boller, Brandt \\& Fink (1996, hereafter BBF), photon indices ranging from 3.5 -- 5 are only found in NLS1. Although the ROSAT photon indices for NLS1s are spread over a wide range ($\\Gamma = $ 2--5), the mean value ($\\Gamma_{\\rm NLS1}\\simeq 3.1$) is significantly larger than that ($\\Gamma_{\\rm Sy1}\\simeq 2.3$) for ordinary Seyfert 1 galaxies with broad Balmer lines (BBF). This may be related to an intrinsic difference in the emission mechanism. Mrk~507, which is an optically-classified NLS1, stands out since it shows an exceptionally flat ROSAT spectrum for any type of Seyfert 1 galaxy ($\\Gamma = 1.6\\pm 0.4$, BBF). Although possible X-ray flux variations were detected from the ROSAT observations, the flat ROSAT spectrum does not fit the general properties of NLS1s (BBF). The best-fit value of the photon index is flatter by $\\sim 0.7$ than even the mean ROSAT value for ordinary Seyfert 1 galaxies. This could be taken to suggest that the spectral slopes of NLS1s vary widely between individual objects, in contrast to those of ordinary Seyfert 1 galaxies, which are found in a relatively tight range. However, it is also possible that the flat spectrum is due to a complicated absorption effect which cannot be resolved by the ROSAT PSPC because of its poor spectral resolution and limited energy range. Indeed, a strong anti-correlation between ROSAT spectral slope and FWHM(H$\\beta$) is found in a sample of quasars studied by Laor et al (1997), which are presumably less absorbed than Seyfert galaxies, i.e., it appears that no flat X-ray spectrum quasars with narrow Balmer lines are found. Since a large fraction of Seyfert 1 galaxies has been known to show evidence for absorption by partially-ionized gas, the so-called ``warm absorber'' (Halpern 1984) from a systematic study of ASCA X-ray spectra (e.g., Reynolds 1997), similar absorption effects can be expected in X-ray spectra of NLS1s. Unlike the absorption by cold material found in many Seyfert 2 galaxies due to the obscuring torus, X-rays are absorbed mainly by partially-ionized oxygen in the energy range between 0.5--2 keV in the case of a warm absorber. The energy where the deepest feature appears depends on the ionization state of the absorber. Our ASCA observation was aimed at investigating whether the remarkably flat ROSAT spectrum of Mrk\\thinspace 507 is the result of such absorption effects or if indeed it has an intrinsically flat X-ray continuum. Previous hard X-ray obsevations of Mrk\\thinspace 507 with non-imaging collimated instruments, such as the Ginga LAC, were confused by the brighter Seyfert 1 galaxy Kaz\\thinspace 163 which is $\\sim 10$ arcmin SW of Mrk\\thinspace 507, so no spectral study of Mrk~507 in the energy band above 2 keV has ever been made. A value of the Hubble parameter of $H_0=50$ km s$^{-1}$ Mpc$^{-1}$ and a cosmological deceleration parameter of $q_0={1\\over 2}$ have been assumed throughout. ", + "conclusions": "We detected significant absorption in the ASCA spectrum of a NLS1, Mrk\\thinspace 507, which plausibly causes the apparently flat spectrum in the ROSAT band measured by BBF. The absorber is neutral or possibly slightly ionized, and the column density is 2--3$\\times 10^{21}$\\psqcm, depending on models. In addition to the cold absorber, marginal evidence for an O{\\sc viii} edge due to another warm absorber with high ionization parameter was found. The spectral slope of the intrinsic continuum is similar to ordinary Seyfert 1 galaxies rather than NLS1s that generally have steeper soft X-ray continua. As one of extreme Fe{\\sc ii} AGNs (Lipari 1994), this is also unusual in the light of the correlation between EW(Fe{\\sc ii}) and steepness of soft X-ray spectra claimed by Wilkes, Elvis \\& McHardy (1987)." + }, + "9708/astro-ph9708239_arXiv.txt": { + "abstract": "Rossi X-ray Timing Explorer (RXTE) All Sky Monitor observations of the transient Be star X-ray source \\src\\ suggest a 34.5 day period. This is apparently confirmed by a serendipitous RXTE Proportional Counter Array (PCA) slew detection of the source on 1997 May 5, near the time of a predicted flux maximum. A subsequent 5ks pointed observation of \\src\\ with the RXTE PCA detector was carried out on 1997 June 7, when \\src\\ was predicted to be bright again, and this revealed pulsations at a period of 103.2 seconds. If the 34.5 day period is orbital, then the pulse period is surprisingly long compared to that predicted by the correlation between orbital period and spin period observed for other Be/neutron star systems. A possible similarity with \\2058\\ is briefly discussed. ", + "introduction": "The transient X-ray source \\src\\ (=4U0728-25, 3A0726-260) has a Be star optical counterpart that was optically identified by Steiner et al. (1984). This star was classified as B0Ve by Corbet \\& Mason (1984) with a corresponding distance of 4.6$\\pm$1.3 kpc. Negueruela et al. (1996), however, derive a slightly earlier spectral type of O8-9Ve with a distance of 6.1$\\pm$0.3 kpc. The X-ray source itself has not been extensively observed, although \\src\\ was also detected with Uhuru (Forman et al. 1978), HEAO-1 (Wood et al. 1984) and ROSAT (Haberl, cited in Negueruela et al 1996), the only published X-ray light curve comes from Steiner et al. (1984) who present a rather sparse long term light curve from the Ariel V Sky Survey Instrument (SSI). Pointed EXOSAT observations on 1984 Oct. 29 and 1984 Nov. 29 failed to detect any X-ray emission, even though simultaneous optical observations from the Isaac Newton Telescope showed the presence of \\halpha\\ emission indicating that a circumstellar envelope was still present. In this paper we present the results of observations of \\src\\ made with two instruments on board the Rossi X-ray Timing Explorer (RXTE): the All Sky Monitor (ASM) and the Proportional Counter Array (PCA). These observations provide information on the orbital period and spin period of the neutron star in the system. ", + "conclusions": "RXTE observations of \\src\\ indicate a pulse period of 103.2 seconds and a most likely orbital period of \\sqig35 days. Additional more extensive pulse timing observations would be valuable to determine the orbital parameters of this system." + }, + "9708/hep-ph9708226_arXiv.txt": { + "abstract": "The dilaton field in string theories (if exists) is expected to have a mass of the order of the gravitino mass $m_{3/2}$ which is in a range of $10^{-2}$keV--1GeV in gauge-mediated supersymmetry breaking models. If it is the case, the cosmic energy density of coherent dilaton oscillation easily exceeds the critical density of the present universe. We show that even if this problem is solved by a late-time entropy production (thermal inflation) a stringent constraint on the energy density of the dilaton oscillation is derived from experimental upperbounds on the cosmic X($\\gamma$)-ray backgrounds. This excludes an interesting mass region, $500{\\rm keV} \\lesssim m_{3/2} \\lesssim 1{\\rm GeV}$, in gauge-mediated supersymmetry breaking models. ", + "introduction": " ", + "conclusions": "" + }, + "9708/astro-ph9708054_arXiv.txt": { + "abstract": "The geometry of Freedman-Roberston-Walker cosmological models is fixed by the mass density parameter, $\\Omega_M$, and the cosmological constant, $\\Omega_\\Lambda$. The classical volume-redshift cosmological relation is a sensitive ${\\bf\\Omega}=[\\Omega_M,\\Omega_\\Lambda]$ indicator but its redshift dependence is observationally degenerate with the luminosity or number density evolution of galaxies. Introducing a measurement of the invariant co-moving mass density of the universe reduces the problem of galaxy evolution to a differential measurement between clustered and field galaxies. The cost is a 25\\% reduction in sensitivity to the $\\bf\\Omega$'s, although this test still remains 50\\% more $\\bf \\Omega$ sensitive than the magnitude-redshift relation. An implementation of the test as the product of the mass-to-light ratio, $M/L$, of some clustered systems such as galaxy groups or clusters, with $j/\\rho_c$, the normalized luminosity density, is considered. Over the zero to one redshift range the apparent $\\Omega_e(z)=M/L\\times j/\\rho_c$ has a zero point and slope related to $\\Omega_M$ and $\\Omega_\\Lambda$, respectively. All quantities are used in a differential sense, so that common selection effects, dynamical scale errors, and galaxy evolution effects will largely cancel. The residual differential galaxy evolution between field and the clustered galaxies can be measured from the sample data. Monte Carlo simulations, calibrated with observational data, show that 20 clusters spread over the zero to one redshift range, each having 100 cluster velocities, allows a 99\\% confidence discrimination between open and closed low density universe models. A similarly distributed sample of about 100 rich clusters, or about 1000 galaxy groups found within a large field survey, will measure $\\Omega_\\Lambda$ to about 7\\% statistical error. ", + "introduction": "The existence of a nonzero cosmological constant, $\\Lambda$, would have profound significance for our understanding of the universe and its physics (\\cite{weinberg}). The expansion history and geometry of the universe, as described by the Freedman-Roberston-Walker (FRW) solution, are completely determined given the density parameter, $\\Omega_M=\\rho_0 8\\pi G/3H_0^2$, and the cosmological parameter, $\\Omega_\\Lambda=\\Lambda c^2/3H_0^2$. Knowledge of ${\\bf \\Omega}=[\\Omega_M,\\Omega_\\Lambda]$ is also of practical concern to interpret the physical properties of objects at large redshifts. The geometrical effects of the cosmological parameters are the basis of a number of classical tests of the world model. These include the redshift dependence of galaxy numbers, sizes or luminosities (\\eg\\ \\cite{s61a,ppc}). The success of any of these tests is in large part dependent on the degree to which the evolution of the intrinsic properties of galaxies is understood so that those effects can removed to leave the cosmological variation of interest (\\eg\\ \\cite{s61b,tinsley,s88,oh,tr}). The number count magnitude relation has long been taken as a hint that the $\\Omega_M=1$ model might not be correct, but this remains locked in the controversies of galaxy formation (\\eg\\ \\cite{ls,koo_kron,ldss,cowie,fukugita,sf}). Although both the observational and theoretical understanding of galaxy evolution is advancing rapidly the fundamental degeneracy between galaxy evolution measurements and the cosmological parameters means that to derive a reliable empirical model requires additional information. There are a number of alternate approaches designed to establish observational constraints on the value of $\\Omega_\\Lambda$. One geometrical test is to compare the redshift and angular extensions of some physically understood shape at a large redshift, such as the quasar-quasar correlation function (\\cite{ap,qsos1,qsos2,qsos3}). The optical depth for multiple gravitational images of distant quasars increases rapidly with positive values of $\\Omega_\\Lambda$. The relatively low frequency of split images argues that $\\Omega_\\Lambda\\lesssim 0.7$ (\\cite{cpt,mr,kochanek}). The high precision of photometry coupled with the growing understanding of supernovae, particularly those of type Ia (\\cite{hamuy,riess}), allows the classical magnitude-redshift test to be implemented. With sufficient data over a large redshift range (\\cite{perlmutter,schmidt}) both the $\\bf \\Omega$ values can be measured. Of particular note is the ``first Doppler peak'' in the angular fluctuation spectrum of the Cosmic Background Radiation which is a measurement over the largest possible path length of the geometry of the universe (\\eg\\ \\cite{bond} and references therein). The cosmological parameters are of sufficient importance that they will be measured with a variety of independent methods to establish their values with confidence, understand the astrophysics of the objects, and, to some degree test the FRW model itself. The purpose of this paper is to note a variant of the classical volume-redshift test which breaks the cosmology-galaxy evolution degeneracy of the volume-redshift relation. That is, the co-moving mass density, which is an invariant at low redshift in all conventional cosmologies, is equal to the product of the total mass-to-light ratio, $M/L$, with the field luminosity density, $j$ (\\cite{oort}). The virialized systems can range from rich clusters to small groups of galaxies. Virialized systems have the considerable benefit that their mass profiles can be inferred from dynamical techniques, independent of the distribution of the galaxies. Furthermore dynamical mass measurements are distance independent (other than the cosmological factors of interest) with no corrections required to compare masses at different redshifts. The following section briefly reviews the volume-redshift relation in the FRW models. The relations are expanded to first order to illustrate the parameter dependencies and their degeneracy in the luminosity function. The mass-density constraint is introduced and the redshift dependence of the apparent $\\Omega_e(z)$, which is a function of the true $\\bf \\Omega$, is shown in section 3. In Section 4 the random errors and data requirements of practical measurements are evaluated, concluding that a high precision measurement is primarily a matter of assembling the appropriate datasets, which is likely to be done anyway for a variety of other purposes. ", + "conclusions": "The classical volume-redshift test, which depends upon an absolute comparison of galaxy numbers or luminosities at different redshifts, can be modified to create a much more reliable, completely differential, test. The extra ingredient is to combine quantities which together give the mean co-moving mass density of the universe, which is a conserved quantity. This benefit comes at the modest cost of a 25\\% reduction in $\\Omega_\\Lambda$ sensitivity. Monte Carlo simulations show that a sample of 2000 cluster galaxies and a comparable field sample will be able to tightly constrain the differential evolution between cluster and field and will measure $\\Omega_\\Lambda$ to a precision of about 25\\%. Moreover, groups of galaxies found within a large field survey will serve the same purpose and provide a second avenue to address differential evolution between cluster and field. Differential evolution between clustered and field galaxies will be addressed using multi-color photometry and imaging. In the longer term, a survey of 100 or so rich clusters will increase the precision of the geometry measurement to about 7\\%. Such data can also give extremely precise measurements of the evolution of the sample galaxies, although these should not be taken as absolute evolutionary measurements unless care is taken to avoid redshift dependent selection effects." + }, + "9708/astro-ph9708262_arXiv.txt": { + "abstract": "A rarity among supernova, SN~1993J in M81 can be studied with high spatial resolution. Its radio power and distance permit VLBI observations to monitor the expansion of its angular structure. This radio structure was previously revealed to be shell-like and to be undergoing a self-similar expansion at a constant rate. From VLBI observations at the wavelengths of 3.6 and 6 cm in the period 6 to 42 months after explosion, we have discovered that the expansion is decelerating. Our measurement of this deceleration yields estimates of the density profiles of the supernova ejecta and circumstellar material in standard supernova explosion models. ", + "introduction": "Supernova SN~1993J in M81 discovered by Francisco Garc\\'{\\i}a of Lugo , Spain (\\cite{RG93}) is a type IIb supernova (SN) whose red giant progenitor probably had a mass of 12-16 \\sun while on the main sequence; at the time of the explosion, 3-5 \\sun likely remained in the He core and $\\stackrel{<}{\\sim}$1 \\sun in the He/H envelope (\\cite{F93}, \\cite{S94}, \\cite{W94}, \\cite{I97}). The first maximum in the supernova optical light curve has been attributed to shock heating of the thin envelope and the second to radioactive decay of $^{56}$Co (\\cite{N93}, \\cite{S94}, \\cite{W94}). Modelling of the X-ray emission (\\cite{SN95}) also implies a relatively low mass envelope due to interaction with a binary companion (\\cite{N93}, \\cite{W94}). \\\\ The standard circumstellar interaction model --hereafter standard model or SM-- for radio supernovae (\\cite{C96} and references therein) suggests that the radio emission arises from a shocked region between the supernova ejecta and the circumstellar material (CSM) that results from the wind of the SN's progenitor star. More specifically, the SM considers SN ejecta with steep density profiles ($\\rho_{ej}\\propto r^{-n}$) shocked by a reverse shock that moves inwards from the contact surface and a CSM with density profile $\\rho_{csm}\\propto r^{-s}$ shocked by a forward shock that moves outwards from the contact surface ($s$=2 corresponds to a steady wind). For $n$$>$5, self-similar solutions are possible (\\cite{C82a}); the radii of the discontinuity surface, forward shock and reverse shock are then related and all evolve in time with a power law R $\\propto t^{m}$ ($t$, time after explosion), where $m$=$(n-3)/(n-s)$. \\\\ SN~1993J is the closest SN that is both young and radio bright (\\cite{W96}) and hence offers a unique opportunity for the study of its radio structure and the test of radio supernova models (\\cite{C82a}, \\cite{SN95}). \\mara found the radio structure to be shell-like. Multiwavelength radio light curves and high resolution radio images of SN~1993J (\\cite{V94}, \\cite{M95b}, respectively) established the self-similar nature of the expansion. \\\\ The technique of VLBI can, in principle, determine $m$ directly by simply observing the angular growth rate of the supernova. \\bartel and \\marb found that $m$=1 was compatible with their results to within their respective uncertainties. In this paper, we present VLBI results for $\\lambda$6 cm through October 1996 (42 months after explosion), combined with those already published for $\\lambda$3.6 cm (\\cite{M95b}), to estimate the deceleration in the supernova expansion and to infer the density profiles of the supernova ejecta and CSM. \\\\ ", + "conclusions": "Our maps give no indication of any structures developing in the shell by the action of either Raleigh-Taylor instabilities (\\cite{CB95}) or interaction with the CSM. There is also no evidence of any departure from circularity as suggested by some authors to explain the action of a putative binary companion. We also do not see any emission above $\\sim$0.5 mJy from any compact source at the center of the structure (i.e., a pulsar as suggested by \\woos and \\shige). \\\\ Within the framework of self-similar models, measurement of the time dependence of the attenuation of the supernova radio emission due to the circumstellar plasma allows us to estimate the exponent of a power law representation of the density profile of the CSM: for free-free absorption as commonly invoked in radio supernova models (Weiler et al. 1996 and references therein), the opacity, $\\tau$, is proportional to the density squared integrated along the line of sight. Given a supernova radius R $\\propto$ $t^{m}$ and $\\rho_{csm} \\propto r^{-s}$, then $\\tau$ $\\propto$ $t^{2m(-s+0.5)}$. \\vandyk found $\\tau$ $\\propto$ $t^{\\delta}$ with $\\delta$=$-1.99^{+0.38}_{-0.16}$ for the homogeneous component of the CSM. Combining this result with $m$=0.86$\\pm$0.02, we obtain $s$=$1.66^{+0.12}_{-0.25}$\\,. This value is lower than the $s$=2 in the SM for a constant stellar wind, but very close to the value $s$=1.7 given by \\fransnusei to explain the X-ray emission (\\cite{Z94}). \\vandyk also obtain a similar time dependence for the attenuation of a clumpy medium and hence argue that the clumpy component is spatially distributed in the same way as the homogeneous component. \\houfran also argue in favor of a clumpy medium based on optical line profiles. \\suznom postulate CSM with homogeneous and clumpy components to explain X-ray data, but they consider two regions: (1) An inner homogeneous region with a density profile described by $s$=1.7 out to radii smaller than $\\sim$ 5$\\times$$10^{15}$cm and (2) an outer clumpy region with density profile described by $s$=3 for the interclump medium at larger radii. Such a model of the CSM allows \\suznom to fit their model to all of the available X-ray data. Specifically, the $s$=3 clumpy medium is needed to account for the hard X-rays and for part of the H$\\alpha$ emission. The supernova explosion model of \\suznom, consisting of ejecta and a clumpy CSM as described above, is very different from self-similar models (\\cite{C82b}). \\\\ The self-similar case with $m$=0.86 and $s$=1.66 gives an ejecta density profile of $n$=$11.2^{+3.5}_{-1.8}$\\,. These values correspond to steep profiles, indeed much steeper than the profiles of white dwarfs ($n$=7), but less steep than those suggested by \\baron from spectral analyses or those used by \\suznom. In the SM for values $n$=11.2 and $s$=1.66, the reverse shock radius is $\\sim$2\\% smaller, and the forward shock radius $\\sim$20\\% larger, than the radius of the contact surface between shocked supernova ejecta and shocked CSM (\\cite{C82b}). \\\\ \\marb estimate that the width of the radio shell is about 0.3 times the size of the outer radius (or, equivalently, about 40\\% of the inner radius). These authors also estimate expansion speeds $\\sim$ 15,000 km $s^{-1}$ which are compatible with the largest velocities ($\\sim$ 11,000 km $s^{-1}$) measured in H$\\alpha$ (\\cite{F94}, \\cite{P95}) if (i) the H$\\alpha$ emission originates in the vicinity of the reverse shock, (ii) a homologous expansion is assumed in the ejecta and shocked regions, and (iii) the shock shell is about twice as large as predicted in the SM . In an attempt to reconcile the SM and the observational results, \\houfran suggest that clumpy ejecta and/or CSM can broaden the shell. On the other hand, the region of the ejecta shocked by the reverse shock in the model of \\suznom is even larger than that of the CSM shocked by the forward shock. However, the maximum speeds of the radio outer shell and of the region of H$\\alpha$ emission in the model of \\suznom match those observed very well, although the density and velocity profiles in the shell are very different from those of the standard model. \\\\ If we consider only VLBI results from epochs more than 500 days after the explosion, we obtain $m$=0.89$\\pm$0.03 and $s$=$1.68^{+0.10}_{-0.27}$\\,. However, such an age range is in the region in which the \\suznom model suggests $s$=3.0. A contradiction is apparent and our results therefore argue against their model. Our estimate of $s$ based on that of $m$ is not dependent on a given explosion model but is a determination from the time dependence of the opacity due to an external medium (\\cite{W86},\\cite{V94}). Furthermore, such time dependence of the opacity has not changed between days 200 and 1000 (Van Dyk, priv. comm.). \\\\ If the physical picture of the radio and H$\\alpha$ emission in the SM were correct, the $\\sim$15\\% decrease in expansion speed measured by VLBI between months 12 and 42 after explosion should be observable in the H$\\alpha$ emission. On the other hand, if the model of \\suznom were correct, a decrease in the maximum speed of H$\\alpha$ would not be expected. \\\\" + }, + "9708/astro-ph9708112_arXiv.txt": { + "abstract": "We report the observation of features near 1 keV in the {\\it ASCA} spectra from three ``Narrow Line Seyfert 1'' (NLS1) galaxies. We interpret these as oxygen absorption in a highly relativistic outflow. If interpreted as absorption edges, the implied velocities are 0.2--0.3~c, near the limit predicted by ``line-locking'' radiative acceleration. If instead interpreted as broad absorption lines, the implied velocities are $\\sim0.57$~c, interestingly near the velocity of particles in the last stable orbit around a Kerr black hole, although a physical interpretation of this is not obvious. The features are reminiscent of the UV absorption lines seen in broad absorption line quasars (BALQSOs), but with larger velocities, and we note the remarkable similarities in the optical emission line and broad band properties of NLS1s and low-ionization BALQSOs. ", + "introduction": "Narrow-line Seyfert 1 galaxies (NLS1s) are defined by their optical line properties (e.g. Goodrich\\markcite{11}~1989): ({\\it i.}) the Balmer lines are only slightly broader than the forbidden lines (H$\\beta$~FWHM$<2000\\,\\rm km/s$); ({\\it ii.}) the forbidden line emission is relatively weak ([O~III]/H$\\beta<3$); ({\\it iii.}) there are often strong emission features from Fe~II and high ionization optical lines. It has recently been discovered that they also have distinctive X-ray properties. {\\it ROSAT} PSPC observations found that the soft X-ray spectra are systematically steeper than ``classical'' Seyfert 1s and that the photon index appears correlated with the optical line width. NLS1s also very frequently exhibit rapid and/or high amplitude X-ray variability (Boller, Brandt \\& Fink\\markcite{2}~1996; Forster \\& Halpern\\markcite{10}~1996 and references therein). NLS1 observations with {\\it ASCA}, which has better energy resolution and a larger band pass, find that the steep spectrum in the soft X-ray band is primarily due to a strong soft excess component with characteristic blackbody temperature in the range 0.1--0.2~keV and a relatively weak hard power law. The hard X-ray power law slope is either remarkably variable (Leighly et al.\\markcite{18}~1996; Guainazzi et~al.\\markcite{12}~1996), or significantly steeper than found in broad-line Seyfert 1s (Pounds, Done \\& Osborne\\markcite{27}~1995; Brandt, Mathur \\& Elvis\\markcite{7}~1997). The combination of strong soft excess and steep power law prompted Pounds, Done \\& Osborne\\markcite{27}~(1995) to postulate that NLS1s represent the supermassive black hole analog of Galactic black hole candidates in the high state. ", + "conclusions": "The outflow velocities of $0.2-0.6\\rm\\,c$ inferred in these objects are very large, larger than those observed the UV spectra of BALQSOs. Such large velocities could be difficult to identify in the UV, because the prominent C~IV line would be shifted out of the band pass or into the Ly$\\alpha$ forest where it might be hard to distinguish. Also, very large velocity dispersions with moderate column densities might result in lines so broad that they blend in with the continuum. It is beyond the scope of this paper to speculate on the mechanism required to accelerate material to these very large velocities, but it is interesting to note that the velocity implied by the edge fits is close to that seen in the Galactic jet object SS~433 of $0.26\\rm\\,c$ and also to the terminal velocity predicted by ``line-locking'' of $0.28\\rm\\, c$ (Shapiro, Milgrom \\& Rees\\markcite{29}~1986). The larger velocities implied by the absorption line fits is intriguingly close to the energy of a particle in circular orbits around a Kerr black hole (e.g. Shapiro \\& Teukolsky\\markcite{30}~1993). However, it seems difficult to relate this fact to a physical outflow mechanism. We obtain lower limits on the equivalent hydrogen column. For the absorption edge model, assuming cross sections of 2.8 and $0.98\\times 10^{-19}\\,\\rm cm^{-2}$ for O~VII and O~VIII, respectively, and an oxygen abundance relative to hydrogen of $8.51\\times 10^{-4}$, the O~VII+O~VIII equivalent hydrogen column densities are in the range $0.4-1.3 \\times 10^{22}\\rm\\,cm^{-2}$. For the absorption lines, assumed to be on the linear part of the curve of growth, oscillator strengths of 0.694 and 0.416 for O~VII and O~VIII were used, giving O~VII+O~VIII equivalent hydrogen column densities of $1.6-2.1 \\times 10^{21}\\rm\\,cm^{-2}$. In each case, the column density upper limits on the absorption edges predicted to accompany these absorption lines were larger, showing that this model is viable. While estimation of the ionization parameter depends on the input continuum and is beyond the scope of this paper, we note that the O~VIII column is always larger than the O~VII column, implying a fairly high ionization parameter. Rest frame absorption features in the X-ray spectra of broad-line hard X-ray selected AGN are common (Reynolds\\markcite{28}~1997), plausibly arising in the same material as $z_{em} \\approx z_{ab}$ absorption lines found in the UV (e.g. Mathur, Elvis \\& Wilkes\\markcite{23}~1995). These `associated' absorption features may be related to the broad UV absorption lines seen in higher luminosity objects (e.g. Kolman et al.\\markcite{16}~1993). Evidence suggests that some aspect of the NLS1 central engine is significantly different compared with broad-line objects; for example, they may be characterized by a higher accretion rate relative to Eddington (Pounds, Done \\& Osborne\\markcite{27}~1995). The rapid, higher amplitude and perhaps characteristically nonlinear X-ray variability may be evidence for strong relativistic effects (Boller et al.\\markcite{1}~1997; Leighly et al.\\markcite{19}~1997). These results may indicate a higher level of activity relative to the black hole mass in NLS1s, and strong relativistic outflows might be expected. The blue-shifted absorption features discussed here are reminiscent of those found in the UV spectra of broad absorption line quasars. A connection between NLS1s and BALQSOs may be quite reasonable considering that they have many optical emission line and broad band continuum properties in common. Many NLS1s and BALQSOs show strong or extreme Fe~II emission and weak [O~III] emission. Objects which have the strongest Fe~II emission and weakest or no [O~III] emission tend to be either low ionization BALQSOs or NLS1s (Boroson \\& Meyers~\\markcite{4}~1992; Turnshek et al.\\markcite{34}~1997; Lawrence et al.\\markcite{17}~1997). Many NLS1s and low ionization BALQSOs have red optical continuum spectra, and relatively strong infrared emission (Boroson \\& Meyers~\\markcite{4}~1992; Moran, Halpern \\& Helfand\\markcite{24}~1996; Turnshek\\markcite{33}~1997). Finally, both classes are predominantly radio quiet (Stocke et al.\\markcite{31}~1992; Ulvestad, Antonucci, \\& Goodrich\\markcite{36}~1995). An interesting possibility is that the low-ionization BALQSOs and NLS1s have a common parent population (e.g. Lawrence et al.\\markcite{20}~1997). If so, perhaps objects with intermediate properties between NLS1s and low-ionization BALQSOs should exist. It was recently reported that NLS1 IRAS~13349+2438 has UV broad absorption lines (Turnshek\\markcite{33}~1997). But while most BALQSOs are X-ray quiet, this object is a bright soft X-ray source and has the very steep hard X-ray spectrum and rapid X-ray variability characteristic of NLS1s (Brinkmann et al.\\markcite{8}~1996)." + }, + "9708/hep-ph9708427_arXiv.txt": { + "abstract": "We present the results of a field theory simulation of networks of strings in the Abelian Higgs model. Starting from a random initial configuration we show that the resulting vortex tangle approaches a self-similar regime in which the length density of lines of zeros of $\\phi$ reduces as $t^{-2}$. We demonstrate that the network loses energy directly into scalar and gauge radiation. These results support a recent claim that particle production, and not gravitational radiation, is the dominant energy loss mechanism for cosmic strings. This means that cosmic strings in Grand Unified Theories are severely constrained by high energy cosmic ray fluxes: either they are ruled out, or an implausibly small fraction of their energy ends up in quarks and leptons. ", + "introduction": " ", + "conclusions": "" + }, + "9708/astro-ph9708106_arXiv.txt": { + "abstract": "We estimate the total predicted Lyc emission rates of OB associations for which the complete census of O star spectral types exists. The results are compared to the observed \\Ha\\ luminosities of the host \\hii\\ regions. We find evidence for substantial leakage of ionizing photons from some \\hii\\ regions, while others appear to be radiation bounded. We estimate that overall for the LMC, 0--51\\% of the ionizing radiation escapes the local nebulae, and would be available to ionize the diffuse, warm, ionized medium (WIM) in that galaxy. This range of values is consistent with the observed 35\\% fraction of \\Ha\\ luminosity emitted by the WIM in the LMC, as well as the corresponding fractions observed in other nearby galaxies. It is therefore possible that photoionization by O stars is indeed the dominant ionization mechanism for the WIM. ", + "introduction": "The balance of evidence currently suggests that the ionization of the diffuse, warm, ionized component of the interstellar medium (ISM) is caused primarily by O stars. From an energetic standpoint, this stellar population has long been targeted as one of the only sources capable of generating the large power requirement (\\eg Reynolds 1984) of this warm, ionized medium (WIM). Models of the WIM that assume ionization by O stars are broadly consistent with its observed properties in the Galaxy (Miller \\& Cox 1993; Dove \\& Shull 1994), and confirm that O stars are more than capable of providing the necessary ionizing power. In fact, these studies suggest that an excess of ionizing luminosity is produced, implying the escape of that radiation from the Galaxy. Observations of the WIM in nearby, external galaxies show concentrations of the diffuse gas near conventional \\hii\\ regions, further suggesting the association of the WIM with O stars (Walterbos \\& Braun 1994; Ferguson {\\etal}1996). However, there are notable complications to the O star ionization hypothesis. Perhaps most problematic is the observed limit to the ratio of the recombination lines \\ion{He}{I} $\\lambda5876$ / \\Ha, implying relative ionizing photon emission rates for He vs. H of $Q({\\rm He^0})/Q({\\rm H^0}) \\lesssim 0.03$ (Reynolds \\& Tufte 1995; Heiles {\\etal}1996). The implied ionizing spectrum in the Galaxy therefore appears to be much softer than anticipated from the O star population, implying an unusually low effective upper-mass cutoff to the stellar initial mass function (IMF) of $\\lesssim 30 \\msol$, and forcing an unrealistically large Galactic star formation rate (Heiles {\\etal}1996). Rand (1997) presents a deep spectrum of the WIM in the edge-on galaxy NGC 891, finding $Q({\\rm He^0})/Q({\\rm H^0}) \\sim 0.08$, which is much more consistent with hot star ionization. Measurements of the He and H recombination lines in the WIM of three dwarf irregular galaxies by Martin \\& Kennicutt (1997) are also consistent with He being mostly ionized in those objects, as expected from phoionization of stars with a normal IMF. However, NGC 891 exhibits an anomalously high ratio of [\\ion{N}{II}]$\\lambda6583$/\\Ha\\ $\\sim 1 - 1.4$, again implying an even harder ionizing spectrum than indicated by Rand's \\ion{He}{I}/\\Ha\\ ratio. There is also some doubt about the ability of the stellar ionizing radiation to travel the required hundreds of parsecs, although models by Dove \\& Shull (1994) and Miller \\& Cox (1993) show tentative compatibility. Investigations of the superbubble structure of the ISM (Heiles 1990; Rosen \\& Bregman 1995; Oey \\& Clarke 1997) also suggest the widespread existence of large voids, which would facilitate radiative transfer over such large distances. Finally, alternative or additional ionizing sources are likely to play a role in the WIM as well. These include turbulent mixing layers (Slavin, Shull, \\& Begelman 1993), neutrino decay from dark matter (Sciama 1990, 1995), white dwarfs, cosmic rays (Liu \\& Dalgarno 1997), and extragalactic ionizing radiation (e.g., Reynolds {\\etal}1995). In general, however, O stars remain as the most likely dominant source of ionization for the WIM. But is the simple, empirical comparison of available ionizing photons vs. ionized gas actually consistent with this scenario? Abbott (1982) predicts that roughly 15\\% of the available ionizing flux of O stars in the solar neighborhood is required to ionize the local WIM, a result supported by Dove \\& Shull (1994) and Miller \\& Cox (1993). This therefore implies that, if the large-scale WIM is similar to that in the solar neighborhood, at least $\\sim15$\\% of the OB Lyman continuum (Lyc) flux must escape the \\hii\\ regions. Photoionization models of the WIM by Domg\\\"orgen \\& Mathis (1994) further suggest that 4\\% of the total flux must escape the Galaxy altogether, in which case a certain fraction of the ionized regions must be density-bounded. However, historically, most \\hii\\ regions have been considered to be essentially radiation-bounded, rather than density-bounded, allowing the nebular luminosities to be used quite successfully as tracers of massive star formation. It is therefore worth examining more closely the comparison of available Lyc and nebular fluxes. We can investigate this effect with a sample of OB associations whose stellar populations have been classified, and whose host nebulae have measured total luminosities. The Large Magellanic Cloud (LMC) provides an ideal sample, with detailed studies of over a dozen OB associations (\\eg Massey {\\etal}1995b; Oey 1996a), and a uniform catalog of nebular photometry (Kennicutt \\& Hodge 1986) at low extinction. We now use this sample of OB/\\hii\\ systems to examine the fraction of Lyc radiation escaping from the \\hii\\ regions, as indicated by the currently available stellar atmosphere models. ", + "conclusions": "While the range of values suggests that a significant escape fraction seems likely, it is also consistent with the nebulae being generally radiation bounded, although Table~\\ref{compare} shows that some individual objects are convincingly density-bounded. But we note that the sample of faint Galactic \\hii\\ regions studied by Hunter \\& Massey (1990) also showed an overall excess of available ionizing radiation. That study predicted the Lyc emission rates from observed, classified stellar populations using the P73 conversions, and compared the results with the Lyc emission rates implied by \\Ha\\ and radio continuum observations. Inspection of their results shows that the median ratio of observed to predicted emission rates is 0.7 for both the \\Ha\\ and radio-derived comparisons, with 19 objects in the \\Ha\\ sample and 21 in the radio sample. The results from this Galactic study are thus in good agreement with the median value from our sample of 0.74. Note that the use of the P73 Lyc predictions by Hunter \\& Massey (1990) imply that the median ratios from their sample are therefore slight overestimates when comparing with our results (cf. Table~\\ref{compare}). An escaping Lyc fraction of up to 51\\% is fully consistent with the estimated fraction of diffuse to total \\Ha\\ flux of $35\\pm 5$\\% estimated by Kennicutt {\\etal}(1995) for the LMC. It is therefore indeed quite possible that the WIM, at least in the LMC, is ionized by hot supergiants. If the \\hii\\ regions in the LMC are typical in their structure, extinction, and relationship to the diffuse ISM, then these results may be compared to observations of the WIM in other galaxies. Measurements for \\Ha\\ luminosity fractions of the WIM in nearby galaxies range from $\\sim20$\\% to 53\\% (see Ferguson {\\etal}1996), which are also consistent with the fraction of ionizing radiation escaping from the LMC \\hii\\ regions. Thus it is a likely possibility that the WIM in these galaxies is dominated by O star photoionization. With better constraints on the large uncertainties in our estimate of the median $L_{\\rm obs}/L_{\\rm SdK}$, it should become possible to better evaluate this hypothesis, as well as the role of alternate ionizing mechanisms, which can contribute to the \\Ha\\ luminosities in both normal \\hii\\ regions and the WIM as well. Our selection of \\hii\\ regions encompasses a range of nebular morphology, including superbubbles, diffuse \\hii\\ regions, and composite objects. The distributions of $L_{\\rm obs}/L_{\\rm SdK}$ for these three subsamples are compared in Figure~\\ref{morph}. On the whole, the superbubbles may have lower ratios of observed to predicted \\Ha\\ luminosities; the mean for the 5 superbubbles in this sample is 0.59, as compared to a mean of 0.89 for the 4 diffuse objects. The mean for the 3 composite objects is 0.85. A preferential leakage of ionizing radiation for the superbubbles might be real if actual holes are present in some of the shell walls, as may be suggested by the morphology of some objects. Such holes could allow the escape of ionizing radiation from those superbubbles. Overall, however, Figure~\\ref{morph} shows that the differing morphologies all span a large range in $L_{\\rm obs}/L_{\\rm SdK}$. In view of possible selection effects and small number statistics, we hesitate at present to attribute any significance to the possible trend with morphology. \\begin{figure*} \\vspace*{-1.9 truein} \\begin{center} \\hspace*{1.0 truein} \\epsfbox{morph.ps} \\end{center} \\vspace*{-15pt} \\caption[]{Distribution of $L_{\\rm obs}/L_{\\rm SdK}$ among the different nebular morphologies. Diffuse \\hii\\ regions are shown with plus symbols, superbubbles with diamonds, and composite objects with squares. There are two points at $L_{\\rm obs}/L_{\\rm SdK} = 0.66$ for the composite objects. \\label{morph}} \\end{figure*} Our results are suggestive that many \\hii\\ regions may be density bounded rather than ionization bounded. As can be seen in Table~\\ref{compare}, this may not be the case for any individual object, but appears likely for many. Density-bounding of the nebulae will probably have a noticeable effect in the observed emission of lower-ionization species such as [\\ion{O}{II}], [\\ion{N}{II}], and [\\ion{S}{II}], that normally dominate in the outer regions of ionization-bounded nebulae (see, \\eg Shields 1990). Emission-line ratios that are based on species originating in different nebular volumes should therefore be interpreted with some caution, if only ionization-bounded models are considered. Photoionization modeling is necessary to fully investigate this effect, which will be explored in a future study. The sample spans over one order of magnitude in \\hii\\ region luminosity, over which there is no obvious correlation between $L_{\\rm obs}$ and fraction of escaping Lyc radiation. Rozas, Beckman, \\& Knapen (1996) have suggested that a change in slope of the \\hii\\ region luminosity function around $10^{38.6} \\ergs$ is due to the density bounding of objects with greater luminosities, whereas the fainter objects are suggested to be primarily ionization bounded. Most of the nebulae in our sample have luminosities below their critical value, and as seen in Table~\\ref{compare}, there is a significant fraction of objects that may be convincingly identified as density-bounded. However, the LMC cannot provide a useful sample of \\hii\\ regions with $\\log L_{\\rm obs} > 38.6$, hence it is impractical to further test the hypothesis with this galaxy. In summary, we find that the \\hii/OB systems in the LMC suggest that up to 51\\% of the ionizing radiation from the hot stars is escaping the local \\hii\\ regions. A significant number of individual \\hii\\ regions reliably appear to be density-bounded. At present, we find no compelling correlation with nebular morphology or luminosity in the fraction of Lyc photons escaping the \\hii\\ regions, although superbubbles might possibly exhibit a greater escaping fraction. Our sample of objects should be fairly representative of the variety of \\hii\\ regions in normal star-forming galaxies, as it includes superbubbles, diffuse nebulae, and complex regions, all with varying luminosities and galactic location. The relative proportion of different nebular types will ultimately depend on the star formation history and environment in any given galaxy. Our estimate of the Lyc escape fraction is consistent with the observed fraction of total \\Ha\\ luminosity emitted by the WIM in the LMC and other star-forming galaxies, and suggests that photoionization by ordinary O stars could indeed be the dominant source of the ionization for the WIM, although alternate ionization mechanisms are not ruled out." + }, + "9708/astro-ph9708089_arXiv.txt": { + "abstract": "Spectral and kinematic studies suggest that the nonthermal radio source Sgr A*, located at the center of the Milky Way, is a supermassive compact object with a mass $\\sim 2-3\\times{10}^6\\msun$. Winds from nearby stars, located $\\approx 0.06$ pc to the east of Sgr A*, should, in the absence of any outflow from the putative black hole itself, be accreting onto this object. We report the results of the first 3D Bondi-Hoyle hydrodynamical numerical simulations of this process under the assumption that the Galactic center wind is generated by several different point sources (here assumed to be 10 pseudo-randomly placed stars). Our results show that the accretion rate onto the central object can be higher than in the case of a uniform flow since wind-wind shocks dissipate some of the bulk kinetic energy and lead to a higher capture rate for the gas. However, even for this highly non-uniform medium, most of the accreting gas carries with it a relatively low level of specific angular momentum, though large transient fluctuations can occur. Additionally, the post-bow-shock focusing of the gas can be substantially different than that for a uniform flow, but it depends strongly on the stellar spatial distribution. We discuss how this affects the morphology of the gas in the inner 0.15 pc of the Galaxy and the consequences for accretion disk models of Sgr A*. ", + "introduction": "Sgr A* may be a massive ($\\sim 2-3\\times 10^6\\;M_\\odot$) point-like object dominating the gravitational potential in the inner $\\la 0.5$ pc region of the Galaxy. This inference is based on the large proper motion of nearby stars (\\cite{HA95}; \\cite{EG97}; \\cite{GEOE97}), the spectrum of Sgr A* (\\cite{MJN92}), its low proper motion ($\\simlt 20 \\kms$; \\cite{B96}), and its unique location (\\cite{LAS91}). The gaseous flows in this region are themselves rather complex, and key constituents appear to be the cluster of mass-losing, blue, luminous stars comprising the IRS 16 assemblage, which is located within several arc seconds ($1^{\\prime\\prime} \\approx$ 0.04 pc in projection at the distance to the Galactic center) from Sgr A*. Measurements of high outflow velocities associated with IR sources in Sgr A West (\\cite{K91}) and in IRS 16 (\\cite{G91}), the $H_2$ emission in the circumnuclear disk (CND) from molecular gas being shocked by a nuclear mass outflow (\\cite{G86}), broad Br$\\alpha$, Br$\\gamma$ and He I emission lines (\\cite{HKS82}; \\cite{AHH90}; \\cite{G91}), and radio continuum observations of IRS 7 (\\cite{YM91}), provide clear evidence of a hypersonic wind, with velocity $v_w \\sim500-1000\\; \\kms$, a number density $n_w\\sim10^{3-4}\\;\\ncm3$, and a total mass loss rate $\\dot M_w\\sim3-4\\times10^{-3}\\mdot$, pervading the inner parsec of the Galaxy. Many of Sgr A*'s radiative characteristics may be due to its accretion of the IRS 16 wind. In the classical Bondi-Hoyle (BH) scenario (\\cite{BH44}), the mass accretion rate for a uniform hypersonic flow is $\\dot M_{BH} = \\pi {R_A}^2 m_H n_w v_w$, in terms of the accretion radius $R_A \\equiv 2 G M / {v_w}^2$. At the Galactic center, with $n_w \\sim 5.5 \\times 10^3 \\ncm3$ and $v_w \\sim700 \\kms$, we would therefore expect an accretion rate $\\dot M_{BH} \\sim 10^{22} \\gms$ onto the black hole, with a capture radius $R_A \\sim .02 \\pc$. Since this accretion rate is sub-Eddington for a one million solar mass object, the accreting gas is mostly unimpeded by the escaping radiation field and is thus essentially in hydrodynamic free-fall starting at $R_A$. Our initial numerical simulations of this process, assuming a highly simplistic uniform flow past a point mass (\\cite{RM94}; \\cite{CM96}) have verified these expectations. On the other hand, the nature of Bondi-Hoyle accretion onto a point like object also presents somewhat of a challenge in understanding what happens to the gas as it settles down into a planar configuration close to the event horizon. Fluctuations beyond the bow shock (located at $\\sim R_A$) produce a transient accretion of net angular momentum that ought to result in the formation of a temporary (albeit small) disk. The circularization radius is $r_c \\approx 2 \\lambda^2 r_g$, where $r_g = 2 G M / c^2$ is the Schwarzschild radius and $\\lambda$ is the specific angular momentum in units of $c r_g$. Our simulations of the BH accretion from a uniform flow suggest that in this configuration $\\langle\\lambda\\rangle\\sim 3-20$. More realistically, the inflow itself carries angular momentum, so that the formation of a disk-like structure at small radii (i.e., $r\\approx 10^{2-3}\\;r_g$) may be difficult to avoid (\\cite{M94}). However, the observations do not appear to favor the presence of a {\\it standard} $\\alpha$-disk at small radii (\\cite{M96}). The current upper limits on the infrared flux from Sgr A* (\\cite{M97}) suggest that either (1) the circularized flow does not form an $\\alpha$-disk, but rather advects most of its dissipated energy through the event horizon (e.g., \\cite{NYM95}), (2) the Bondi-Hoyle flow merges into a massive, fossilized disk, storing most of the deposited matter at large radii (\\cite{FM97}), or (3) Sgr A* is not a point-like object (\\cite{HA95}). Added to this is the fact that in reality the flow past Sgr A* is not likely to be uniform. For example, one might expect many shocks to form as a result of wind-wind collisions within the IRS 16 comples, even before the plasma reaches $R_A$. With this consequent loss of bulk kinetic energy, it would not be surprising to see the black hole accrete at an even larger rate than in the uniform case. The implications for the spectral characteristics of Sgr A*, and thus its nature, are significant. We have therefore undertaken the task of simulating the BH accretion from the spherical winds of a distribution of 10 individual point sources located at an average distance of a few $R_A$ from the central object. As we shall see below, the accretion rate depends not only on the distance of the mass-losing star cluster from the accretor but also on the relative spatial distribution of the sources. As suggested by related work involving linear gradients (\\cite{RA95}), the average value of $\\lambda$ is also larger than that of the uniform case, exhibiting large temporal fluctuations, but still not as large as one might expect. In this {\\sl Letter} we present the results of the first 3D hydrodynamical calculations of multiple-source stellar winds accreted by a point mass. ", + "conclusions": "A variety of accretion scenarios for Sgr A* have been proposed over the years (\\cite{M94}; \\cite{FO94}; \\cite{NYM95}; \\cite{Be96}), each with its own restrictions on $\\dot M$ and $\\lambda$. While the accreted specific angular momentum determined in the present simulations is an order of magnitude too small to support the fossil disk scenario (since then the energy liberated as the wind impacts the fossil disk should be visible; \\cite{FM97}), it is still large enough that any standard $\\alpha$-disk would be easily detectable (\\cite{M94}). The advection dominated disk scenario (\\cite{NYM95}), while permitting a large range of values for $\\lambda$, requires an accretion rate of $\\simlt 10^{-5}\\mdot$ or roughly $0.1 \\dot M_{BH}$, which is 10-20 times smaller than the value derived here. In addition, our $\\dot M$ is more than $10$ times larger than that permitted by the ``mono-energetic'' electron model of Sgr A* (\\cite{Be96}). It appears that additional work is needed to reconcile disk models with the fact that the observed multiple wind sources result in a large mass accretion rate onto the central engine, if its mass is $\\sim {10}^6\\msun$. In view of this, it may not be unreasonable to conjecture that in fact no flattened disk actually forms in Sgr A*, but rather that the excess angular momentum is dissipated in a quasi-spherical configuration and that the thermalized energy is then advected inwards through the event horizon before the gas settles onto a plane (\\cite{Me92}). These simulations suggest that the $\\sim 0.1 \\pc$ region of the Galaxy, centered on the wind sources, is swept clear of gas, leaving a hot, low density gas filling the central cavity. This is consistent with observations of the region within the CND (\\cite{YW93}), and may be acting in concert with other mechanisms to produce the sharp inner edge of the CND (e.g., the abrupt change in gravitational potential; \\cite{D89}). Additionally, a tongue of hot, dense gas has been observed that connects members of the IRS 16 cluster to Sgr A* (\\cite{G91}). It is worthwhile noting that the images in Figures 1 and 2 show ridges of dense gas connecting the sources in the figures to the accretor." + }, + "9708/astro-ph9708150_arXiv.txt": { + "abstract": "\\baselineskip 24pt We report optical imaging and spectroscopy of 41 galaxies in a 22 arcmin square region surrounding Cygnus A. The results show that there is an extensive rich cluster associated with Cygnus A of Abell richness at least 1 and possibly as high as 4. The velocity histogram has two peaks, one centered on Cygnus A, and a more significant peak redshifted by about 2060 km s$^{-1}$ from the velocity of Cygnus A. The dynamical centroid of the spatial distribution is also shifted somewhat to the NW. However, statistical tests show only weak evidence that there are two distinct clusters. The entire system has a velocity dispersion of 1581 km s$^{-1}$ which is slightly larger than other, well studied, examples of rich clusters. ", + "introduction": "The well-known, radio luminous, prototypical FR II radio galaxy, Cygnus A has been known for some time to be embedded in an extensive, cluster-like cloud of hot gas. The gas distribution was first imaged well by the {\\it Einstein} x-ray satellite (Arnaud \\etal \\markcite{arnaud84} 1984). The best high resolution (ROSAT HRI) image was made by Carilli, Perley, \\& Harris \\markcite{carill194} 1994, which shows a complex interaction between the radio jets and lobes of Cygnus A and the x-ray emitting gas in the region immediately surrounding the galaxy. The overall x-ray structure has been interpreted as a cluster-scale cooling-flow, as is usually found in very rich clusters of galaxies. However, only four other galaxies in this region have published redshifts close to that of Cygnus A (Spinrad \\& Stauffer \\markcite{spinrad82} 1982) and thus it has appeared that Cygnus A does not lie in a rich cluster. However, the low galactic latitude (5$^{\\circ}$) of Cygnus A certainly hampers the detection of a cluster because of the dense, confusing galactic star field and the relatively high galactic extinction in this region. Given the apparent anomaly of a cooling flow in a poor group surrounding the most luminous radio galaxy known with a redshift less than 0.1, we thought this problem was worth another look. Below we report the detection of an extensive cluster of optical galaxies surrounding Cygnus A and briefly discuss its properties as currently known. ", + "conclusions": "The Cygnus A radio galaxy lies within a rich, possibly very rich, high velocity dispersion cluster. It appears to be offset from the cluster center both in physical space and in velocity. The existence of a very powerful radio galaxy/AGN in a dense, hot gaseous core is at least very interesting and could well be a key to the origin of such systems. We thank Bill Oegerle for comments on the text. M.J.L. acknowledges partial support from NSF Grant AST-9317596. \\clearpage" + }, + "9708/astro-ph9708220_arXiv.txt": { + "abstract": "We develop a method for interpreting faint galaxy data which focuses on the integrated light radiated from the galaxy population as a whole. The emission history of the universe at ultraviolet, optical, and near-infrared wavelengths is modeled from the present epoch to $z\\approx 4$ by tracing the evolution with cosmic time of the galaxy luminosity density, as determined from several deep spectroscopic samples and the {\\it Hubble Deep Field} (HDF) imaging survey. In a $q_0=0.5$, $h_{50}=1$ cosmology, the global spectrophotometric properties of field galaxies can be well fit by a simple stellar evolution model, defined by a time-dependent star formation rate (SFR) per unit comoving volume and a universal initial mass function (IMF) extending from 0.1 to 125 $M_\\odot$. While a Salpeter IMF with a modest amount of dust reddening or a somewhat steeper mass function, $\\phi(m)\\propto m^{-2.7}$, can both reproduce the data reasonably well, a Scalo IMF produces too much long-wavelength light and is unable to match the observed mean galaxy colors. In the best-fit models, the global SFR rises sharply, by about an order of magnitude, from a redshift of zero to a peak value at $z\\approx 1.5$ in the range 0.12--0.17 $\\sfrd$, to fall again at higher redshifts. After integrating the inferred star formation rate over cosmic time, we find a stellar mass density at the present epoch of $\\Omega_sh_{50}^2\\gta 0.005$, hence a mean stellar mass-to-light ratio $\\gta 4$ in the $B$-band and $\\gta 1$ in $K$, consistent with the values observed in nearby galaxies of various morphological types. The models are able to account for the entire background light recorded in the galaxy counts down to the very faint magnitude levels probed by the HDF. Since only $\\sim 20$\\% of the current stellar content of galaxies is produced at $z>2$, a rather low cosmic metallicity is expected at these early times, in good agreement with the observed enrichment history of the damped Lyman-$\\alpha$ systems. The biggest uncertainty is represented by the poorly constrained amount of starlight that was absorbed by dust and reradiated in the IR at early epochs. A ``monolithic collapse'' model, where half of the present-day stars formed at $z>2.5$ and were shrouded by dust, can be made consistent with the global history of light, but overpredicts the metal mass density at high redshifts as sampled by QSO absorbers. ", + "introduction": "In the past few years two different approaches have been widely used to interpret faint galaxy data (see Ellis 1997 for a recent review). In the simplest version of the ``traditional'' scheme, a one-to-one mapping between galaxies at the present epoch and their distant counterparts is assumed: one starts from the local measurements of the distribution of galaxies as a function of luminosity and Hubble type and models their photometric evolution assuming some redshift of formation and a set of parameterized star formation histories (Tinsley 1980; Bruzual \\& Kron 1980; Koo 1985; Guiderdoni \\& Rocca-Volmerange 1990; Metcalfe \\etal 1991; Gronwall \\& Koo 1995; Pozzetti, Bruzual, \\& Zamorani 1996). These, together with an initial mass function (IMF) and a cosmological model, are then adjusted to match the observed number counts, colors, and redshift distributions. Beyond the intrinsic simplicity of assuming a well defined collapse epoch and pure-luminosity evolution thereafter, the main advantage of this kind of approach is that it can easily be made consistent with the {\\it classical} view that ellipticals and spiral galaxy bulges formed early in a single burst of duration 1 Gyr or less (see, e.g. Ortolani \\etal 1995 and references therein). Because much of the action happens at high-$z$, however, these models predict far more Lyman-break ``blue dropouts'' than are seen in the {\\it Hubble Deep Field} (HDF) (Ferguson \\& Babul 1997; Pozzetti \\etal 1997), and cannot reproduce the rapid evolution -- largely driven by late-type galaxies -- of the optical luminosity density with lookback time observed by Lilly \\etal (1996) and Ellis \\etal (1996). Less straighforward models which include, e.g., a large population of dwarf galaxies that begin forming stars at $z\\approx 1$ (Babul \\& Ferguson 1996), or do not conserve the number of galaxies due to merger events (Broadhurst, Ellis, \\& Glazebrook 1992; Carlberg \\& Charlot 1993) also appear unable to match the global properties of present-day galaxies (Ferguson 1997; Ferguson \\& Babul 1997). A more physically motivated way to interpret the observations is to construct semianalytic hierarchical models of galaxy formation and evolution (White \\& Frenk 1991; Lacey \\& Silk 1991; Kauffmann \\& White 1993; Kauffmann, White, \\& Guiderdoni 1993; Cole \\etal 1994; Baugh \\etal 1997). Here, one starts ab initio from a power spectrum of primordial density fluctuations, follows the formation and merging of dark matter halos, and adopts various prescriptions for gas cooling, star formation, feedback, and dynamical friction. These are tuned to match the statistical properties of both nearby and distant galaxies. In this scenario, there is no period when bulges and ellipticals form rapidly as single units and are very bright: rather, small objects form first and merge continually to make larger ones. While reasonably successful in recovering the counts, colors, and redshift distributions of galaxies, a generic difficulty of such models is the inability to simultaneously reproduce the observed local luminosity density and the zero-point of the Tully-Fisher relation (White \\& Frenk 1991). In this paper we shall develop an alternative method, which focuses on the emission properties of the galaxy population {\\it as a whole}. It traces the cosmic evolution with redshift of the galaxy luminosity density -- as determined from several deep spectroscopic samples and the HDF imaging survey -- and offers the prospect of an empirical determination of the global star formation history of the universe and initial mass function of stars independently, e.g., of the merging histories and complex evolutionary phases of individual galaxies. The technique relies on two basic properties of stellar populations: a) the UV-continuum emission in all but the oldest galaxies is dominated by short-lived massive stars, and is therefore a direct measure, for a given IMF and dust content, of the instantaneous star formation rate (SFR); and b) the rest-frame near-IR light is dominated by near-solar mass evolved stars that make up the bulk of a galaxy's stellar mass, and can then be used as a tracer of the total stellar mass density. By modeling the ``emission history'' of the universe at ultraviolet, optical, and near-infrared wavelengths from the present epoch to $z\\approx 4$, we will shed light on some key questions in galaxy formation and evolution studies: Is there a characteristic epoch of star and metal formation in galaxies? What fraction of the luminous baryons observed today were already locked into galaxies at early epochs? Are high-$z$ galaxies obscured by dust? Do spheroids form early and rapidly? Is there a universal IMF? Let us point out some of the limitations of our approach at the outset. (1) We shall study the emission properties of ``normal'', optically-selected field galaxies which are only moderately affected by dust -- a typical spiral emits 30\\% of its energy in the far-infrared region (Saunders \\etal 1990). Starlight which is completely blocked from view even in the near-IR by a large optical depth in dust will not be recorded by our technique, and the associated baryonic mass and metals missed from our census. The contribution of infrared-selected dusty starbursts to the present-day total stellar mass density cannot be very large, however, for otherwise the current limits to the energy density of the mid- and far-infrared background would be violated (Puget \\etal 1996; Kashlinsky, Mather, \\& Odenwald 1996; Fall, Charlot, \\& Pei 1996; Guiderdoni \\etal 1997). Locally, infrared luminous galaxies are known to produce only a small fraction of the IR luminosity of the universe (Soifer \\& Neugebauer 1991). (2) Our method bypasses the ambiguities associated with the study of morphologically-distinct samples whose physical significance remains unclear, but, by the same token, it does not provide any {\\it direct} information on the processes which shaped the Hubble sequence. Similarly, this approach does not specifically address the evolution of particular subclasses of objects, like the oldest ellipticals or low-surface brightness galaxies, whose star formation history may have differed significantly from the global average (e.g. Renzini 1995; McGaugh \\& Bothun 1994). (3) Although in our calculations the IMF extends from 0.1 to 125 $\\msun$, by modeling the rest-frame galaxy luminosity density from 0.15 to 2.2 \\micron\\ we will only be sensitive to stars within the mass range from $\\sim 0.8$ to about 20$\\msun$. This introduces non-negligible uncertainties in our estimate of the total amount of stars and metals produced. (4) No attempt has been made to include the effects of cosmic chemical evolution on the predicted galaxy colors. All our population synthesis models assume solar metallicity, and thus will generate colors that are slightly too red for objects with low metallicity, e.g. truly primeval galaxies. (5) The uncertanties present in our estimates of the UV luminosity density from the identification of Lyman-break galaxies in the HDF are quite large, and the data points at $z>2$ should still be regarded as tentative. This is especially true for the faint blue dropout sample at $z\\approx 4$, where only one spectroscopic confirmation has been obtained so far (Dickinson 1997). On the other hand, there is no evidence for a gross mismatch at the $z\\approx 2$ transition between the photometric redshift sample of Connolly \\etal (1997) and the Madau \\etal (1996) UV dropout sample. The initial application of this method was presented by Lilly \\etal (1996) and Madau \\etal (1996, hereafter M96). A complementary effort -- which starts instead from the analysis of the evolving gas content and metallicity of the universe -- can be found in Fall \\etal (1996). Unless otherwise stated, we shall adopt in the following a flat cosmology with $q_0=0.5$ and $H_0=50\\,h_{50}\\kmsmpc$. ", + "conclusions": "" + }, + "9708/astro-ph9708016_arXiv.txt": { + "abstract": "Because the baryon-to-photon ratio $\\eta_{10}$ is in some doubt, we drop nucleosynthetic constraints on $\\eta_{10}$ and fit the three cosmological parameters $(h, \\Omega_{\\mathrm{M}}, \\eta_{10})$ to four observational constraints: Hubble parameter $h_{\\mathrm{o}} = 0.70 \\pm 0.15$, age of the universe $t_{\\mathrm{o}} = 14^{+7}_{-2}$ Gyr, cluster gas fraction $f_{\\mathrm{o}} \\equiv f_{\\mathrm{G}}h^{3/2} = 0.060 \\pm 0.006$, and effective shape parameter $\\Gamma_{\\mathrm{o}} = 0.255 \\pm 0.017$. Errors quoted are $1\\sigma$, and we assume Gaussian statistics. We experiment with a fifth constraint $\\Omega_{\\mathrm{o}} = 0.2 \\pm 0.1$ from clusters. We set the tilt parameter $n = 1$ and the gas enhancement factor $\\Upsilon = 0.9$. We consider CDM models (open and $\\Omega_{\\mathrm{M}} = 1$) and flat $\\Lambda$CDM models. We omit HCDM models (to which the $\\Gamma_ {\\mathrm{o}}$ constraint does not apply). We test goodness of fit and draw confidence regions by the $\\Delta\\chi^2$ method. CDM models with $\\Omega_{\\mathrm{M}} =1$ (SCDM models) are accepted only because the large error on $h_{\\mathrm{o}}$ allows $h < 0.5$. Baryonic matter plays a significant role in $\\Gamma_{\\mathrm{o}}$ when $\\Omega_{\\mathrm{M}} \\sim 1$. Open CDM models are accepted only for $\\Omega_{\\mathrm{M}} \\gtrsim 0.4$. The combination of the four other constraints with $\\Omega_{\\mathrm{o}} \\approx 0.2$ is rejected in CDM models with 98\\% confidence, suggesting that light may not trace mass. $\\Lambda$CDM models give similar results. In all of these models, $\\eta_{10}$ $\\gtrsim 6$ is favored strongly over $\\eta_{10}$ $\\lesssim 2$. This suggests that reports of low deuterium abundances on QSO lines of sight may be correct, and that observational determinations of primordial $^4$He may have systematic errors. Plausible variations on $n$ and $\\Upsilon$ in our models do not change the results much. If we drop or change the crucial $\\Gamma_{\\mathrm{o}}$ constraint, lower values of $\\Omega_{\\rm M}$ and $\\eta_{10}$ are permitted. The constraint $\\Gamma_{\\mathrm{o}} = 0.15 \\pm 0.04$, derived recently from the IRAS redshift survey, favors $\\Omega_{\\rm M} \\approx 0.3$ and $\\eta_{10} \\approx 5$ but does not exclude $\\eta_{10} \\approx 2$. ", + "introduction": "\\label{Sec:Introduction} In a Friedmann-Lema\\^\\i tre big bang cosmology, the universal baryonic mass-density parameter $\\Omega_{\\mathrm{B}}\\; (\\,\\equiv 8 \\pi G \\rho_{\\mathrm{B}}/3H_0^2\\,)$ may be calculated from \\begin{equation} \\begin{split} \\Omega_{\\mathrm{B}}\\,h^2 & = 3.675 \\times 10^{-3}(T/2.73\\,\\mathrm{K})^3 \\; \\eta_{10} \\\\ & = 3.667 \\times 10^{-3} \\; \\eta_{10}, \\label{Eq:Omega_B} \\end{split} \\end{equation} where $h$ is defined by the present Hubble parameter $H_0 \\; [\\, h \\equiv H_0/(100$ km s$^{-1}$ Mpc$^{-1})\\,]$, $T$ is the present microwave background temperature, and $\\eta_{10}$ is the baryon-to-photon number ratio in units $10^{-10}$. The last member of equation~(\\ref{Eq:Omega_B}) is obtained by setting $T = 2.728$ K (Fixsen et al. 1996). In principle, $\\eta_{10}$ is well determined (in fact overdetermined) by the observed or inferred primordial abundances of the four light nuclides D, $^3$He, $^4$He, and $^7$Li, if the number of light-neutrino species has its standard value $N_\\nu =3$. For some years it has been argued that $\\eta_{10}$ is known to be $3.4 \\pm 0.3$ (Walker et al. 1991; these error bars are about ``1$\\sigma$''; cf. Smith, Kawano, \\& Malaney 1993) or at worst $4.3 \\pm 0.9$ (Copi, Schramm, \\& Turner 1995a; cf. Yang et al. 1984), and that equation~(\\ref{Eq:Omega_B}) is a powerful constraint on the cosmological parameters $\\Omega_{\\mathrm{B}}$ and $h$. In practice, it seems recently that $\\eta_{10}$ may not be so well determined, and even that the standard theory of big bang nucleosynthesis (BBN) may not give a good fit. With improved abundance data, it appears that the joint fit of the theory to the four nuclide abundances is no longer good for any choice of $\\eta_{10}$ (Hata et al. 1995). These authors offer several options for resolving the apparent conflict between theory and observation. Although they suggest that some change in standard physics may be required (e.g., a reduction in the effective value of $N_\\nu$ during BBN below its standard value 3), they note that large systematic errors may compromise the abundance data (cf. Copi, Schramm, \\& Turner 1995b). The nature of such errors is unclear, and this remains controversial. Other authors have reacted to the impending crisis in self-consistency by simply omitting one or more of the four nuclides in making the fit (Dar 1995; Olive \\& Thomas 1997; Hata et al. 1996, 1997; Fields et al. 1996). This controversy has been sharpened by new observations giving the deuterium abundances on various lines of sight to high-redshift QSOs. These data should yield the primordial D abundance, but current results span an order of magnitude. The low values, D/H by number $\\approx 2 \\times 10^{-5}$ (Tytler, Fan, \\& Burles 1996; Burles \\& Tytler 1996), corresponding to $\\eta_{10} \\approx 7$ in the standard model, have been revised slightly upward [D/H $\\approx (3-4) \\times 10^{-5}$ (Burles \\& Tytler 1997a,b,c); $\\eta_{10} \\approx 5$], but it still seems impossible to reconcile the inferred abundance of $^4$He [Y$_{\\rm P} \\approx 0.234$; Olive \\& Steigman 1995 (OS)] with standard BBN for this large value of $\\eta_{10}$ (which implies Y$_{\\rm BBN} \\approx 0.247$) unless there are large systematic errors in the $^4$He data (cf. Izotov, Thuan, \\& Lipovetsky 1994, 1997). Such low D/H values have also been challenged on observational grounds by Wampler (1996) and by Songaila, Wampler, and Cowie (1997), and deuterium abundances nearly an order of magnitude higher, D/H $\\approx 2\\times10^{-4}$, have been claimed by Carswell et al. (1994), Songaila et al. (1994), and Rugers and Hogan (1996) for other high-redshift systems with metal abundances equally close to primordial. Although some of these claims of high deuterium have been called into question (Tytler, Burles, \\& Kirkman 1997), Hogan (1997) and Songaila (1997) argue that the spectra of other absorbing systems require high D/H (e.g., Webb et al. 1997). If these higher abundances are correct, then D and $^4$He are consistent with $\\eta_{10} \\approx 2$, but modellers of Galactic chemical evolution have a major puzzle: How has the Galaxy reduced D from its high primordial value to its present (local) low value without producing too much $^3$He (Steigman \\& Tosi 1995), without using up too much interstellar gas (Edmunds 1994, Prantzos 1996), and without overproducing heavy elements (cf. Tosi 1996 and references therein)? It appears that $\\eta_{10}$, though known to order of magnitude, may be among the less well-known cosmological parameters at present. Despite this, large modern simulations which explore other cosmological parameters are often limited to a single value of $\\eta_{10} = 3.4$ (e.g., Borgani et al. 1997). In this situation it may be instructive, as a thought experiment, to abandon nucleosynthetic constraints on $\\eta_{10}$ entirely and ask: If we put $\\eta_{10}$ onto the same footing as the other cosmological free parameters, and apply joint constraints on all these parameters based on other astronomical observations and on theory and simulation, what values of $\\eta_{10}$ and the other parameters are favored? This may indicate the most promising avenue to a resolution of the controversy over $\\eta_{10}$. We discuss the following popular CDM models: (1) Open or closed cold dark-matter model with cosmological constant $\\Lambda = 0$ (CDM model). The ``standard'' (flat) CDM model (SCDM), which is an Einstein-de Sitter model, is covered as a special case of this. (2) Flat CDM model with nonzero $\\Lambda$ ($\\Lambda$CDM model). In a flat model with both hot and cold dark matter, with $\\Lambda = 0$ (HCDM model), the constraints will be different; we defer these HCDM models to a later paper. Nonflat models with nonzero $\\Lambda$ are not necessarily ruled out by ``fine-tuning\" arguments and may be of interest (Steigman \\& Felten 1995), but at the moment we are not compelled to resort to these. Our approach will be to let three parameters range freely, fit the constraints (observables) other than nucleosynthetic constraints, test goodness of fit by $\\chi^2$, and draw formal confidence regions for the parameters by the usual $\\Delta\\chi^2$ method. Because statistical results of this kind are sometimes controversial, we intend to keep the work conceptually simple, review the constraints in a helpful way, and discuss our method carefully. Error bars are $\\pm 1\\sigma$ unless stated otherwise. The $\\Delta\\chi^2$ approach is revealing because, in the linear approximation, the confidence regions obtained are rigorous as probability statements and require no ``a priori\" probability assumptions about the unknown parameters. Most of our results are not surprising, and related work has been done before (Ostriker \\& Steinhardt 1995, White et al. 1996, Lineweaver et al. 1997, White \\& Silk 1996, Bludman 1997), but not with these three free variables and the full $\\chi^2$ formalism. It is well known that recent cosmological observations and simulations, particularly related to the ``shape parameter'' $\\Gamma$ and the cluster baryon fraction (CBF), pose a challenge to popular models, and that there is some doubt whether any simple model presently fits all data well. Our work, which begins by discarding nucleosynthetic constraints, provides a new way of looking at these problems. The CBF and $\\Gamma$ constraints have not been applied jointly in earlier work. We find that, given our conservative (generous) choice of error bar on $h$, the SCDM model is disfavored somewhat but by no means excluded, if we are willing to accept $\\eta_{10} \\gtrsim 9$. But even with the generous error on h, and allowing $\\Omega_{\\rm M}$ to range freely, large values ($\\gtrsim 5$) of $\\eta_{10}$ are favored over small values ($\\lesssim 2$). This suggests that the low D abundances measured by Burles and Tytler may be correct, and that the observed (extrapolated) primordial helium-4 mass fraction [$Y_{\\mathrm{P}} \\approx 0.23$; cf. OS and Olive, Skillman, \\& Steigman 1997 (OSS)], thought to be well determined, may be systematically too low for unknown reasons. ", + "conclusions": "\\label{Sec:Conclusions} If BBN constraints on the baryon density are removed (or relaxed), the interaction among the shape-parameter $(\\Gamma)$ constraint, the $f_{\\mathrm{G}}$ (cosmic baryon fraction) constraint, and the value of $\\eta_{10}$ assumes critical importance. These constraints still permit a flat CDM model, but only as long as $h < 0.5$ is allowed by observations of $h$. The $f_{\\mathrm{G}}$ constraint means that large $\\Omega_{\\mathrm{M}}$ implies fairly large $\\Omega_{\\mathrm{B}}$. Therefore the exponential term in $\\Gamma$ becomes important, allowing $\\Omega_{\\mathrm{M}} = 1$ to satisfy the $\\Gamma$ constraint. Values of $\\eta_{10} \\approx 8-15$ are required (Fig.~\\ref{Fig:H-eta_OM1}). The best-fit SCDM model has $h \\approx 0.43$ and $\\eta_{10} \\approx 13$, which is grossly inconsistent with the predictions of BBN and the observed abundances of D, $^4$He, and $^7$Li. For $h > 0.5$ a fit to SCDM is no longer feasible (Fig.~\\ref{Fig:H-eta_OM1}). The SCDM model is severely challenged. The $\\Gamma$ and age constraints also challenge low-density CDM models. The $\\Gamma$ constraint permits $\\Omega_{\\mathrm{M}} < 0.4$ only for high $h$, while the age constraint forbids high $h$, so $\\Omega_{\\mathrm{M}} \\gtrsim 0.4$ is required. Values $\\eta_{10} \\gtrsim 6$ are favored strongly over $\\eta_{10} \\lesssim 2$. The bound $\\Omega_{\\mathrm{M}} \\gtrsim 0.4$ conflicts with the added cluster constraint $\\Omega_{\\mathrm{o}} = 0.2 \\pm 0.1$ at the 98\\% CL, suggesting strongly that there is additional mass not traced by light. Although a few plausible variations on the CDM models do not affect the constraints very much (Figs.~\\ref{Fig:H-Omega_M_L0_var} -- \\ref{Fig:H-eta_L0_G+-0.05}), removing the $\\Gamma$ constraint would have a dramatic effect. Both high and low values of $\\Omega_{\\mathrm{M}}$ would then be permitted. Adopting a smaller observed $\\Gamma_{\\mathrm{o}} \\approx 0.15$ from the IRAS redshift survey also makes a difference. Values $\\Omega_{\\mathrm{M}} \\approx 0.3$ and $\\eta_{10} \\approx 5$ are then favored, but even $\\eta_{10} \\approx 2$ is not excluded. At either low or high density, the situation remains about the same for the $\\Lambda$CDM models (Figs.~\\ref{Fig:H-Omega_M_k0} \\& \\ref{Fig:H-eta_k0}). Because the ages are longer, we can tolerate $\\Omega_{\\mathrm{M}} \\approx 0.3$ for $h = 0.85$. The $\\Lambda$CDM model therefore accepts (barely) the added constraint $\\Omega_{\\mathrm{o}} = 0.2 \\pm 0.1$ at the 7\\% CL, even with the larger $\\Gamma_{\\mathrm{o}} \\approx 0.255$. Improved future constraints on $\\Omega_{\\Lambda}$ will come into play here. Having bounded the baryon density using data independent of constraints from BBN, we may explore the consequences for the light element abundances. In general, our fits favor large values of $\\eta_{10}$ ($\\gtrsim 5$) over small values ($\\lesssim 2$). While such large values of the baryon density are consistent with estimates from the Ly-$\\alpha$ forest, they may create some tension for BBN. For deuterium there is no problem, since for $\\eta_{10} \\gtrsim 5$ the BBN-predicted abundance, (D/H)$_{\\mathrm{P}} \\lesssim 4 \\times 10^{-5}$ (2$\\sigma$), is entirely consistent with the low abundance inferred for some of the observed QSO absorbers (Tytler et al. 1996; Burles \\& Tytler 1996; Burles \\& Tytler 1997a,b,c). Similarly, the BBN-predicted lithium abundance, (Li/H)$_{\\mathrm{P}} \\gtrsim 1.7 \\times 10^{-10}$, is consistent with the observed surface lithium abundances in the old, metal-poor stars (including, perhaps, some minimal destruction or dilution of the prestellar lithium). However, the real challenge comes from $^4$He where the BBN prediction for $\\eta_{10} \\gtrsim 5$, Y$_{\\mathrm{P}} \\gtrsim 0.246$ (2$\\sigma$), is to be contrasted with the \\hii region data which suggest Y$_{\\mathrm{P}} \\lesssim 0.238$ (OS, OSS)." + }, + "9708/hep-ph9708303_arXiv.txt": { + "abstract": "\\noindent We show that the presence of primordial stochastic (hypercharge) magnetic fields before the electroweak (EW) phase transition induces isocurvature fluctuations (baryon number inhomogeneities). Depending on the details of the magnetic field spectrum and on the particle physics parameters (such as the strength of the EW phase transition and electron Yukawa couplings) these fluctuations may survive until the Big Bang nucleosynthesis (BBN). Their lenghtscale may exceed the neutron diffusion length at that time, while their magnitude can be so large that sizable antimatter domains are present. This provides the possibility of a new type of initial conditions for non-homogeneous BBN or, from a more conservative point of view, stringent bounds on primordial magnetic fields. ", + "introduction": " ", + "conclusions": "" + }, + "9708/astro-ph9708002_arXiv.txt": { + "abstract": "We report results of analysis of the ASCA observation of the Low Mass X-ray Binary dipping source XB\\thinspace 1916-053 made on 1993, May 2nd, during which dipping was very deep such that in the deepest parts of dips, the X-ray intensity in the band 0.5 - 12.0 keV fell to zero, demonstrating that all emission components were completely removed. The best-fit orbital period of the binary system determined from the X-ray data was found to be $\\rm {3005 \\pm 10}$~s. The high quality ASCA data allowed spectral evolution in dipping to be systematically investigated by spectral analysis in intensity bands covering the full range of dipping from intensities close to zero to non-dip values. We have shown that the spectra can be well fitted by the same two-component model previously used to give good explanations of the very different dip sources X\\thinspace 1755-338 and X\\thinspace 1624-490, consisting of point-source blackbody emission from the neutron star plus extended Comptonised emission probably from the accretion disk corona. In the case of XB\\thinspace 1916-053 we show that all levels of dipping can be fitted using $\\rm {kT_{bb}}$ = $\\rm {2.14\\pm 0.28}$ keV and power law photon index = $\\rm {2.42\\pm 0.21}$ which are the best-fit values for non-dip data, together with the corresponding non-dip normalisations. Dipping is shown to be due to large increases of column density for the point-like blackbody, combined with the extended power law component being progressively covered by absorber until in the deepest parts of dips, the partial covering fraction approaches unity. This approach differs radically from the ``absorbed plus unabsorbed'' approach previously used in spectral modelling of XB\\thinspace 1916-053 and similar sources, in which the normalisation of the unabsorbed component is allowed to decrease markedly in dipping, behavior generally attributed to the effects of electron scattering. Thus we have shown that spectral evolution in XB\\thinspace 1916-053 can be explained simply in terms of photoelectric absorption without the need for substantial electron scattering. This explanation is supported by calculation of the relative importance of photoelectric absorption and electron scattering in the absorbing region which shows that little electron scattering is expected in the ASCA energy band. ", + "introduction": "XB\\thinspace 1916-053 is an important member of the class of $\\sim $ 10 Low Mass X-ray Binary (LMXB) sources that exhibit decreases in X-ray intensity at the orbital period, generally accepted as being due to absorption in the bulge in the outer accretion disk where the flow from the companion impacts on the outer disk (White and Swank, 1982). XB\\thinspace 1916-053 is unusual in several respects: it has the shortest orbital period of the dipping sources, 50 min (Walter et al. 1982). It has a depth of dipping that is highly variable which at times reaches 100\\%. The source is also notable because of the difference between the X-ray period and the optical period of $\\sim $ 1\\% (Grindlay et al. 1988). \\medskip \\noindent Three Exosat observations were made and the results of Smale et al. (1988) showed that the average extent of dipping changed markedly between the observations as shown by comparing the light curves folded on the X-ray period for the 3 observations. \\begin{figure*}[t] \\leavevmode\\epsffile{fig1.eps} \\caption{ASCA GIS2 light curve for the complete 18~hr observation in the energy band 0.5 - 12.0 keV with 16~s binning. \\label{fig1}} \\end{figure*} Spectral analysis showed that non-dip data for all 3 observations was best-fitted by a simple power law model with photon index $\\Gamma $ close to 1.80. Dip spectra were selected in intensity bands from all 3 observations, for which the best fit was a model consisting of two power laws, each with the index fixed at the above value. The column density of the one component could be held at the quiescent value, while that of the other component increased strongly in dipping; ie there was an absorbed and an unabsorbed term. For the unabsorbed component, the normalisation decreased strongly in dipping. This ``absorbed plus unabsorbed'' approach has been used for several of the dip sources. It has been applied to the sources: XBT\\thinspace 0748-676 (Parmar et al. 1986), X\\thinspace 1254-690 (Courvoisier et al. 1986) and X\\thinspace 1624-490 (Jones and Watson, 1989). It is clear that the parameters of the source emitting regions cannot change coincidentally with dipping and the strong decrease in normalisation of the unabsorbed component has been attributed to the effects of electron scattering in the absorber (Parmar et al 1986; Smale et al. 1988); ie there is a decrease in flux from the source due to scattering which does not reveal itself as low energy absorption. More recently, GINGA data on XB\\thinspace 1916-053 has also been fitted by the absorbed plus unabsorbed approach (Smale et al. 1992; Yoshida et al., 1995). Yoshida et al. showed that the variation in normalisation can be reproduced as absorption by cold matter if electron scattering in the absorber is taken into account. \\medskip \\noindent XB\\thinspace 1916-053 shows increases in hardness during dip ingress. It was originally expected that all of the dipping sources should show hardening during dipping, due to photoelectric absorption of the X-ray spectrum in the absorbing bulge which preferentially removes the lower energy X-rays. However, the dip sources do not in general follow this expectation; some sources show a hardening, but some show a marked softening in dipping, eg X\\thinspace 1624-490 (Church and Balucinska-Church, 1995) which is totally unexpected on simple physical models, ie with a single emission component. Furthermore X\\thinspace 1755-338 has dipping which is independent of energy in the band 1 - 10 keV (White et al. 1984). \\medskip \\noindent Several types of spectral model have been used in fitting the dipping sources. The Exosat spectra of several dipping sources were fitted by a Comptonisation model (White et al. 1988) represented by a single component absorbed cut-off power law. Other workers have used two-component models, notably Mitsuda and co-workers (see for example, Mitsuda et al. 1984). Apart from this difference in approach, individual dipping sources have generally been fitted by different spectral models. In particular, the absorbed plus unabsorbed approach has involved different spectral forms when applied to different sources; ie power law plus power law, cut-off power law plus cut-off power law, etc. More recently, Church and Balucinska-Church (1993, 1995) have proposed a two-component or ``complex continuum'' model which has been able to explain the softening in dipping in X\\thinspace 1624-490 and the energy-independence of dipping in X\\thinspace 1755-338 using the same model. In this model, X-ray emission originates as blackbody radiation from the boundary layer at the surface of the neutron star, plus power law emission representing Comptonisation in an accretion disk corona (valid at energies much lower than the Comptonisation break). In the above sources, the model shows that dipping is primarily due to absorption of the point-source blackbody emission with comparatively little absorption of the extended power law emission. Whether there is softening, hardening or energy-independence depends mostly on the blackbody temperature. In X\\thinspace 1624-490, this is higher than in X\\thinspace 1755-338 ($\\rm {kT_{bb}}$ = 1.39 keV compared with 0.88 keV), so that removal of the blackbody leaves the residual spectrum softer than in non-dip emission. \\vfill \\eject \\noindent The main aim of the present work was to test the complex continuum model in the case of XB\\thinspace 1916-053 using high quality ASCA data, and to test whether this physical model can offer an alternative to the absorbed plus unabsorbed approach. ", + "conclusions": "We have demonstrated that the two-component model can give a good description of spectral evolution in XB\\thinspace 1916-053. In this model, emission originates as point-source blackbody emission from the neutron star plus extended power law emission probably from the accretion disk corona. In the case of XB\\thinspace 1916-053, the source intensity in the band 1.0 - 10.0 keV often actually becomes zero in the deepest parts of dips. In X\\thinspace 1755-338 and X\\thinspace 1624-490, dipping was not 100\\% and spectral evolution during dips was explained by the 2-component model with dipping being primarily due to absorption of the blackbody (Church and Balucinska-Church 1993, 1995). In the case of XB\\thinspace 1916-053, it was necessary to allow the spectral modelling to take account of the fact that the power law component must be totally removed in deepest dipping. The modelling we performed showed that it was not sufficient to use a spectral form AB*BB + AB*PL; the data could not be fitted by this model with normalisations fixed. It was necessary to allow the extended power law component in our model to be progressively covered by the absorber; extended emission would not be covered essentially instantaneously as is the point-source component. With the inclusion of the partial covering term, the two-component model provides very good fits to the spectra. \\medskip \\noindent \\medskip \\noindent The fluxes of the blackbody and power law components in the non-dip data in the energy range 1 - 10 keV are $\\rm {0.98\\cdot 10^{-10}}$ erg $\\rm {cm^{-2}\\;s^{-1}}$ and $\\rm {1.83\\cdot 10^{-10}}$ erg $\\rm {cm^{-2}\\;s^{-1}}$ respectively, so that the blackbody contributes 34\\% to the total energy flux in this band. In X\\thinspace 1755-338 and X\\thinspace 1624-490, in which the two-component model showed that dipping was due primarily to absorption of the blackbody, the spectral evolution in dipping is determined by $\\rm {kT_{bb}}$; in X\\thinspace 1624-490 $\\rm {kT_{bb}}$ was 1.39 keV, the blackbody peaking at $\\sim $4.5 keV, ie it was relatively hard such that the residual power law spectrum when the blackbody was absorbed was softer. XB\\thinspace 1916-053 has even higher $\\rm {kT_{bb}}$ of 2.14 keV; however this does not determine the spectral evolution during dips since both components are absorbed. The low energy cut-off of the spectrum is determined by the power law component, and the hardening observed at dip ingress is clearly simply due to absorption of the low energy part of the spectrum. \\medskip \\noindent Perhaps the most interesting question is whether we expect electron scattering to be important in the absorbing region producing the dips. In other dip sources in which the absorbed plus unabsorbed approach was not used, electron scattering was not thought to be important. Electron scattering may take place in XB\\thinspace 1916-053 between the source regions and the absorbing bulge in the outer disk, but this will occur in both non-dip and dip cases and so is not relevant. In the absorbing region, we can determine the state of ionization by estimating the ionization parameter $\\xi $ as follows, using the column density of the point source blackbody component as a probe of density along a track through the absorbing bulge as dipping develops. As dipping develops, parts of the absorber at different radial positions from the center of the absorber will contribute to attenuation of the incident radiation, having different lengths along the line of sight. We can write $\\xi = {L\\epsilon/N_H r}$ where L is the luminosity, r is radial distance from the source and the thickness of the absorber along the line of sight is a fraction $\\epsilon $ of the accretion disk radius, assumed to fill approximately the Roche lobe of the neutron star. From our best-fit modelling, the unabsorbed flux of the source in the band 1 - 20 keV is $\\rm {4.1\\cdot 10^{-10}}$ erg $\\rm {cm^{-2}\\;s^{-1}}$, and the corresponding luminosity is $\\rm {3.4\\cdot 10^{36}}$ erg $\\rm{s^{-1}}$, using a distance of 8.4 kpc (Smale et al. 1988). As dipping commences, in the intensity band 3.0 - 4.0 c/s, the blackbody has column density $\\rm {4.3\\cdot 10^{22}}$ H atom $\\rm {cm^{-2}}$ and estimating $\\epsilon $ as 0.03, we find that $\\xi $ = 70 erg cm $\\rm {s^{-1}}$. For this value, ionization will not be complete. In deepest dipping, $\\rm {N_H}$ rises to at least $\\rm {6\\cdot 10^{24}}$ H atom $\\rm {cm^{-2}}$, and the thickness of absorber along the line of sight will be maximum, so if we take $\\epsilon $ $\\sim $ 0.3 we get $\\xi $ = 5, such that some elements only can be, at most, singly ionized. \\medskip \\noindent In both cases described above where ionization is not complete, the relative importance of photoelectric absorption and electron scattering is given by the ratio ${\\rm {N_H\\cdot \\sigma _{PE}/ N_e\\cdot \\sigma_{T}}}$, where $\\sigma_{PE}$ is the total photoelectric absorption cross section, $\\sigma_{T}$ is the Thomson scattering cross section, and $\\rm {N_e}$ is the electron column density. This follows from the dependences of the processes on $\\rm {exp -(N_H\\sigma_{PE})}$ and $\\rm {exp-(N_e\\sigma_{T})}$. For a completely ionized plasma of a medium with Solar abundances, it can be calculated that the electron density $\\rm {n_e}$ is related to the ion density $\\rm {n_i}$ via $\\rm {n_e \\; \\simeq \\; 1.2\\cdot n_i}$, since elements other than H contribite more than 1 electron, but with small abundances. Thus in the above cases $\\rm {N_H \\; \\simeq \\; N_e}$ and the above ratio is dominated by the ratio of cross sections. At 1 keV, $\\sigma _{PE}/\\sigma _T$ $\\simeq $ 500, at 4 keV, $\\sigma _{PE}/\\sigma _T$ $\\simeq $ 10, and at 10 keV, $\\sigma _{PE}/\\sigma _T$ $\\simeq $ 2. Thus photoelectric absorption strongly dominates over electron scattering throughout most of the ASCA band. We should also consider how this depends on $\\rm {N_H}$. As the column density increases in deep dipping corresponding to the central regions of the absorber, at $\\rm {N_H}$ = $\\rm {1.5\\cdot 10^{24}}$ H atom $\\rm {cm^{-2}}$, the optical depth for electron scattering becomes unity. However the product $\\rm {N_H \\sigma_{PE}}$ will be very much greater than 1; ie the probability of photoelectric absorption is still much higher than that of electron scattering. If incident photons are allowed to be simultaneously absorbed and scattered, in deep dipping the factor exp-$\\rm {(N_e \\sigma_T)}$ implies that an appreciable fraction of the radiation incident upon the absorber could be scattered. The average column density of $\\rm {35.6\\cdot 10^{22}}$ H atom $\\rm {cm^{-2}}$ from spectral fitting implies a loss of 20\\% by scattering, and it could be argued that the spectral model we have applied should be modified for deep dipping. However if there is a density gradient in the absorber, with $\\rm {N_H}$ = $\\rm {10^{22}}$ - $\\rm {10^{23}}$ H atom $\\rm {cm^{-2}}$ in the outer layers, the optical depth for absorption will be $>$ 1, but for scattering $<$ 1. Thus photons will be preferentially removed before they reach the higher density central regions where scattering would play a part. At 10 keV, this preferential absorption will be much weaker, so that in deepest dipping, electron scattering of the high energy photons will take place. The blackbody emission is by this stage of dipping already highly absorbed leaving flux only in the range just below 10 keV. Thus our value of $\\rm {N_H}$ for the blackbody in deepest dipping of $>$ $\\rm { 600\\cdot 10^{22}}$ H atom $\\rm {cm^{-2}}$ may be an overestimate as some decrease of normalisation may be appropriate for this term, which would imply a smaller $\\rm {N_H}$ value. \\medskip \\noindent Thus, in summary, we have shown that the physical model used by Church and Balucinska-Church to explain the dipping sources X\\thinspace 1755-338 and X\\thinspace 1624-490 also provides a good explanation for XB\\thinspace 1916-053. In all of these cases, the model can be expressed as point source blackbody emission plus extended Comptonised emission which is modified by partial covering. For the first two sources the partial covering fraction f is small, but for XB\\thinspace 1916-053 f becomes large in dipping. Consequently it is possible to explain the dipping in XB\\thinspace 1916-053 purely in terms of photoelectric absorption in a bulge in the outer accretion disk. There is no need for there to be substantial electron scattering, and we have shown that in the ASCA band little scattering is, in fact, expected. The explanation of 3 very different members of the dipping class by the two-component model makes it increasingly likely that it will be able to explain all members of the class." + }, + "9708/astro-ph9708234_arXiv.txt": { + "abstract": "The global structure of optically thin advection dominated accretion flows which are composed of two-temperature plasma around black holes is calculated. We adopt the full set of basic equations including the advective energy transport in the energy equation for the electrons. The spectra emitted by the optically thin accretion flows are also investigated. The radiation mechanisms which are taken into accout are bremsstrahlung, synchrotron emission, and Comptonization. The calculation of the spectra and that of the structure of the accretion flows are made to be completely consistent by calculating the radiative cooling rate at each radius. As a result of the advection domination for the ions, the heat transport from the ions to the electrons becomes practically zero and the radiative cooling balances with the advective {\\it heating} in the energy equation of the electrons. Following up on the successful work of Narayan et al. (1995), we applied our model to the spectrum of Sgr A*. We find that the spectrum of Sgr A* is explained by the optically thin advection dominated accretion flow around a black hole of the mass $\\MBH=10^{6}\\Ms$. The parameter dependence of the spectrum and the structure of the accretion flows is also discussed. ", + "introduction": "Ever since the pioneering studies of steady thin accretion disks by Shakura \\& Sunyaev (1973, hereafter SS), the model of thin accretion disks has been applied successfully to low energy emission from astrophysical objects, such as dwarf novie and pre-main-sequence stars. However, the model has been less successful in modeling of high energy emission from Galactic black hole candidates and active galactic nuclei which are considered to be powered by accreting black holes. The main problem lies in the fact that the thin accretion disks can not reproduce the observed spectra of such systems. The thin accretion disk model which SS originally proposed assumes that the disk is optically thick in the vertical direction and radiates the energy generated by the viscosity locally. This model predicts that the generated spectrum is multi-colored black body, which cannot explain the observed power-law spectra of X-rays from AGNs and Galactic black hole candidates even though it can explain the UV bump of the AGN or the soft state spectrum of the Galactic black hole candidates. More fundamentally, although it is generally believed that QSOs and Seyfert galaxies are powered by the gas accretion onto a super-massive black hole of the mass $\\MBH\\sim 10^8\\Ms$, the thin accretion disks onto such super-massive black holes are too cool to generate high energy photons which are observed in many QSOs and Sayfert galaxies. The model which is investigated by Shapiro, Lightman \\& Eardley (1976, hereafter SLE) is quite attractive in that the accreted gas is optically thin and is much hotter than that in the SS solution and is hot enough to produce high energy photons. SLE considered two-temperature plasma with ions being much hotter than electrons, which enabled quantitative studies of the non-blackbody spectra. SLE model has been applied to explain the spectrum of X-ray binaries and active galactic nuclei successfully. However, it is known that the optically thin hot accretion disk which is considered in SLE is thermally unstable (Piran 1978). If the accretion disk is heated up, then the disk expands and the density decreases, so that the bremsstrahlung cooling rate decreases. The reduced cooling then causes the gas to become even hotter, leading to a runaway thermal instability. For this reason it is not likely that such hot accretion disks exist in real systems for much longer than the thermal timescale. The models introduced so far are local solutions in the sense that the heat generated via viscosity is locally radiated away efficiently, which corresponds to neglecting the advective energy transport in the energy equations. When plasma cannot emit radiation efficiently, the heat generated via viscosity is advected inwards as the internal energy of the plasma. Abramowicz et al. (1988) investigated the effect of advection term in their ``slim disk'' model in the optically thick case, and made it clear that there exists advection dominated branch where the viscous heating is balanced with the advection term rather than the radiative cooling term at high mass accretion rates. The optically thin advection dominated solution at low mass accretion rates is studied intensively by Abramowicz et al. (1995) and Narayan \\&Yi (1995) (see also Ichimaru 1977, Matsumoto et al. 1985). Although they claimed that optically thin advection dominated solution is thermally stable for the long wavelength perturbations, Kato et al. (1996) showed the possibilities of the instability against the short wavelength perturbations. Manmoto et al. (1996) demonstrated that such an instability is favored for explaining rapid X-ray fluctuations from Galactic black hole candidates and does not affect the global stability of the accretion flows. The optically thin advection dominated accretion flows are applied to explain the observed spectra of accreting black holes. Narayan \\& Yi (1995) investigated the self-similar solution which is used later to calculate the spectrum from several low-luminosity accreting systems with great success. As a next step, the calculations of full global steady solutions were awaited for further investigations. Chen et al. (1997) solved optically thin advection dominated solution globally, but their solutions are those of one-temperature plasma and do not include detailed radiation mechanisms. Narayan et al. (1997) derived global solution for optically thin advection dominated accretion flows and showed that the self-similar solution is a good approximation at the radius far enough from the outer and the inner boundaries. This means that the spectra which are derived by using the self-similar solution may be modified, because considerable amount of photons may come from the hot region near the inner boundary. Narayan et al. (1997) did calculate the spectra from two-temperature accretion flows with their global solutions, but they treated the electron energy equations locally, neglecting the effect of the electron advection. Nakamura et al. (1996) were the first to solve the energy equations for ions and electrons and obtained the global two-temperature advection dominated solutions, and showed that the temperature profiles which are crucial to the generated spectra are largely modified when the effect of electron advection is taken into account. However, Nakamura et al. (1996) focused their attentions on the structure of the optically thin accretion disks and did not investigate the spectra from the disks. The study on the spectrum from optically thin advection dominated accretion flows with full global treatment of the basic equations is yet to be done. Thus we are motivated to consistently solve full set of equations including the energy equation for ions and electrons with detailed radiation mechanisms and obtain the spectra from the optically thin accretion flows. In section 2, we present the physical assumptions and the basic equations of our model. We show the results of our calculations in section 3. We then apply our model to Sgr A* (the central core of our Galaxy) in section 4. We conclude in section 5 with a summary and discussion. ", + "conclusions": "In this paper, we have calculated the global structure of optically thin advection dominated accretion flows in the context of two-temperature plasma, adopting the full set of basic equations including the energy equation for the electrons. We have also calculated the spectra emitted by the optically thin accretion flows which we calculated. We have made the calculation of the spectra and that of the structure of the accretion flows to be completely consistent by calculating the radiative cooling rate at each radius by numerically integrating the whole spectrum emitted at the radius. As a result of the advection domination for the ions, the heat transport from the ions to the electrons becomes practically zero and the radiative cooling balances with the advective {\\it heating} of the electrons. This means that the electron cools itself by releasing the stored internal energy as a radiation. Hence the energy equation for the electrons play an important role for the calculation of the spectra, where the temperature profile of the electron is the important factor. An interesting feature of the advection dominated flow, which is known already, is that the azimuthal velocity becomes highly sub-Keplerian and of the same order as the radial velocity and the sound velocity. The point worth noting is that in the innermost hot luminous region, the divergence of the velocities from those in the self-similar solution is fairly large. The accreting gas becomes very hot. The electron temperature even exceeds the rest mass energy of an electron. We have not taken into account the effect of the pair production and the annihilation, which is an important issue. For such hot accretion flows, the synchrotron emission and the Compton scattering are very important. The spectrum is composed by 1) the synchrotron peak which comprises optically thin synchrotron emission and the self-absorbed Rayleigh-Jeans slope and 2) the unsaturated Comptonized photons which forms some bumps and 3) the bremsstrahlung emission and 4)the saturated Comptonized photons. The dependence of the each component on the model parameters is complex. Among them, the position of the Rayleigh-Jeans slope is almost solely determined by the mass of the central black hole. When we make the magnetic field stronger, the temperature of the entire flow decreases, which has significant effect on the Comptonization and the bremsstrahlung emission, but has little effect on the synchrotron emission. When we make the viscosity smaller, the surface density increases and the bremsstrahlung emission increases, but the synchrotron emission and the Comptonization decreases. The simplest relation is the dependence on the mass accretion rate. When we reduce the mass accretion rate, the entire emission is reduced. However the bremsstrahlung emission is much more sensitive to the change of the mass accretion rate than the synchrotron emission. We find that the spectrum of Sgr A* is explained by the optically thin advection dominated accretion flow around a black hole of the mass $\\MBH=10^{6}\\Ms$. Narayan et al. (1995) also calculated the spectrum of Sgr A* using an optically thin advection dominated accretion flow model. Their best fit parameters are different from ours. For instance the mass of the central black hole is $\\MBH=7.0\\times 10^{5}\\Ms$ according to their model. The different points in our model are 1) full global treatment of the basic equations and 2) inclusion of the electron energy equation and 3) calculation of the innermost region where the flow is supersonic and the effects of the redshift are important. 2) and 3) have very important effect on the emitted spectrum, which is not considered in Narayan et al. (1995). We conclude that the X-ray data obtained by various satellite observations are explained by the variation of the mass accretion rate by a factor of $\\sim 2$, if we allow $\\alpha$ to have small value. If we set $\\alpha\\sim0.1$, which is considered to be a standard value for the advection dominated accretion flows, it is not likely that the X-rays come from Sgr A*. We have computed the model using height-integrated equations with fixed structure in the vertical direction. We have also adopted simplified form of equations for the radiation field. Our immediate goal is to solve the basic equations including equations for the radiation field in at least two dimensional space. However, our successful result presented in this paper tells us that the basic idea is correct. The full treatment of Schwarzschild or Kerr metric is also an important issue." + }, + "9708/astro-ph9708144_arXiv.txt": { + "abstract": "On the basis of the best available member list and duplicity information, we have studied the radial distribution of 270 stars and multiple systems earlier than K0 in the Pleiades. Five new long period spectroscopic binaries have been identified from the CORAVEL observations. We have found a clear mass segregation between binaries and single stars, which is explained by the greater average mass of the multiple systems. The mass function of the single stars and primaries appears to be significantly different. While the central part of the cluster is spherical, the outer part is clearly elliptical, with an ellipticity of 0.17. The various parameters describing the Pleiades are (for a distance of 125 pc): core radius $r_{c} = 0\\fdg6$ (1.4 pc), tidal radius $r_{t} = 7\\fdg4$ (16 pc), half mass radius $r_{m/2} = 0\\fdg88$ (1.9 pc), harmonic radius $\\overline{r} = 1\\fdg82$ (4 pc). Low-mass stars (later than K0) probably extend further out and new proper motion and radial velocity surveys over a larger area and to fainter magnitudes would be very important to improve the description of the cluster structure and complete mass function. ", + "introduction": "The unity of star cluster structures has been discussed about thirty years ago by Kholopov (1969) as a generalization of the results obtained first on the basis of star counts, and later of proper motion studies, covering wide areas in several nearby open clusters made by Artyukhina in the Pleiades (Artyukhina 1969; Artyukhina \\& Kalinina 1970), Praesepe (Artyukhina 1966) and other nearby clusters. He clearly showed the existence of a core - halo structure and got a tentative relation between the core and tidal radii. The star count methods give only statistical results because no real knowledge of the member stars is obtained. Results based on proper motion studies are more sound, because cluster members are individually selected. The core - halo structure of open clusters based on the direct identification of individual member stars has so far been established for a small number of open clusters, among them Alpha Per (Artyukhina 1972), the Hyades (Oort 1979), NGC 6705 (Mathieu 1984; Solomon \\& McNamara 1980), NGC 2682 (Mathieu \\& Latham 1986). This is easily explained by the difficulty of identifying the true members in the outer part of the clusters where the density of cluster members is low with respect to the projected density of field stars, i.e. one star per square degree at 2$\\fdg$5 from the center in Praesepe. Due to the considerable area occupied over the sky by nearby clusters with diameters of about 10${\\degr}$, the number of stars to measure for proper motions reaches several times 10$^4$. However, a definitive understanding of the cluster structure and dynamical evolution can only be obtained through a detailed knowledge of the properties of the member stars. Pleiades and Praesepe are, among the nearby clusters, primary candidates to study their overall structure, because of the possibility of obtaining very detailed information on individual stars. Mass segregation and concentration of binary stars towards the cluster centers are predicted by theories and numerical simulations of cluster dynamical evolution (Spitzer \\& Mathieu 1980, Kroupa 1995, de la Fuente Marcos 1996), but observational evidences have so far also been difficult to gather because of the generally too limited cluster surface coverage and the lack of systematic radial velocity survey or detection of binaries in the range of 0\\farcs05 to 1\\farcs0. This traditional gap in binary period coverage is now beginning to be filled by speckle interferometry observations obtained at CHARA (Mason et al. 1993) and direct imaging in the near-IR with adaptive optics (Bouvier et al. 1997). But it is still true that radial velocity surveys of the upper main sequences in the Pleiades, Praesepe and $\\alpha$ Persei clusters are badly needed to improve the knowledge of binarity and orbit characteristics to the level reached by the surveys of the solar-type stars (Duquennoy et al. 1992). Previous studies of the structure of the Pleiades (Artyukhina 1969; Peikov 1990) relied on proper motions to detect the member stars in the outer part of the cluster, but without photometric data to enable an analysis of the membership in the colour-magnitude diagram. van Leeuwen's very comprehensive study was based on photometric observation in the Walraven photometric system and independent new proper motions (van Leeuwen 1983, van Leeuwen et al. 1986). However, the studies of the coronas of the Pleiades (Rosvick et al. 1992a) and of Praesepe (Mermilliod et al. 1990) showed that accurate proper motions, reliable photometry and radial velocities are needed to determine reliable membership estimates for the stars in the very outer part of the clusters. New data are now available for the Pleiades due to the systematic radial velocity survey undertaken by one of the authors (JCM) with the CORAVEL scanner on the lower main sequence, the speckle interferometry survey done at CHARA (Mason et al. 1993; Mason 1996) and the near-IR imaging (Bouvier et al. 1997). This unique material describes for the first time the binary characteristics of the cluster stars and offers the opportunity of studying the radial structure and mass distribution in the Pleiades with much more detail. The first and third papers in this series on the Pleiades (Rosvick et al. 1992a; Mermilliod et al. 1997) have specified the membership of stars in the corona of this cluster. In addition to proper motions and photometry from the literature, we have used radial velocity to improve the selection. We have shown that about 50\\% of the cluster members in the spectral range F5-K0 are located outside the classical area studied by Hertzsprung (1947) and member stars have been found as far as nearly 5$\\degr$ from the center. This makes the real size of the cluster much larger. The improved list of cluster members bears a lot of importance in the determination of the mass function of the Pleiades. Searches for lower mass stars, extending those conducted in the central part (Stauffer et al. 1991, Jameson 1993), will certainly reduce the discrepancies between the field star mass function and those observed in star clusters, usually explained by cluster evaporation. This paper will discuss the new mass function and the radial structure of the Pleiades, with emphasise on mass segregation for the brighter and fainter stars, or between the single and binary members. We first present the results of the radial velocity survey for the central part of the Pleiades and the orbit of a new binary (Sect. 2), discuss the binarity of the upper main sequence (Sect. 3), describe the catalogue of member stars (Sect. 4) and present the results on the cluster structure and the mass function in Sect. 5. ", + "conclusions": "A study of the Pleiades structure has been performed on the basis of the presently available data which limits the sample to stars brighter than $V$ = 12.5. We used the best present knowledge on duplicity in the Pleiades. Using a multi-component analysis applied to the apparent stellar positions we find an ellipticity in the cluster outer part, in agreement with theoretical expectations and van Leeuwen's results (1983). We have observed a clear mass segregation, which depends on the mass of the stars or systems and not on the binary periods. Consequently, binaries are more concentrated than single stars and massive binaries are more concentrated than less massive ones. The mass segregation is significant down to 1 M$_{\\odot}$. Different radii have been computed to characterize the radial distributions of various cluster star populations. For the first time, to our knowledge, the mass function of single stars and primaries of multiple systems have been determined separately and compared. They turned out to be different, in agreement with predictions made by Vanbeveren (1982). Such a result, if confirmed by subsequent studies, may have important implications for star formation models. It shows the extreme interest of detailed studies of young, or very young, open clusters and, especially, of their binary populations. We review the great difficulty in deriving an estimation of the cluster total mass. The mass estimates span an order of magnitude, from $\\sim$500 to 8000 M$_\\odot$, once we consider parameter 1$\\sigma$ errors. Considering only the stars included in our sample we derived a projected spatial mass density of 17.7 M$_\\odot$ pc$^{-2}$, or 9.5 stars pc$^{-2}$, for the 2-pc radius central disk of the cluster. This region was chosen because it is widely used by Kroupa (1995) for comparisons between numerical models. Central density values for different clusters should be compared only if they are computed in the same manner, for example inside this central 2-pc radius disk. There seems to be no other alternative to determine the Pleiades total mass and complete luminosity function than identifying all members in an area even larger than that investigated here and to fainter magnitude. The deepest surveys made in the Pleiades have been limited to the central region within a radius of 3$\\degr$. They should be extended to at least 6$\\degr$ from the center." + }, + "9708/astro-ph9708059_arXiv.txt": { + "abstract": "Some aspects of gravitational lensing by large scale structure (LSS) are investigated. We show that lensing causes the damping tail of the cosmic microwave background (CMB) power spectrum to fall less rapidly with decreasing angular scale than previously expected. This is due to a transfer of power from larger to smaller angular scales which produces a fractional change in power spectrum that increases rapidly beyond $\\ell \\sim 2000$. We also find that lensing produces a nonzero mean magnification of structures on surfaces of constant redshift if weighted by area on the sky. This is a result of the fact that light-rays that are evenly distributed on the sky oversample overdense regions. However this mean magnification has a negligible affect on the CMB power spectrum. A new expression for the lensed power spectrum is derived and it is found that future precision observations the high-$\\ell$ tail of the power spectrum will need to take into account lensing when determining cosmological parameters. ", + "introduction": "Previous discussions of gravitational lensing by large-scale structure have concentrated on calculating the shear and convergence along unperturbed light-paths, i.e. what the geodesics would be were there no fluctuations (e.g. \\cite{seljak}, and references cited therein). Three basic methods have been adopted. The first is by numerical simulation (e.g. \\cite{fuk94}). This method often suffers from limited resolution and overly idealized cosmological models. Another method has been to use a model where light travels freely in a constant background density between clumps of localized mass densities \\cite{fuk94,bess94}. This is not considered to be a realistic cosmological model, because of the wide range of length scales on which galaxy clustering is observed. What appears to be the best method thus far is to take a smooth field of density fluctuations and calculate the shear and convergence along unperturbed light-paths. This can be done with the use of optical scalars \\cite{gunn67,blan91} or equivalently by using methods based on those of \\cite{kais92}. In particular, \\cite{seljak} has applied the techniques of Kaiser \\cite{kais92} to the lensing of the CMB. He found that lensing results in a relatively small smoothing of the CMB power spectrum which makes peaks and troughs somewhat less distinct. This smoothing is due to fluctuations in the magnification of structures on the surface of last scattering. The average magnification was assumed to be zero, as it is to first order. Seljak also found that evolving the deflecting density fluctuations by linear or nonlinear theory makes little difference in the results for $\\ell < 1000$. We show here that deviations of the light-paths from their form in an unperturbed universe result not only in fluctuations in the magnification around a mean of zero, but also a shift in the mean to a positive value. Light-paths are attracted by regions of overdensity and repelled by regions of underdensity. This means that the column density of mass seen by the observer is larger on average than what would be expected using unperturbed light-paths. The predominantly positive second derivatives of the potential in overdense regions produces a shear between light paths which acts to magnify images. At the same time, the average shear between light-paths is, to a lesser extent, reduced by the increase in the density of light paths in overdense regions. The net result is that objects on surfaces of equal redshift or cosmological time will on average appear larger than in an unperturbed universe. The apparent violation of flux conservation can be resolved by realizing that the area of a surface of constant redshift is smaller when light-paths are perturbed. In angular size coordinates, light travels ``slower'' in regions of low potential. The other and more important aim of this paper is to show that after lensing the CMB power spectrum will be enhanced over the unlensed power spectrum at small angular scales or large $\\ell$. Power is transferred upward in $\\ell$ in the damping tail. This result is independent of the existence of a nonzero mean magnification. The paper is organized as follows: In the next section we introduce the formalism used to calculate the lensing effects of LSS. In section 3 it is shown how lensing will change a generic CMB power spectrum. In section 4 the formalism is applied to some specific cosmological models. ", + "conclusions": "We have shown that gravitational lensing can have a significant effect on the CMB power spectrum at small angular scales. The mean magnification will probably be too small to detect, but variations in the magnification will cause the damping tail to decrease less rapidly with increasing $\\ell$. We have also found that nonlinear structure formation and anisotropic contributions the transformation of the power spectrum are important at large $\\ell$. Acoustic peaks at large $\\ell$ may be smoothed to such an extent that they are unidentifiable. The effects of lensing can be removed from the spectrum, but a model for both the lensing potential and the unlensed CMB power spectrum must be assumed. In addition the transformation of the power spectrum is nonlinear although it seems well behaved. This increases the amount of potential information in the damping tail, but makes the interpretation of future small-scale observations more ambiguous. Perhaps the power spectrum of density perturbations will be more tightly constrained by other means. In this paper we have only shown results for flat cosmological models with no cosmological constant, $\\Lambda$. Lensing effects will be somewhat smaller in both low density and $\\Lambda$ models, because of the appearance of the mass density in Poisson's equation, $P_{\\phi}(k,\\tau)=9 a(\\tau)^{-2}\\Omega^2_o \\mbox{H}_o^4 k^{-4} P(k,\\tau)/4$. This factor overcompensates for the increase in path length and growth in fluctuations with lookback time \\cite{seljak}. With reasonable values for $\\Omega_o$ and $\\Lambda$ the change in the spectrum is still significant. Interferometers are under construction that will be capable of probing the predicted CMB fluctuations from 150 to 3500 in $\\ell$. The window size in $\\ell$-space for an interferometer is $\\sim 2\\pi D$ where $D$ is the diameter of the dishes in units of the wavelength. The proposed instruments will operate at around $30$GHz and have dish diameters of tens of centimeters so an $\\ell$-space resolution of $100$ should be achievable. This should allow for many independent measurements of the rate at which the tail falls with $\\ell$. In addition, mosaicing over the sky can further narrow the window. These experiments will use multiple frequencies so that foregrounds such as the Sunyaev-Zel'dovich effect can be removed. We have shown that lensing corrections increase the amplitude by a factor of $\\sim 2$ or more above $\\ell\\sim 3000$ in flat CDM models with Hubble constants in the observed range. With $\\ell$-space windows in the above range any experiment that is capable of detecting the unlensed spectrum at these high $\\ell$'s will be measurably affected by lensing. It will then be essential to include the lensing contribution to the CMB fluctuations in order to utilize the tail beyond $\\ell\\sim 2000$ for cosmological parameter estimation." + }, + "9708/astro-ph9708090_arXiv.txt": { + "abstract": "We develop a prescription for characterizing the strengths of metal lines associated with \\lya forest absorbers (LYFAs) of a given neutral hydrogen column density $N_{\\rm HI}$ and metallicity [O/H]. This {\\em Line Observability Index} (LOX) is line-specific and translates, for weak lines, into a measure of the equivalent width. It can be evaluated quickly for thousands of transitions within the framework of a given model of the \\lya forest, providing a ranking of the absorption lines in terms of their strengths and enabling model builders to select the lines that deserve more detailed consideration, i.e. those that should be detectable in observed spectra of a given resolution and signal-to-noise ratio. We compute the LOX for a large number of elements and transitions in two cosmological models of the \\lya forest at $z \\sim 3$ derived from hydrodynamic simulations of structure formation. We present results for a cold dark matter universe with a cosmological constant; an $\\Omega=1$ cold dark matter model yields nearly identical results, and we argue more generally that the LOX predictions are insensitive to the specific choice of cosmology. We also discuss how the LOX depends on redshift and on model parameters such as the mean baryonic density and radiation field. We find that the OVI (1032 \\AA, 1038 \\AA) doublet is the best probe of the metallicity in low column density LYFAs $(N_{\\rm HI} \\approx 10^{14.5} {\\rm cm}^{-2})$. Metallicities down to [O/H] $\\sim$ -3 yield OVI absorption features that should be detectable in current high-quality spectra, provided that the expected position of the OVI feature is not contaminated by HI absorption. The strongest transitions in lower ionization states of oxygen are OV(630 \\AA), OIV(788 \\AA), and OIII(833 \\AA). These absorption lines are all predicted to be stronger than the OVI feature, but even at redshifts $3-4$ they will have to be observed in the ultraviolet, and they are extremely difficult to detect with present UV instruments, such as the Space Telescope Imaging Spectrograph (STIS). At lower redshifts, detection of these lines may be possible in STIS spectra of the very brightest QSOs, while one may have to wait for next-generation instruments such as the Cosmic Origins Spectrograph (COS) to detect such lines in a number of high-redshift QSOs. The strongest metal lines with restframe wavelength larger than $912 {\\rm \\AA}$ associated with higher column density LYFAs at $z \\sim 3$ are CIII (977 \\AA) and SiIII (1206.5 \\AA), which peak at $N_{\\rm HI} \\sim 10^{17} {\\rm cm}^{-2}$. Of the lines with rest wavelengths $\\lambda_r > 1216 {\\rm \\AA}$, which can potentially be observed redwards of the \\lya forest, the CIV(1548,1551) doublet is expected to dominate in all LYFAs, regardless of the value of $N_{\\rm HI}$. We argue that CIV and CII absorption may peak in different spatial regions, and that comparison of single-phase models of the CIV/CII ionization ratios with observed CIV/CII column density ratios can lead to an overestimate of the ionization parameter in the central parts of the absorbers. ", + "introduction": "Over the past decade, observations have demonstrated that strong \\lya forest absorbers (LYFAs) generally show associated metal line absorption (Meyer \\& York 1987; Lu 1991; Cowie et al. 1995; Womble et al. 1995; Songaila \\& Cowie 1996, hereafter SC). SC find CIV absorption in nearly all LYFAs at $z \\sim 3$ with neutral column densities $N_{\\rm HI}\\simgt 1.6 \\times 10^{15} \\, {\\rm cm}^{-2}$ and in $\\simeq 75 \\%$ of systems with $N_{\\rm HI}\\simgt 3.2 \\times 10^{14}$. The observed CIV column densities are consistent with the absorbers having a mean metallicity ${\\rm [C/H]} \\sim -2.5$ and an intrinsic scatter in metallicity of about an order of magnitude (Rauch et al. 1996; Hellsten et al. 1997). An important question is whether or not a chemically pristine or extremely metal poor population of LYFAs exists, and if so, what the characteristic HI column densities of this population are. For $\\log{N_{\\rm HI}} \\simlt 14.5$ the associated CIV lines are close to the detection limits of SC, and it cannot yet be determined if these absorbers have the same metallicity distribution as those with higher $N_{\\rm HI}$. Some theoretical models of early metal enrichment predict that the mean metallicity declines steadily with decreasing $N_{\\rm HI}$ (Gnedin \\& Ostriker 1997). Considerable progress has been made in the theoretical understanding of LYFAs. From the work of several groups (e.g. Cen et al. 1994; Zhang et al. 1995; Petitjean et al. 1995; Hernquist et al. 1996; Bi \\& Davidsen 1997) a picture has emerged in which the LYFAs are interpreted in terms of ``Gunn-Peterson'' absorption in an inhomogenous IGM pervaded by a background ionizing radiation field, presumed to originate from QSOs and perhaps young, star-forming galaxies. This picture links the \\lya forest directly to cosmological structure formation. The analysis of cosmological models of the LYFAs consists of producing artificial spectra by evaluating the absorption properties of the baryonic IGM along lines of sight through N-body and hydrodynamical simulations of structure formation. These spectra can then be analyzed in much the same way as observed QSO spectra, and the resulting distribution functions in $N_{\\rm HI}$ and linewidths are found to agree quantitatively with those observed, lending considerable support to the cosmological LYFA picture (e.g. Miralda-Escud\\'{e} et al. 1996; Dav\\'e et al. 1997; Zhang et al. 1997). The interpretation of the metal line data within the framework of these models provides a strong test of this picture. If a metal enrichment pattern of the baryonic IGM is specified, the properties of selected absorption lines are readily calculated from the knowledge of densities, temperatures, and UV radiation field along the lines of sight (Haehnelt et al. 1996; Rauch et al. 1996; Hellsten et al. 1997). Until now, the metal lines considered have been limited to OVI(1032,1038), NV(1239,1243), CIV(1548,1551), SiIV(1394,1403), and CII(1335). These lines have been chosen because they have already been identified with LYFAs in real QSO spectra. With the exception of the OVI doublet they all have rest wavelengths $\\lambda _0 > 1216$ \\AA{}, which means that for LYFAs at sufficiently high redshifts, they appear redward of the \\lya forest and hence are relatively easy to detect. Instead of making an a priori selection of a few metal lines to include in a cosmological model for LYFAs, a more satisfactory approach would be to make the models {\\em predict} which metal lines, out of thousands of candidates, should be observable (in spectra of a given resolution and S/N) in LYFAs of different column densities. Such an approach would allow a comprehensive and swift selection of lines that deserve a more detailed treatment, depending on the specific purpose of the modeling. It would provide a sharper test of the model and place more stringent constraints on the metallicities and UV radiation spectrum, by making it possible to compare the complete list of metal lines predicted to be observable to the list of lines actually detected in QSO spectra, as well as the observed relative strengths of these lines. This paper presents and applies a technique for implementing this comprehensive, a priori approach to metal line absorption predictions in cosmological models. In Section 2 we define a line observability index, which can easily be evaluated for a database of metal lines, within a particular model for LYFAs. This index is a measure of column density and hence allows the lines to be sorted by strength. It also predicts which lines should be detectable in LYFAs for a wide range of column densities and metallicities. In Section 3 we apply this technique to two specific cosmological models of the \\lya forest, based on hydrodynamic simulations of structure formation. In Section 4 we discuss the dependence of the LOX on model parameters such as the mean baryon density and radiation field and estimate how it changes with redshift, and we argue that the LOX has only a very weak dependence on the underlying cosmology. We discuss the results and summarize the conclusions in Section 5. ", + "conclusions": "In this paper we have introduced the LOX, a measure that allows a comprehensive selection of the metal lines that are useful for testing the detailed predictions of models for LYFAs. We have applied this method to a particular model, a cosmological $\\Lambda$CDM simulation, and we have quantitatively discussed the dependence of the results on the model parameters $J_{\\nu}$, $\\Omega _b$, and $z$. We have also performed this analysis for a different cosmological simulation, a standard CDM model, but the results differ negligibly from those of the $\\Lambda$CDM model. We find that very few detectable metal lines are associated with the low to moderate density regions of the IGM that give rise to absorbers with $\\log{N_{\\rm HI}} \\simlt 15$. Most of the baryons in the Universe at $z=3$ are believed to reside in this part of the IGM (e.g., Zhang et al. 1997), and the OVI(1032,1038) doublet is found to be the best probe of the metallicity in these regions. Other strong O lines, such as OV(630), OIV(788), and OIII(833) are technically difficult to observe, and have yet to be detected. The CIII(977) and SiIII(1206) lines are the strongest metal lines in high column density LYFAs. The CIV(1548) line is the strongest metal line with $\\lambda _r > 1216 {\\rm \\AA}$ for all LYFAs with $\\log{N_{\\rm HI}} \\simlt 17$. While the specific values of the LOX as a function of column density will depend on redshift and on the adopted cosmological model and radiation field, we have argued that the qualitative conclusions presented above are insensitive to the choice of cosmological model of high redshift LYFAs, and they should hold generally within this scenario. The relatively uncommon absorbers with $\\log{N_{\\rm HI}} \\simgt 16$ should offer the strongest test of cosmological simulations. These systems should have more than a dozen observable metal lines associated with them (for ${\\rm [O/H]} \\sim -2.5$). The relative strengths of these lines will depend on the assumed radiation field and abundance pattern, so a first check would be whether it is possible to make the line strengths in the model spectra match those of lines associated with high column density LYFAs for reasonable values of model parameters. Then one should make a detailed examination of the relative locations in velocity space between absorption components of individual species. A detailed comparison to observations will ultimately determine if the resolution of the simulations is adequate in these high-density regions, and to what extent local dynamical effects associated with the production of metals can be neglected. There are some hints of discrepancies between models and observations for these high column density systems. For example, the models in some cases predict the velocity components of CIV and CII absorption lines to differ, whereas there seems to be some new observational evidence that CII and CIV components are found at the same positions (A. Boksenberg, private communication). Cosmological simulations depict the typical low column density LYFAs as simple structures: low density, smooth, and governed by the straightforward physics of gravity, cosmic expansion, and photoionization. Studies of absorption along parallel lines of sight provide direct observational support for this point of view, especially recent HIRES observations of gravitationally lensed QSOs that set stringent lower limits to the scale of any substructure within low column density absorbers (Rauch 1997; see also Smette et al. 1992, 1995). The simulations do not predict a sudden change in LYFA properties at any particular column density, but they do predict a steady increase of gas density with $N_{\\rm HI}$ (Fig.\\ 1). For ${\\rm log}N_{\\rm HI} \\simgt 16$, the typical absorbers are partially collapsed structures, and gas within them has begun to experience radiative cooling. In this regime, the finite mass and spatial resolution of the simulations may limit the accuracy of the results, and local astrophysical processes such as star formation and supernova explosions might influence the structure of these systems. The standard CDM model studied here produces only about 1/3 of the observed number of Lyman limit systems (${\\rm log}N_{\\rm HI} \\geq 17.2$), even though it matches the abundance of much stronger, damped \\lya absorbers quite well (Katz et al.\\ 1996b; Gardner et al.\\ 1997). While further work is needed to assess the sensitivity of this result to cosmological parameters --- especially the baryon density parameter $\\Omega_b$ --- the discrepancy of numbers suggests that a population of absorbers unresolved by the simulations may become important at column densities ${\\rm log}N_{\\rm HI} \\sim 17$. The high column density LYFAs in the cosmological simulations represent the transition between the weak IGM fluctuations traced by low column density lines and the cold gas concentrations in virialized dark halos traced by damped \\lya absorption. The high-LOX lines listed in Table~1 can reveal many details of the structure and physical conditions in these absorbers, testing the accuracy of the simulations in this regime, providing clues to the nature of any additional absorber populations, and capturing snapshots of gas as it makes its way from the IGM into high redshift galaxies." + }, + "9708/astro-ph9708023_arXiv.txt": { + "abstract": "HIPPARCOS astrometric and kinematical data of oxygen-rich Mira variables are used to calibrate absolute near-infrared magnitudes and kinematic parameters. Two sets of near-infrared magnitudes compiled from different authors are used: broad-band K and narrow-band photometric measurements at 1.04 $\\mu$m (104 filter). Three distinct classes of stars with different kinematics and scale height have been identified. The two most significant groups present characteristics close to the ones usually assigned to extended/thick disk--halo population and old disk population respectively, and thus they might differ by their metallicity abundance. They exhibit different period distributions, as expected if these two groups actually correspond to populations of distinct initial masses, ages and metallicities. Two parallel period--luminosity relations are found in K as well as in 104, one for each significant population. The shift between these relations is interpreted as the consequence of the effects of metallicity abundance on the luminosity. ", + "introduction": "Mira variables, due to their intrinsic brightness and large range of their ages, mark a unique stage in stellar evolution of intermediate-mass stars and thus are important in the study of stellar populations in our Galaxy. Knowledge of their distances is crucial to understand the Galactic structure evolution as well as the pulsational properties of these stars. The existence of infrared and bolometric period--luminosity (PL) relations in the Large Magellanic Cloud (LMC) for Mira variables (see, e.g., Feast et al.\\ 1989) has allowed us to estimate Galactic distances for a large number of Miras (Jura \\& Kleimann 1992; Jura et al.\\ 1993; Alvarez \\& Mennessier 1997). But such works have been limited by the (unavoidable untill recently) assumption on the choice of the zero point for the Galactic Mira PL relation. Now, the release of HIPPARCOS data enables one to proficiently investigate this particular point.\\\\ The results presented in this paper constitute the application of the LM (Luri, Mennessier et al.\\ 1996a; hereafter Paper I) method to HIPPARCOS data concerning oxygen-rich Miras. This method is based on a maximum-likelihood estimation using apparent magnitudes, trigonometrical parallaxes, proper motions and radial velocities. It has been applied to two different samples of about one hundred oxygen-rich Miras for which two sets of near-infrared (K and 104) magnitudes at maximum have been compiled from different authors. These apparent magnitudes complement the kinematical data: the trigonometric parallaxes and proper motions are obtained from the recently available HIPPARCOS Catalogue (ESA 1997) and the radial velocities from the HIPPARCOS Input Catalogue (Turon et al.\\ 1992).\\\\ A preliminary approach to the fundamental problem of absolute magnitudes and distances determinations was made by Luri et al.\\ (1996b) using HIPPARCOS Input Catalogue kinematics data and visual magnitudes. Infrared data are maybe more suitable to study the very cool stars as they emit most of their energy in the infrared. At least in K band, these measurements are less sensitive than visual magnitudes to interstellar and circumstellar extinction. The effect of molecular absorption, which considerably varies in visual region over a cycle, is also weaker in K or 104. Furthermore, the existence of period--luminosity relations is well attested in infrared bands. All these properties contribute to the interest of applying the LM method to near-infrared data. ", + "conclusions": "In this work we have made use of a powerful tool -- the maximum-likelihood LM method -- applied to HIPPARCOS data and near-infrared apparent magnitudes for a large sample of Mira variables. In K as well as in 104, we separate three populations which exhibit different kinematics and exponential scale heights. The two most significant populations can be interpreted as the old disk population and the extended/thick disk--halo population respectively. So they probably differ by their metallicities. They also exhibit clearly distinct period distribution: this important result corroborates that the two populations are certainly composed of stars of different masses, age and metallicity abundance.\\\\ The LM method enables us to derive individual distances. They were compared to the ones derived from visual magnitudes (Mennessier et al.\\ 1997a,c). Distances $r_{\\iK}$ are in good agreement with those obtained with visual data. There is a discrepancy found between the distances $r_{104}$ and the distances $r_{\\iV}$, which might be due to sampling effects.\\\\ Two parallel period--luminosity fit lines are obtained in K as well as in 104 for the Mira variables samples. The slope in K is very similar to the one observed in the LMC (Feast et al.\\ 1989). The Galactic PL relation calibrated by van Leeuwen et al.\\ (1997) lies between our two fits. The shift between our PL relations is probably due to metallicity and mass effects. We stress the necessity to take into account a possible distinction between populations when deriving PL relations for the Galaxy." + }, + "9708/astro-ph9708165_arXiv.txt": { + "abstract": "We have produced deep radio maps of the double quasar {\\sf 0957+561} from multiple-epoch VLA observations. To achieve high sensitivity to extended structure we have re-reduced the best available 1.6~GHz observations and have combined 5~GHz data from multiple array configurations. Regions of faint emission approximately 15\\arcsec\\,north and south of the radio source G are probably lobes associated with the lensing galaxy. An arc 5\\arcsec\\,to the east of G may be a stretched image of emission in the background quasar's environment. 1.4\\arcsec\\,southwest of G we detect a source that we interpret as an image of emission from the quasar's western lobe, which could provide a constraint on the slope of the gravitational potential in the central region of the lens. We explore the consequences of these new constraints with simple lens models of the system. ", + "introduction": "Astrophysicists have anticipated the use of gravitational lensing as an observational tool for 60 years (\\cite{zwicky37a}; \\cite{schneiderbook}), and in the case of the double quasar {\\sf 0957+561} (\\cite{walsh79}), after 18 years of study the promise is closest to fulfillment. If one knew the details of the gravitational-lensing potential, and the time delay among the images of flux-variable components, one could make an estimate, albeit cosmology-dependent, of Hubble's constant H$_0$ (\\cite{refsdal64.2}). Efforts to measure the time delay in this system have converged recently (417$\\pm$3, \\cite{kundic97}; 420$\\pm$13, \\cite{haarsma-th}). However, models of the lensing potential have been less well constrained (\\cite{falco91}; \\cite{kochanek91}; \\cite{grogin96}, 1996b) despite detailed observations of the cluster of galaxies providing the lensing mass (\\cite{young81}; \\cite{angonin94}; \\cite{fischer97}). In an effort to produce a definitive radio map of the object we undertook to re-reduce VLA\\footnote{The VLA is part of the National Radio Astronomy Observatory, which is operated by Associated Universities, Inc. under co-operative agreement with the National Science Foundation.} data gathered by the M.I.T. group, discovering several new features in the field (\\cite{avruch93}). In this letter we present improved maps and identify features that may be useful as model constraints. ", + "conclusions": "To illustrate our interpretation of these new VLA components, we used the {\\tt LENSMOD} software (\\cite{lehar93}) to model the lensing mass with a softened power-law potential (\\cite{blandford87}). The model parameters were: the lens position ($\\Delta\\alpha$, $\\Delta\\delta$), the critical radius ($b$), a core radius ($\\theta_c$), the power index $P$ ($P=1$ is isothermal, while $P=2$ is a Hubble profile), the isodensity ellipticity ($e=1-\\frac{\\rm minor\\;axis}{\\rm major\\;axis}$), and the major axis orientation ($\\phi$). As constraints we used the new HST quasar and G1 positions (\\cite{bernstein97}) and required that the quasar images have a magnification ratio of 0.75$\\pm$0.02 (\\cite{schild90}). We required that any third image of the quasar near G be at least 30 times fainter than B (as a 1$\\sigma$ limit). We also added constraints from the new HST ``blobs'' and ``arc.'' We required that blob2 and blob3 be images of each other, and that the two knots in the arc share a common source. Note that the HST arc is probably caused by the eastern end of the same object that gives rise to blobs 2 and 3, and this could be used to further constrain lens models. To account for the possibility that the HST objects are at a different redshift than the quasar, we added a uniform scale factor $Q_2$ to the deflection angles for those components, as an extra model parameter. The lens model parameters were varied until the source plane position and magnitude differences for each pair were minimized, with a resultant reduced $\\chi^2$ for the fit of 1.1. The best fit model parameters are given in Table~3, with uncertainties determined by varying the model parameters until the reduced $\\chi^2$ increased by 1. Note that the $Q_2$ range corresponds to HST component redshifts of $z_{\\rm HST}\\approx1.3\\pm0.1$ for an $\\Omega=1$ cosmology, which is fully consistent with the quasar and HST objects being at the same redshift. Figure~3 shows the best fit model for $Q_2=1$, with components added to show the modeled radio emission. We do not attempt to account for the VLBI structures (\\cite{garrett94}) in this model, and thus make no claims about the time delay or Hubble's constant based upon our model. We interpret the component GE as the counter-image to the low surface brightness tail of the quasar's western radio lobe E\\@. GE's peak surface brightness and spectral index ($\\alpha^{\\rm 18cm}_{\\rm \\phm{0}6cm} \\sim$ $-$1.0) matches that of component E's northeastern extension, so the brighter parts of the lobe are not multiply imaged. The Bernstein {\\it et~al.\\ }\\,(1996) HST blobs 2 and 3, almost certainly multiple images of a background object, are very close to the positions of GE and the northeast end of E; therefore we expect an image of E near where we have found GE\\@. Because not all of E is multiply imaged, the detailed structure of GE can yield strong constraints on the central region of the lens: either the mass distribution is non-singular, in which case GE comprises two merging images of the eastern end of E, or, if the mass has a central singularity, GE will have a sharp cusp at its western end. High resolution radio observations of GE may be able to distinguish these two possibilities, or at least determine an upper limit on the size of the central mass concentration in G. This is also important because, for a given lens mass, the potential near the quasar B~image is generally deeper for singular models, yielding a longer predicted time delay and thus a lower H$_0$ estimate. The arc-like feature R1 may be a stretched image of background emission. As there is no clear counterpart to the west of G, it is unlikely to be multiply imaged. If the background source is circular, the axial ratio of R1 yields a lower limit of about 5 for its magnification. Jones~{\\it et~al.\\ }(1993), in {\\it Einstein} HRI data, have detected an apparent x-ray arc about 3\\arcsec\\,northwest of R1. The positions are formally consistent, but seem unlikely to be coincident judging from the relative positions of A and B\\@. An association is not ruled out, however. The authors claim the extended x-ray source is an image of thermal emission from a cooling flow in the cluster hosting the lensed quasar at z = 1.41. There are examples of diffuse non-thermal radio emission associated with x-ray-emitting clusters (\\cite{deiss97}), and in this case the lensing magnification may have helped to make it observable. Of course this emission could be foreground; if G has radio lobes, N and S, it could as well have jets. R1 might be back flow from the lobe S, and GN might be a faint jet feeding the lobe N\\@. GE is well explained as an image of the quasar's E lobe, but it's not impossible for it to be the counter-jet of GN, feeding lobe S\\@. The features R2, GN, and GNE are not readily explained by a lensing hypothesis. R2 is in the position of the western half of the HST arc, but all the models we have investigated would produce an eastward extension of this arc which is not detected. We could appeal, {\\it ad hoc}, to source size and spectral index morphology causing the image to be unobservable. The component GN should have a brighter image 5\\arcsec\\,south of G, which is not seen, though we could make the same appeal and note that it might be difficult to separate visually from S\\@. GNE should also have a counter-image to the south of G, which is not seen. However, given the interpretation of R1 as lensed, and the faintness of these features, it is not ruled out that at least some of the emission is due to structure in the background quasar's environment. N and S are certainly not multiply imaged, but whether they are foreground or background is less clear. They could be the radio lobes of the galaxy G. At the lens redshift ($z = 0.36$, and assuming $\\Omega=1$, $h=0.75$) N and S would have a (projected) proper separation of 120 kpc, and luminosities at 178~MHz of about 10$^{24}$ WHz$^{-1}$, typical values for low power, limb darkened radio galaxies. The optical classification of G as a cD galaxy, and the fact that N and S are aligned within 30\\arcdeg\\, of its optical minor axis are also consistent (\\cite{miley80}). N and S might be old lobes of the background quasar, in which case the numbers are 170 kpc and 10$^{26}$WHz$^{-1}$, more appropriate for powerful, limb brightened sources. If N and S are associated with the quasar, the relatively small lobe separation (56 kpc) and the high core-to-lobe flux ratio ($R=0.22$) suggest that the jet axis is moderately inclined towards the line-of-sight (\\cite{muxlow91}). This inclination readily explains the seemingly large rotation of the jet from the axis defined by N and S to that defined by C and E. The performance of the VLA at $\\lambda$18cm has improved markedly since 1980, and new observations should detect or exclude these features with high significance. We are aware of a very deep VLA observation (\\cite{harvanek96}) at $\\lambda$18cm and $\\lambda$3.6cm; the longer wavelength data should be able to confirm GE, GN, and GNE, and if GE is detected at $\\lambda$3.6cm it may be possible to determine whether the mass model is singular, or whether GE consists of two merging images." + }, + "9708/astro-ph9708171_arXiv.txt": { + "abstract": "We present data on the monitoring of the Galactic X-ray transient GRS\\,1915+105 at 15 GHz with the Ryle Telescope. We have found quasi-periodic oscillations with periods in the range 20--40 min which are tentatively associated with the soft-X-ray variations on the same time-scale. The overall behaviour of the radio emission is shown to vary in a strong association with the X-ray emission as recorded by the {\\it RXTE} all-sky monitor. ", + "introduction": "Castro-Tirado et al.\\ (1992) reported the discovery of the hard-X-ray transient source GRS\\,1915+105, using the WATCH instrument on the {\\it GRANAT} satellite. Its X-ray emission at both high and low energies has proved to have a rich structure -- see, e.g., Paciesas et al.\\ (1996) and Greiner et al.\\ (1996). In the radio regime, the source is no less remarkable; Mirabel \\& Rodr\\'\\i guez (1994) discovered a double-sided relativistic ejection of radio-emitting material. Adopting a model based on symmetrical ejection, they derive an angle to the line of sight of 70 degrees and a velocity of ejection of 0.92$c$. The distance, from 21-cm H\\,{\\sc i} absorption, is estimated to be 12.5 kpc, consistent with the relativistic-expansion model. Rodr\\'\\i guez et al.\\ (1995) and Foster et al.\\ (1996) have presented flux-density monitoring data at a range of frequencies. From this monitoring, it was apparent that the flux density in the radio regime varies on many time-scales, and we started a systematic monitoring program at 15 GHz with the Ryle Telescope in 1995 August. One surprising feature, apparently periodic oscillations with periods in the range 20--40 min, has been reported in two IAU Circulars (Pooley 1995, 1996). This phenomenon, and other patterns of variation including the relationship to the {\\it RXTE} monitoring data, are considered in more detail in this paper. ", + "conclusions": "A wide range of previously unknown phenomena have been recorded in the high-frequency radio emission from GRS\\,1915+105. A strong link has been established between the radio and X-ray emission. Fender et al.\\ (1997) also report emission in the infrared varying on similar time-scales to those of the radio oscillations reported here. They present evidence that each oscillation is associated with an ejection event, and interpret the infrared flux as the high-frequency tail of a synchrotron spectrum. Belloni et al. (1997) have shown that the soft-X-ray dips on timescales near 30 min can well be explained by the removal of the inner 200 km of the accretion disc. The coincidence of the rise in the radio flux density with the X-ray dip in Fig. 5(b) therefore suggests that at least part of the inner disc is ejected from the system during the oscillations. Important observations remain to be made, including simultaneous observations with high time resolution at as many frequencies as possible. We suggest also that the major radio outburst starting near MJD 50275, and the simultaneous nearly-constant X-ray flux, may have been the results of the removal of a larger part of the inner accretion disc." + }, + "9708/gr-qc9708070_arXiv.txt": { + "abstract": " ", + "introduction": "The observable quantum gravity effects are usually considered as a subject of study for a very remote future. In any case, they are normally expected to take place on microscopic, rather than macroscopic, scales. It is remarkable that we are probably facing these effects right now, and on extremely large scales, in the form of very long-wavelength cosmological perturbations responsible for the observed~[1] large-angular-scale anisotropies in the cosmic microwave background (CMB) radiation. It is also likely that we will have direct access to these effects on smaller scales, and in the near future, by detecting the relic background of (squeezed) gravitational waves~[2] with the help of gravity-wave detectors currently under construction or in the phase of the design study (for a recent review, see~[3]). Since the mentioning of the subject of quantum gravity often instigates the feeling of disbelief and panic among the audience, we should first clarify our intentions. We are not going to quantize, create and destroy, universes. We will be discussing a conceptually simpler problem of quantizing cosmological perturbations in the Universe. This approach is quite analogous to, say, quantization of excitations in a sample of a condensed matter, or to the quantum field description of the squeezed light generation in quantum optics experiments~[4,5]. The cosmological perturbations (density perturbations, rotational perturbations, gravitational waves) consist of a gravitational field component and a matter component. The gravitational field and matter variables are governed together by a system of coupled differential equations~[6]. In case of gravitational waves, only the gravitational field component is present. An attempt of quantization of the matter part of cosmological perturbations without quantizing the associated gravitational field part would be as inconsistent as an attempt of quantization of the electric components of the electromagnetic field without quantizing the magnetic components. The central point of our effort is realization of the fact that even if there were no cosmological perturbations initially, they must have been generated later, under quite general conditions, as a result of parametric (superadiabatic) amplification of the zero-point quantum oscillations of the perturbation field~[7]. The source of amplification, which ``pumps'' energy into the zero-point (vacuum) fluctuations, is the strong variable gravitational field of the very early Universe. Although we are quantizing only the perturbed part of the gravitational field, we are still quantizing gravity. We expect our results to depend explicitly on all the three fundamental constants: $G$-representing gravity, $c$-representing relativity, and $\\hbar$-representing quantum theory. Specifically, the $G$ and $c$ participate in the definition of the gravitational energy-momentum tensor, and the $\\hbar$ enters because of our quantum normalization of the energy of each mode to ${1\\over 2}\\hbar\\omega$, {\\it i.e.,} to a ``half of the quantum in each mode''. These constants combine naturally in the Planck length $l_{Pl} =\\sqrt{G\\hbar/c^3}$ or in the Planck mass $m_{Pl}=\\sqrt{\\hbar c/G}$. The vacuum fluctuations of the field in question can be visualized as oscillations with a nonzero amplitude proportional to $\\sqrt{\\hbar}$. Since everything in this formulation of the problem depend on the presence of the vacuum fluctuations, our results must vanish when $\\hbar$ is sent to zero (in other words, no quantum theory --- no nonzero effects). It is important to know why the Planck constant $\\hbar$ participates in our expressions, and it is helpful to write and keep track of all the fundamental constants. The unfortunate practice of setting outright $G=c=\\hbar =1$ is sometimes an indication that the author is not quite sure why and which fundamental constants should be present, and whether, say, the Planck constant should appear in the nominator or in the denominator of the final expression. So, for safety, so to say, all the fundamental constants are being set to 1. Similarly, the common practice of ``measuring'' some physical quantity (for instance, a scalar field) in units of the Planck mass is not at all an element of quantum physics, like the possibility of expressing the distance $L$ between the Earth and the Moon in units of the Planck length, $L\\approx 4\\times 10^{43}l_{Pl}$, is not an indication of quantum nature of that distance. In our study, the fundamental constants $G$, $c$, $\\hbar$ are present because we are quantizing the (perturbed part of) gravitational field and the associated matter fields (unless we are considering specifically gravitational waves in which case only gravity is being quantized). However, the $G$, $c$, $\\hbar$ can also participate in the final expressions if a discussed problem does not belong to the realm of quantum gravity, but involves gravity plus quantum theory. A beautiful example is provided by the expression for the maximal masses and radii of cold white dwarfs and neutron stars. These objects owe their existence to the interplay between gravity and quantum mechanics. It is the pressure of the degenerate Fermi gas which through the Pauli exclusion principle enables the cold white dwarfs and neutron stars to resist gravitational collapse~[8]. If there were no quantum mechanics $(\\hbar =0)$, these objects would not exist. This statement in words is of course supported by formulae. The maximal masses of cold white dwarfs (the Chandrasekhar mass) and neutron stars are approximately equal to \\[ M_{\\rm max} \\approx m_{Pl} \\left( {m_{Pl}\\over m_B}\\right)^2 \\approx 1.5M_{\\odot} \\] where $m_B$ is the mass of a baryon. The equilibrium radius of a white dwarf with the mass $M_{\\rm max}$ is approximately equal to \\[ R \\approx {\\hbar \\over m_e c}\\left( {m_{Pl}\\over m_B} \\right) \\approx 5\\times 10^8~{\\rm cm} \\] where $m_e$ is the mass of an electron. These formulae explicitly demonstrate that $M_{\\rm max}$ and $R$ would vanish if the Planck constant were sent to zero. It is important to note that the sign of approximate equality in these formulae means the approximate equality. It is not a sign of proportionality, which would allow many other dimensional quantities to be present, and it is not a sign of an estimate based on the dimensional analysis only, which would allow the numerical coefficient to be, say, 50 times larger or 50 times smaller than 1. A beautiful physics is usually expressed in simple and beautiful formulae: the final result depends essentially on the fundamental constants and a couple of parameters, such as $m_B$ and $m_e$, characterizing the system. As we will see below, our predictions for quantum-mechanically generated cosmological perturbations in simple cosmological models do also depend essentially on the fundamental constants, combined in $l_{Pl}$, and a couple of parameters characterizing the model. The inspection of positions and powers in which the fundamental constants enter the claimed result may serve as a first test of whether we are prepared to believe the result. For instance, when you read that ``inflation generates'' density perturbations because it ``stretches the vacuum fluctuations beyond the horizon'', plus accompanying words about something quantum becoming something classical ``upon the horizon crossing'', this explanation may raise your doubts even before any further reading. If everything is so simple, why do we consider vacuum fluctuations of the ``inflaton'', and not the vacuum fluctuations of the electromagnetic field which we definitely know to exist? Since the vacuum fluctuations of the electromagnetic field are being ``stretch beyond the horizon'' in exactly the same manner, we should apparently be able to generate enormous amount of photons as well. However, the result for photons is known: strict zero, independently on whether the expansion is inflationary or not. This example simply shows that it is the dynamical property of the system, that is, whether and how the vacuum fluctuations are coupled to the pumping field, and not the universal kinematical property of ``stretching'' the wavelength ``beyond the horizon'', which is really important. As a first test of the inflationary predictions you may wish to check whether and how the Planck constant enters the claimed result. Surprisingly, you will not find the Planck constant, or when it is present in the form of the Planck mass, it appears in the denominator, and not in the nominator, of the final expression. Taken for the face value, this position of the Planck constant would suggest that the effect, owing its existence to the quantum fluctuations, goes to infinity when the Planck constant is sent to zero. Of course, the correct position of the fundamental constants is not a guarantee that the result is correct numerically. And the lack, or a strange position, of the fundamental constants in the final expression is not an indication that the result is necessarily wrong. The author could have set $G=c=\\hbar=1$ everywhere, or could have done this in one part of the participating equations but forgot to do the same in the other. However, an unclear physical explanation of the expected effect and a chaotic placing of the fundamental constants in the resulting expression does raise a suspicion that the numerical result may also be wrong. We will return to this point later. ", + "conclusions": "" + }, + "9708/astro-ph9708037_arXiv.txt": { + "abstract": "We present results from an optical--infrared photometric study of early--type (E+S0) galaxies in 19 galaxy clusters out to $z = 0.9$. The galaxy sample is selected on the basis of morphologies determined from {\\it HST} WFPC2 images, and is photometrically defined in the $K$--band in order to minimize redshift--dependent selection biases. Using new ground--based photometry in five optical and infrared bands for each cluster, we examine the evolution of the color--magnitude relation for early--type cluster galaxies, considering its slope, intercept, and color scatter around the mean relation. New multiwavelength photometry of galaxies in the Coma cluster is used to provide a baseline sample at $z \\approx 0$ with which to compare the distant clusters. The optical--IR colors of the early--type cluster galaxies become bluer with increasing redshift in a manner consistent with the passive evolution of an old stellar population formed at an early cosmic epoch. The degree of color evolution is similar for clusters at similar redshift and does not depend strongly on the optical richness or x--ray luminosity of the cluster, suggesting that the history of early--type galaxies is relatively insensitive to environment, at least above a certain density threshold. The slope of the color--magnitude relationship shows no significant change out to $z = 0.9$, providing evidence that it arises from a correlation between galaxy mass and metallicity, not age. Finally, the intrinsic scatter in the optical--IR colors of the galaxies is small and nearly constant with redshift, indicating that the majority of giant, early--type galaxies in clusters share a common star formation history, with little perturbation due to uncorrelated episodes of later star formation. Taken together, our results are consistent with models in which most early--type galaxies in rich clusters are old, formed the majority of their stars at high redshift in a well--synchronized fashion, and evolved quiescently thereafter. We consider several possible effects which may be introduced by the choice of morphologically recognizable elliptical and S0 galaxies in dense environments as a subject for study. In particular, the inclusion of S0 galaxies, which might be undergoing morphological transformation in clusters as part of the Butcher--Oemler effect, may influence the results of our investigation. ", + "introduction": "The elliptical galaxy formation scenario proposed by, e.g., Eggen, Lynden--Bell, \\& Sandage (1962, hereafter ELS), Searle, Sargent, \\& Bagnuolo (1973), and Tinsley \\& Gunn (1976) postulates a single burst of star formation at high redshift, followed by passive stellar evolution. An elliptical galaxy is assumed to form the vast majority of its stellar mass during the initial starburst. Several observations suggest that early--type galaxies in present--day clusters may have formed and evolved in this fashion. For example, the color--magnitude ($c-m$) relation seen in nearby clusters can be readily explained in terms of a single star formation episode. More massive galaxies would retain supernovae ejecta more effectively, resulting in higher metallicities for the succeeding generations of stars within the initial burst, and hence in redder colors for more luminous galaxies (Larson 1974; Arimoto \\& Yoshii 1987; Franx \\& Illingworth 1990). Indeed, the $c-m$ relation in nearby clusters evidently implies a close tie between metallicity and galaxy mass, as seen in the tightness of the Mg$_2$--$\\sigma$ correlation (Bender, Burstein, \\& Faber 1993). Furthermore, the scatter in the $UVK$ colors of present--epoch cluster E+S0s is observed to be very small (Bower, Lucey, \\& Ellis 1992; Eisenhardt et al.\\ 1997). This has been used to argue for a high degree of synchronization in their star formation histories, and for relatively old ages. Late, episodic bursts of star formation, young galaxy ages, or a wide range in galaxy formation redshifts would be expected to lead to larger color scatter than is observed. An early formation epoch coupled with passive evolution simply predicts the observed homogeneity in the colors of most elliptical galaxies in clusters today. There is, however, spectroscopic evidence for younger stellar populations in some elliptical and S0 galaxies (O'Connell 1980; Worthey 1996), as well as morphological signs of disturbance attributed to past merger events (Toomre 1978; Schweizer \\& Seitzer 1988). Moreover, Schweizer et al.\\ (1990) and Schweizer \\& Seitzer (1992) have shown that the spectroscopic and photometric indicators of younger starlight correlate with the degree of morphological disturbance in E and S0 galaxies, implying a connection between the two phenomena. The galaxy samples considered in such studies have consisted primarily of {\\it field} galaxies, and have not included galaxies in the cores of rich clusters such as Coma. It is therefore unclear whether these results apply to all early--type galaxies or are a function of their environment. The traditional picture of massive galaxy formation in a monolithic collapse episode at high redshift does not sit comfortably within the context of modern hierarchical merging scenarios for galaxy formation and evolution, such as those based on the cold dark matter (CDM) model. Such models predict that massive galaxies form late, at $z \\le 2$, from the gradual merging of smaller galaxies (e.g., White \\& Frenk 1991; Kauffmann et al.\\ 1993; Cole et al.\\ 1994). Although this scenario predicts a wide range of ages for elliptical galaxies, their small color scatter in the present epoch can be accomodated because sufficient time elapses since the epoch of extensive merging for the resulting color variations to damp out (Kauffmann 1996; see also Schweizer \\& Seitzer 1992). The mass--metallicity correlation implied by the slope of the color--magnitude relation has been recently examined in the context of hierarchical merging models by Kauffmann \\& Charlot (1997). Using a multi--metallicity spectral synthesis code and semi--analytical galaxy evolution models, they are able to reproduce this correlation by forming more massive ellipticals from mergers of more massive progenitor disk galaxies, which themselves are better able to retain metals during star formation. The aforementioned observational constraints on E/S0s are well--known only at $z = 0$, which allows a wide range of possible star formation histories, including both the ELS and CDM scenarios. Recent observational advances both on the ground and in space are beginning to provide detailed information on the properties of early--type galaxies at higher redshift, e.g., through study of the Fundamental Plane and its projections (Dokkum \\& Franx 1996; Dickinson 1995 and 1997; Pahre et al.\\ 1995), and the Mg--$\\sigma$ relation (Ziegler \\& Bender 1997). By examining the properties of galaxies in clusters at high redshifts, constraints on the nature of early--type galaxy evolution might be developed so as to favor one formation scenario over the other. In practice, several complications have arisen in investigations of galaxy evolution in distant clusters. Photometric studies based on optical imaging (e.g., Butcher \\& Oemler 1978; Dressler \\& Gunn 1992; Rakos \\& Schombert 1995; and Lubin 1996) may be hampered by selection effects due to the redshifting of blue and ultraviolet rest frame wavelengths into the observed bands. This is a significant concern for galaxy selection in the most distant known clusters from optical and x--ray samples ($z \\approx 0.9$), where even the $I$--band measures blue wavelengths in the cluster rest frame. Selecting galaxy samples in the near infrared alleviates this problem. Even at $z \\sim 1$, the $K$--band measures galaxy light emitted in the rest frame near--IR. Because the infrared spectral energy distributions of all but the most vigorously star forming galaxies are very similar (e.g.\\ Johnson 1966), a galaxy sample defined in the near--IR should be free of redshift--dependent biases regarding galaxy type out to $z \\approx 1$ and beyond. Optical--to--infrared photometry, particularly where the optical band measures light emitted shortward of 4000 \\AA\\ in the rest frame, ensures a long wavelength baseline for color measurements, providing in effect a measurement of the luminosity ratio of main sequence stars to evolved red giants in a galaxy's stellar population (e.g. Bruzual \\& Charlot 1993). Arag\\'on--Salamanca et al.\\ (1993) used this approach to study color evolution in a study of 10 $z \\ge 0.5$ clusters, and found that the modal optical--IR colors of the galaxies, presumed to be primarily cluster ellipticals, become bluer with redshift. Without the ability to distinguish between morphological classes of galaxies in distant clusters, purely ground--based photometric studies run the risk of mixing galaxy types. This results in a particular hazard for studying the evolution of early--type galaxies, since the increasing proportion of blue, late--type galaxies in distant clusters (the ``Butcher--Oemler effect'') may introduce greater contamination at higher redshifts. Imaging data from the Hubble Space Telescope ({\\it HST}) largely solves this problem by providing the capability to select galaxies purely by morphology (Dressler et al.\\ 1994; Couch et al.\\ 1994). In Stanford, Eisenhardt \\& Dickinson (1995; henceforth SED95) we made a first attempt to combine multi--band optical--IR photometry with {\\it HST} imaging of two clusters to evaluate the evolutionary state of early--type cluster galaxies at $z \\approx 0.4$. In that paper, the rest--frame $V-H$ colors of distant, morphologically--selected early--type galaxies were found to be similar ($< 0.2$ mag bluer) to those at $z = 0$, and the slope and scatter of the color--magnitude relation for those galaxies were also found to be indistinguishable from the present--day values. Ellis et al.\\ (1997) have recently made a similar test using multi--color optical WFPC2 data to study E+S0 galaxies in three clusters at $z \\approx 0.54$, with similar results. In the present paper, we extend the analysis of SED95 to a much larger sample of galaxy clusters, spanning a wide redshift range ($0 < z < 0.9$), using substantially improved optical--IR photometric data. For $\\Lambda = 0$ cosmologies with $0.05 < q_0 < 0.5$, our cluster sample spans 50--62\\% of the lookback time to the Big Bang, affording us a broad view of the evolution of early--type galaxies throughout the second half of the universe's lifespan. Our major results are presented here; the photometric data set itself will be presented in a supplemental paper (Stanford et al.\\ 1997). Except where noted, the assumed cosmology is $H_0 = 65$ km s$^{-1}$ Mpc$^{-1}$, $q_0 = 0.05$, and $\\Lambda = 0$, which results in a present--day age for the Universe of 13.5~Gyr. In this paper, we use the term ``early--type'' galaxy to refer to those galaxies classified morphologically as having Hubble classes E, E/S0, or S0. We have made no attempt here to further subdivide the early--type galaxy population in high redshift clusters, and in particular no effort has been made to separate elliptical from S0 galaxies. Making such morphological distinctions is sometimes difficult even for nearby galaxies, and is dependent on orientation: a face--on S0, for example, can be hard to distinguish from a ``true'' elliptical galaxy. It has yet to be established how well this subclassification can be achieved for WFPC2 images of high redshift galaxies (cf. Smail et al.\\ 1997 for a discussion), and for the purposes of this paper we have preferred to avoid the resulting uncertainties which might arise from misclassification. However, grouping elliptical and S0 galaxies together for an evolutionary study such as this one may have undesirable consequences for the interpretation of the results. We will return to this point in \\S4. Throughout the text, except where we wish to make the distinction between elliptical and S0 galaxies explicit, we will generally use the terms ``early--type'' and E+S0 interchangeably to refer to the overall population of elliptical and S0 galaxies in a cluster. In addition, we note here that the subject of our investigation is {\\it giant} galaxies, not dwarf ellipticals or ``spheroidals.'' As described in \\S2, our photometric data and the analysis thereof are limited to galaxies not more than $\\sim 2$ magnitudes fainter than present--day (unevolved) $L^{\\ast}$ at the redshift of each cluster examined. For objects with the typical colors of E+S0s, we are therefore discussing galaxies with $M_B \\simlt -18.5$ (for our adopted cosmology). Finally, the results reported here concern only cluster galaxies, and do not necessarily bear on the evolution of field ellipticals. Moreover, because we use WFPC2 images which cover only the core regions of the distant clusters, we cannot address the evolution of galaxies located at larger cluster--centric radii. Some studies of nearby clusters indicate that early--type galaxies located outside of the core regions may have experienced recent star formation (Caldwell et al.\\ 1993; Caldwell \\& Rose 1997). Almost all of these objects seem to be S0s. ", + "conclusions": "The color evolution seen in Fig.\\ 3 is consistent with the idea that luminous early--type cluster galaxies have evolved with time, and that their stellar populations were younger at high redshifts. In itself this is no surprise; it is the uniformity and regularity of the spectrophotometric evolution which is the more important result. The degree of color evolution which we find is mild. At $z = 0.8$ to 0.9, the observed $R-K$ colors correspond approximately to rest--frame $U-J$. Early--type cluster galaxies at that epoch were only 0.4 to 0.6 magnitudes bluer at these wavelengths than are galaxies of the same morphological types in the Coma cluster today. Most of this change in color occurs because of an increase in the near--UV flux --- the IR--IR colors change very little. The mean ages of the stars in E+S0 galaxies in the distant clusters are thus younger than those in their present day counterparts, but their intrinsic colors are still quite red out to nearly $z=1$. Spectrophotometric models of passive evolution predict that UV--to--IR color evolution should be rapid during the first $\\sim$2 Gyrs after star formation ceases, quickly reaching red colors which evolve in a slow and steady fashion thereafter. Looking back to $z = 0.9$ (5.6 Gyr ago for our adopted cosmology), we evidently have not yet closely approached the time at which the bulk of the stars in early--type cluster galaxies were young. This conclusion is reinforced by the consistently small scatter in the galaxy colors out to $z = 0.9$. The rms dispersion in the $blue - K$ colors remains fixed at approximately 0.1 magnitudes over the redshift range of our sample, without evidence for the substantial change that might be expected if the early--type cluster galaxies formed over a broad range of redshifts. It is reasonable to expect some component of the intrinsic color scatter at all redshifts is due to metallicity variations, and that overall the evolution of early--type galaxies within a given cluster is highly synchronized. In addition, we see little evidence of substantial variations in the average galaxy colors from cluster to cluster at a fixed redshift, despite the fact that the clusters span a wide range of richnesses and x--ray luminosities. The constancy with redshift of the slope in the color--magnitude relation of the luminous E+S0 cluster galaxies strongly favors a scenario in which this slope arises from a mass--metallicity correlation, rather than one involving age. It appears that, on average, the evolutionary history of the majority of luminous early--type cluster galaxies has been very similar. It is important to recall that the redshift dependence of the galaxy colors shown in Figure 3 describes the evolution of the {\\it stellar populations} within those galaxies rather than the evolution of the {\\it galaxies} themselves. The histories of galaxies and of the stars from which they are made may, in principle, differ substantially. It is for this reason that hierarchical merging models, as described in the introduction, are able to match the color properties of nearby cluster galaxies, despite the fact that substantial merging takes place at relatively late times. In those models, the bulk of the {\\it stars} which comprise today's elliptical galaxies form early --- the assembly of the galaxy itself occurs fairly late. However, this merging is unlikely to occur without some impact on the galaxy colors due to episodes of star formation during the merging process. In particular, it has often been noted that mergers of disk galaxies cannot produce the high phase--space density of stars observed in the central regions of ellipticals unless some amount of star formation takes place in gas funneled to the nucleus of the merger remnant (e.g. Carlberg 1986; Hernquist, Spergel \\& Heyl 1993; Barnes \\& Hernquist 1996). If most elliptical galaxies in clusters form by a process of late mergers, there should be some detectable signature on their overall photometric properties such as the scatter in their colors, or the slope of the color--magnitude relation. By extending the redshift baseline for cluster photometric studies to $z = 0.9$, a redshift at which hierarchical models predict that galaxy formation is only partially complete, we have provided a data set against which detailed models can be tested. It seems that it is not difficult to reconcile the observed color evolution with traditional models of monolithic collapse and passive evolution for early--type galaxies, but it remains to be seen whether other scenarios can account for the observations. In this context, it will be important to examine the luminosity functions of cluster galaxies as a function of redshift and of morphological type, as these should be more sensitive to the merging history of the galaxies than the colors themselves. The upper redshift limit of our cluster sample just reaches the point at which serious constraints can be placed on the evolutionary history of early--type galaxies. Beyond $z\\sim1$ in cosmologically flat CDM models, the amount of merging occurring within the prior $\\sim$1 Gyr is sufficiently large so as to seriously inflate the locus of early--type galaxy colors, even if the recent, merger--induced starbursts are small (Kauffmann 1996). The identification of clusters at $z > 1$ and the characterization of their galaxy populations should provide a powerful means of testing galaxy formation theories." + }, + "9708/astro-ph9708201_arXiv.txt": { + "abstract": "The orbital dynamics of Cyg X-3 are a key to understanding this enigmatic X-ray binary. Recent observations by the RXTE ASM and the OSSE instrument on GRO enable us to extend the baseline of arrival time measurements and test earlier models of orbital period evolution. We derive new quadratic and cubic ephemerides from the soft X-ray data (including ASM). We find a significant shift between the predicted soft X-ray phase and the light curve phase measured by OSSE from $\\sim 44$ to 130 \\keV. Some of the apparent phase shift may be caused by a difference in light curve shape. ", + "introduction": " ", + "conclusions": "" + }, + "9708/astro-ph9708185_arXiv.txt": { + "abstract": "Whether or not metal-rich HB stars are the dominant UV source in giant elliptical galaxies (gEs) is an important question in current astronomical research. We follow up our previous evolutionary population synthesis study with quantitative tests to answer this question affirmatively under the following three conditions: (1) Reimers' empirical mass loss formula is proper, (2) the mass loss efficiency parameter ($\\eta$) in metal-rich stars is somewhat larger than the value estimated from the metal-poor star studies, and (3) the true value of the helium enrichment parameter ($\\Delta$$Y$$\\!$/$\\!$$\\Delta$$Z$) is positive. All three important empirical characteristics of the UV upturn (i.e., the fact that strong UV upturns are restricted to gEs, the positive UV upturn-metallicity correlation, and the narrow range of the \\Teff of the UV sources) are closely reproduced for reasonable ranges of input parameters. We discuss the major sources of uncertainties in the models, such as the production and role of hot horizontal-branch stars in gEs, and the importance of galactic nucleosynthesis. ", + "introduction": "The ultraviolet (UV) upturn phenomenon in the spectra of giant elliptical galaxies (gE's) has been known since early space observations with UV capability became available (\\cite{cw79}). It is defined as the increase in flux with decreasing wavelength in the range $\\approx$ 1,000 -- 2,500~\\AA, as shown in Figure 1. Several important discoveries have been made related to the UV upturn. Firstly, strong UV upturns are found only in the spectra of gEs\\footnote{Metal-poor Galactic globular clusters show high ratios of UV-to-$V$ flux (e.g., van Albada, de Boer, \\& Dickens 1981) mainly because of opacity effects (Dorman, O'Connell, \\& Rood 1995; Yi, Demarque, \\& Oemler 1997 - hereafter YDO). However, their UV spectra are either flat or declining with decreasing wavelength, producing low ratios of far-UV-to-near-UV flux. Thus, it is correct to say that only gEs show a UV upturn with a steep slope in the UV spectrum.}. Secondly, IUE observations suggest a positive correlation between the magnitude of the UV upturn and Mg$_2$ index (\\cite{f83}; \\cite{b88}). If the Mg$_2$ index traces metallicity (although there is reason for caution [Worthey, Faber, \\& Gonzalez 1992]), this implies that a more metal-rich galaxy shows a stronger UV upturn. Lastly, Hopkins Ultraviolet Telescope (HUT) observations suggest that the sources of the UV photons are hot stars with a narrow range of temperature, i.e., \\Teff $\\approx$ 20,000 -- 23,000 K (Brown, Ferguson, \\& Davidsen 1995). Since the dominant light sources (main sequence [MS], red giant branch [RGB], and horizontal branch [HB] stars) all tend to become cooler as metallicity increases, the unexpectedly high UV flux in such old, metal-rich systems has been a puzzle. \\placefigure{fig1} Understanding the cause of the UV upturn is important for the following reasons: (1) it provides insight into the hot stellar component in elliptical galaxies, (2) it tests the stellar evolution theory, (3) it constrains the age and metallicity of the majority of stars in gEs, if the UV upturn is sensitive to age and metallicity as some models suggest (e.g., \\cite{gr90}; Bressan, Chiosi, \\& Fagotto 1994; \\cite{dor95}; YDO). The age-dependence of the UV upturn is particularly noteworthy because such models predict that the UV upturn becomes significant only at large ages when optical spectral evolution is hardly detectable. Finally, (4) the UV upturn implies significant corrections to model-predicted optical colors of distant (high redshift) galaxies (\\cite{grv87}; \\cite{bcf94}). The origin of the UV upturn has been controversial since the first observations were made, and several interpretations have been proposed. Young MS stars were among the favorite candidates as the UV sources in many studies (e.g., Gunn, Stryker, \\& Tinsley 1981; \\cite{grv87}; \\cite{r88}; \\cite{mb93}). However, no evidence of recent star formation has been found in the UV-strong galaxies (\\cite{o92}; \\cite{ber93}). Using the HUT, Ferguson et al. (1991) also found that a lack of C\\,IV absorption and the shape of the continuum were inconsistent with flux from a MS population having a standard initial mass function. Moreover, such hot MS stars (\\Teff $\\approx$~20,000 K: spectral type B) are short-lived. If the UV upturn were caused by young MS stars, it would be a transient feature, suggesting that all these UV-strong galaxies had experienced a secondary starburst recently, nearly at the same time, which is very unlikely. Post asymptotic giant branch (PAGB) stars were the next to attract attention (\\cite{bc93}; \\cite{mb93}). However, PAGB stars are also thought to be so short-lived that the number needed to reproduce the UV upturn in the UV-strong gEs would exceed that allowed by the fuel consumption theorem (\\cite{ct91}). In addition, during most of their lifetimes, PAGB stars are much hotter than the suspected UV sources in gEs. Core helium-burning stars (HB and evolved HB stars) soon became an attractive candidate because they also can be hot and bright (\\cite{gr90}, and references therein). In addition, their mean temperature can match the estimated temperature of the dominant UV source in gEs easily and does not change rapidly with time, thus having advantages in explaining the narrow range of the \\Teff of the UV sources. Since the HB in Galactic globular clusters tends to become hotter as metallicity decreases, {\\it metal-poor} HB stars have been suggested as the cause of the UV upturn (\\cite{aar78}; \\cite{ay87}; \\cite{l94}; \\cite{pl97}). However, even the most metal-poor $and$ oldest Galactic globular clusters do not show UV upturns that are as strong as those in UV-strong gEs (Dorman et al. 1995; YDO). Moreover, gEs are metal-rich. Thus, if the metal-poor HB stars were the major UV sources in gEs, the positive UV upturn-metallicity relation would be puzzling, unless even the metal-rich gEs contain a substantial number of metal-poor stars $and$ the metal-poor stars in the UV-strong galaxies are significantly older than the oldest Galactic globular clusters (\\cite{pl97}). Instead, Demarque \\& Pinsonneault (1988) suggested that, under the conventional assumptions of mass loss\\footnote{Horch et al. (1992) proposed that if mass loss on the RGB increases with metallicity, a more metal-rich population would contain more hot (low-mass) HB stars. This assumption is not empirically proven yet but consistent with Reimers' empirical formula of mass loss for a fixed efficiency (see YDO).} and galactic helium enrichment, low-mass HB stars evolve into UV-bright objects instead of becoming AGB stars. They found that this phenomenon, the so-called ``slow blue phase'' (SBP, [Horch, Demarque, \\& Pinsonneault 1992])\\footnote {The significance of the SBP is in its positive metallicity dependence. In some sense, the SPB phenomenon states the metallicity dependence of the combined AGB-manqu\\'{e} (\\cite{gr90}) and post-early-AGB (\\cite{ct91}) evolutionary phases, a point which had not been addressed before. Yi, Demarque, \\& Kim (1997, hereafter YDK) presented a mathematical analysis of the SBP and clarified the general confusion between the SBP and other terms.}, occurs more easily when metallicity is higher if $Z \\gtrsim$ \\Zsun. Then, the classical metallicity dependence of HB morphology (i.e., HB becomes redder as metallicity increases) should be reversed in the metal-rich regime ($Z \\gtrsim$ \\Zsun). Several evolutionary population synthesis (EPS) studies qualitatively showed that the hypothesis that such metal-rich, UV bright, core helium-burning stars are likely to be the dominant UV source in gEs is consistent with empirical data (\\cite{gr90}; \\cite{bcf94}; \\cite{dor95}; \\cite{yado95}; \\cite{bfdd97}; YDO). We call this {\\it the metal-rich HB hypothesis}. In this quantitative study, we show, following YDO, that EPS models based on the metal-rich HB hypothesis reproduce quite well the empirical discoveries related to the UV upturn phenomenon. We explore the sensitivity of the UV upturn in the models to the input parameters. We compare single abundance models and a few composite models to observations of gE's. We then discuss major uncertainties in the EPS models and the origin of the discrepancies between various EPS studies. Finally, the implications of the UV upturn for understanding galaxy evolution are also discussed. ", + "conclusions": "The models based on the metal-rich HB hypothesis seem to satisfy all the empirical constraints related with the UV upturn phenomenon for reasonable input parameters. Under the conventional assumptions of stellar evolution theory, evolved low-mass, metal-rich ($Z \\gtrsim$~\\Zsun) core helium-burning stars are likely to be the dominant UV source in gEs, as suggested by earlier studies (\\cite{gr90}; \\cite{bcf94}; \\cite{dor95}; \\cite{yado95}; \\cite{bfdd97}; YDO). The UV upturn is an intricate phenomenon played by an orchestra of various instruments among which the following two phenomena have the most profound effects. (1) More metal-rich red giants experience higher mass loss even for a $fixed$ mass loss efficiency parameter $\\eta$, according to Reimers' formula. This is because a more metal-rich red giant has a higher opacity in the atmosphere, and the higher opacity causes a larger stellar radius, and a smaller surface gravity, which results in a larger mass loss (see discussion in Horch et al. (1992) and references therein). After the large mass loss, a more metal-rich red giant becomes a lower-mass HB star. However, a large opacity in metal-rich stars causes even low-mass HB stars to become red (low \\Teff). So mass loss cannot reproduce by itself the observed magnitude of the UV upturn. (2) The SBP (slow blue phase) phenomenon, the UV bright phase of the core helium-burning stars, is more prominent in more metal-rich stars when a positive \\DYDZ is assumed. Consequently, the magnitude of the UV upturn increases with increasing metallicity under the assumption of a positive \\DYDZ, which is consistent with the empirical discovery (\\cite{f83}; \\cite{b88}). Composite models seem to reproduce the observed range of the magnitude and the characteristic temperature of the UV upturn (\\Tuv) reasonably. However, their match is not as good as those of single abundance models. This may indicate a large uncertainty in galactic chemical evolution models. If such small mismatchs are considered acceptable, we may claim that a simple instantaneous burst model of the formation of gEs naturally develops a UV upturn with the observed characteristics. The full range of the observed magnitude of the UV upturn (or, \\fnuv) can be produced either by a metallicity dispersion (e.g., the majority of stars in UV-strong galaxies are $Z \\approx$ 2~\\Zsun and those in UV-weak galaxies are $Z \\approx$ \\Zsun) or by an age dispersion (UV-strong galaxies being older than UV-weak galaxies) among gEs, or perhaps both combined. At least some age difference seems to be favored if a factor of two difference in metallicity among gEs is unlikely. Models (single abundance models) also match the observed range of \\Tuv (the characteristic temperature of the UV upturn) rather precisely. It is interesting to note the apparently reverse correlation between the magnitude of the UV upturn and \\Tuv (with the exception of NGC\\,1399), although it is not yet statistically significant. Under the current assumption of input parameters, this phenomenon is also understandable by the same scenario that explains the UV upturn-metallicity relation. The UV upturn may serve as a relative-age indicator, provided the mean metallicity in a gE can be independently determined either through spectroscopic or photometric studies. However, it seems premature to use the UV upturn as an absolute-age indicator until input parameters are much better constrained. A more secure calibration of the UV upturn (e.g., for the precise dating of galaxies) will require improvements in our understanding of various things including mass loss during stellar evolution and the origin and evolutionary status of sdB stars in gEs. Willson et al. (1996) (also \\cite{bw91}) have recently claimed that their hydrodynamical models describe the mass loss better than Reimers' formula with a fixed mass loss efficiency parameter. It would be a great step forward if such new approaches can provide the astronomical community with a parameterized mass loss formula for single stars that matches the observations and is easy to use. In the same line, the role and frequency of binary stars, which appear responsible for at least a fraction of the sdB's in the stellar population of gE's, has to be studied further. More fundamentally, it is crucial to obtain a larger and reliable sample of spectra of nearby gEs. It is almost unbelievable that there is hardly any gE whose well-calibrated spectrum is available from far-UV to infrared. It is important to acquire the whole wavelength range of spectrum in order to find a unique solution using the EPS technique. The far-UV (including the $Lyman$ break) spectrum is particularly important because it provides important clues to the properties of the UV sources. We have not even touched other complexities, such as uncertainties in the stellar evolutionary calculation, which can be tested best by observing nearby stellar populations, and in the spectral library. The role of dust and of galaxy merging history could be as important as the ones that are discussed here. So, the question about the UV upturn is still open. However, we feel that the current EPS studies are going toward its solution; at least, this model (the metal-rich HB hypothesis) has survived so far." + }, + "9708/astro-ph9708246_arXiv.txt": { + "abstract": "It is assumed, in general, that the electromagnetic spectrum in the Primordial Universe was a blackbody spectrum in vacuum. We derive the electromagnetic spectrum, based on the {\\it Fluctuation-Dissipation Theorem} that describes the electromagnetic fluctuations in a plasma. Our description includes thermal and collisional effects in a plasma. The electromagnetic spectrum obtained differs from the blackbody spectrum in vacuum at low frequencies. In particular, concentrating on the primordial nucleosynthesis era, it has more energy for frequencies less than $3$ to $6\\omega_{pe}$, where $\\omega_{pe}$ is the electron plasma frequency. ", + "introduction": "\\label{sec:In} It is usually assumed in Cosmology that the primordial plasma was a homogeneous plasma and that the electromagnetic field was a blackbody spectrum in vacuum. These assumptions are used, for example, in standard Big Bang Nucleosynthesis calculations. Deviations from a blackbody spectrum in vacuum and a homogeneous plasma can affect Primordial Nucleosynthesis and in this letter we concentrate on this epoch. In the epoch of Big Bang Nucleosynthesis, the universe was a thermal bath of photons, electrons, positrons, baryons and neutrinos. The usual treatment considers the universe as a homogeneous plasma in thermal equilibrium and the electromagnetic field as a blackbody spectrum in vacuum. For example, in the energy density calculation, the energy density of the photons is given by the energy density of the blackbody spectrum in vacuum. In this manner, the Primordial Universe is treated as an ideal gas: {\\it collective effects} are assumed to be negligible. However, a plasma differs from an ideal gas due to correlations, for example, two-particle correlations. In an ideal gas, the distribution function for two-particles is: $F_{2}(x_{1},x_{2})=F_{1}(x_{1})F_{2}(x_{2})$. In a plasma we have an additional term $F_{2}(x_{1},x_{2})=[1+P_{12}(x_{1},x_{2})]F_{1}(x_{1})F_{2}(x_{2})$, where $P_{12}$ is the Debye-Huckel screening. ($P_{12} \\propto exp(-\\mid x_{2} - x_{1} \\mid/\\lambda_{D})/\\mid{x_{2}-x_{1}}\\mid$, where $\\lambda_{D}$ is the Debye length.) A plasma, even in thermal equilibrium, has fluctuations, that is, the physical variables such as temperature, density and electromagnetic fields, fluctuate. Even for a non-magnetized plasma, where the average magnetic field is zero, $\\langle B \\rangle = 0$, the squared average is not zero: $\\langle B^{2} \\rangle \\neq 0$. The study of the electromagnetic fluctuations in a plasma has been made in numerous studies, including those of Dawson, Rostoker et al., Sitenko and Gurin and Akhiezer et al. \\cite{daw}. Most of the results are compiled in Sitenko and Akhiezer et al. \\cite{sit}. The electromagnetic fluctuations are described by the {\\it Fluctuation- Dissipation Theorem} \\cite{sit}. The intensity of such fluctuations is highly dependent on how the plasma is described, for example, on the dissipation mechanisms present in it. It is necessary to describe the plasma in the most complete way. Cable and Tajima \\cite{ct} (see also \\cite{tc,t}) studied the magnetic field fluctuations, for several cases. Two of there descriptions concern the primordial plasma, which is an isotropic, non-magnetized and non-degenerate plasma. In particular, they studied: a) a cold, gaseous plasma and b) a warm, gaseous plasma described by kinetic theory. In their study, Cable and Tajima \\cite{ct} in case (a) used the {\\it cold plasma} description with a constant collision frequency. In case (b) they analyzed the spectrum only for low frequencies, with the {\\it warm plasma} description for phase velocity $\\omega/k$ less or equal to the thermal velocity of the electrons, $v_{e}$, and the ions, $v_{i}$, in a collisionless description. Through the study of the electromagnetic fluctuations, that is, the magnetic and electric field transverse fluctuations, given for example by the {\\it Fluctuation-Dissipation Theorem}, a reliable electromagnetic spectrum can be obtained. In this letter, we present a model that extends the work of Cable and Tajima \\cite{ct}. It includes in the same description collisional and thermal effects. By using the {\\it Fluctuation-Dissipation Theorem} relations, we derive the electromagnetic spectrum in a plasma. We concentrated, in particular, in the primordial nucleosynthesis era, in the case of the electron-positron plasma at high temperatures and the electron-proton plasma at low temperatures. The first case is the plasma at the beginning of the primordial nucleosynthesis, when the number of electrons and positrons was comparable to the number of photons. The second case is the plasma at low temperatures, after the annihilation of the electrons and positrons where the number of electrons and protons was $\\eta \\sim 10^{-10}$ smaller than the number of photons. In Section II we present the expression of the electromagnetic fluctuations, in Section III our model, and in Section IV a discussion of our results and conclusions. ", + "conclusions": "\\label{sec:con} The unique manner to obtain the electromagnetic transverse spectrum is analyzing the magnetic and electric field transverse fluctuations. This is the only manner to obtain information, not only about modes that propagate, like photons, but also modes that do not propagate. These modes appear, not only at low frequencies but also at high frequencies, resulting from the correlations in the plasma. Only at very high frequencies, the photons contribute uniquely to the magnetic and electric field transverse spectrae. We present a model that incorporates, in the same description, the thermal and collisional effects and used the {\\it Fluctuation-Dissipation Theorem} that describes the electromagnetic fluctuations. We use the Vlasov equation with the BGK collision term. The final electromagnetic spectrum for the primordial plasma at the epoch of Big Bang Nucleosynthesis behaves like a blackbody spectrum in vacuum for high frequencies. However, for low frequencies, it is distorted. It has more energy than the blackbody spectrum in vacuum. In the case of the high-temperature plasma ($T=10^{10}~K$), this range is for frequencies $\\omega \\leq 3\\omega_{pe}$ and for the low-temperature plasma ($T=10^{9}~K$) the range is for $\\omega \\leq 6\\omega_{pe}$, where $\\omega_{pe}$ is the electron plasma frequency. This additional energy is due to the collective modes of the plasma. The reason why the collective modes of the plasma can have more energy for $\\omega \\leq \\omega_{pe}$ than the photons in vacuum, can be understood as follows. Photons are massless bosons with the dispersion relation $\\omega^{2}=k^{2}c^{2}$. For the energy interval $0 \\leq \\omega \\leq \\omega_{pe}$, the wave number interval is $k=0$ to $k=\\omega_{pe}/c$. A relatively small amount of phase space is involved. For the collective motions of the plasma, in general, we have a larger amount of phase space. For example, for plasmons with energy $\\omega \\sim \\omega_{pe}$, the amount of phase space extends to a maximum $k$ of $k_{D} \\cong \\omega_{pe}/v_{T}$, where $v_{T}$ is the thermal electron velocity, which is greater than $\\omega_{pe}/c$ for the photons. In general, for a given frequency for $\\omega < \\omega_{pe}$, the greater phase space available to the collective modes of the plasma (than that of the photons) implies more energy, or a higher spectrum. This result (the additional energy that appears in the electromagnetic spectrum compared to the blackbody energy usually assumed) can affect several fields in Cosmology, in particular, Big Bang Nucleosynthesis. This extra energy, that has not been previously taken into account, causes the Universe to expand more rapidly at a given temperature. In particular, it causes the neutrinos to decouple earlier (at a higher temperature) and the neutron to proton ratio to freeze-in at a higher value. (In order to estimate the total additional energy, we have to add also, the longitudinal energy due to the longitudinal electric field spectrum. For $T=0.8~MeV$, we obtain for the additional energy, $\\Delta \\rho \\cong 1\\% \\rho_{\\gamma}$, where $\\rho_{\\gamma}=\\rho_{BB}$, the energy density of the blackbody photon spectrum in vacuum. A complete study, estimating the additional energy for diverse temperatures and densities, is in preparation \\cite{op}). Another interesting aspect is how the additional energy affects the spectrum of the microwave background. The additional energy at $z_{DEC}$ (the redshift when the spectrum was formed) occurred at frequencies $\\sim 10^{-9}~\\omega_{peak}$ ($\\omega_{peak}=2.8k_{B}T/\\hbar$). However, non-linear effects in plasma can bring the additional energy to higher frequencies. This is very interesting and should be investigated in the future. The authors would like to thank Swadesh Mahajan for useful suggestions, especially concerning the BGK collision term. The authors also would like to thank Arthur Elfimov for helpful discussions and the anonymous referees for helpful comments. M.O. would like to thank the Brazilian agency FAPESP for support and R.O. would like to thank the Brazilian agency CNPq for partial support." + }, + "9708/astro-ph9708070_arXiv.txt": { + "abstract": "We present results of N-body/gasdynamical simulations designed to investigate the evolution of X-ray clusters in a flat, low-density, $\\Lambda$-dominated cold dark matter (CDM) cosmogony. The simulations include self-gravity, pressure gradients and hydrodynamical shocks, but neglect radiative cooling. The density profile of the dark matter component can be fitted accurately by the simple formula originally proposed by Navarro, Frenk \\& White to describe the structure of clusters in a CDM universe with $\\Omega=1$. In projection, the shape of the dark matter radial density profile and the corresponding line-of-sight velocity dispersion profile are in very good agreement with the observed profiles for galaxies in the CNOC sample of clusters. This suggests that galaxies are not strongly segregated relative to the dark matter in X-ray luminous clusters. The gas in our simulated clusters is less centrally concentrated than the dark matter, and its radial density profile is well described by the familiar $\\beta$-model. As a result, the average baryon fraction within the virial radius ($r_{\\rm vir}$) is only $85$-$90 \\%$ of the universal value and is lower nearer the center. The total mass and velocity dispersion of our clusters can be accurately inferred (with $\\sim 15\\%$ uncertainty) from their X-ray emission-weighted temperature. We generalize Kaiser's scalefree scaling relations to arbitrary power spectra and low-density universes and show that simulated clusters generally follow these relations. The agreement between the simulations and the analytical results provides a convincing demonstration of the soundness of our gasdynamical numerical techniques. Although our simulated clusters resemble observed clusters in several respects, the slope of the luminosity-temperature relation implied by the scaling relations, and obeyed by the simulations, is in disagreement with observations. This suggests that non-gravitational effects such as preheating or cooling must have played an important role in determining the properties of the observed X-ray emission from galaxy clusters. ", + "introduction": "Galaxy clusters, the largest virialized systems in the universe, are useful cosmological probes. For example, the abundance of massive clusters (characterized either by mass or X-ray temperature) depends sensitively on $\\Omega_0$, the cosmological density parameter, and on $\\sigma_8$, the rms amplitude of density fluctuations on the fiducial scale $8 \\hmpc$ (White, Efstathiou \\& Frenk 1993, Eke, Cole \\& Frenk 1996, Viana \\& Liddle 1996). Thus, the present-day abundance of clusters and its redshift evolution may be used to place constraints these two fundamental cosmological parameters. Similarly, the observed baryon fraction in clusters places strong constraints on the value of $\\Omega_0$ (White \\etal 1993). Recent applications of these ideas tend to favor low values of $\\Omega_0\\simeq 0.3$ (White \\& Fabian 1995, Evrard 1997, Henry 1997) and, for flat models with this $\\Omega_0$, values of $\\sigma_8\\simeq 1$ which are broadly consistent with the amplitude of the microwave background anisotropies measured by COBE (Smoot \\etal 1992). To exploit fully the cosmological information encoded in the cluster population, it is necessary to understand their evolutionary history in some detail. This requires modeling the coupled evolution of the dark matter and gas, which together constitute the dominant contribution to the cluster mass. In its full generality, this problem is best approached by direct simulation and a variety of numerical techniques have now been developed for this purpose. Many of the techniques currently in use (including both Eulerian and Lagrangian hydrodynamics methods) have been recently compared by means of a test calculation of the formation of a cluster by hierarchical clustering in which the gas was assumed to be non-radiative (Frenk \\etal, in preparation). The different simulations resolved the cluster to different degrees, but in the regions resolved by each calculation, they generally gave remarkably similar results for most cluster properties of interest. Already the first N-body/gasdynamic simulations showed that in the non-radiative approximation, the X-ray properties of individual clusters formed in flat cold dark matter (CDM) cosmologies resemble those of real clusters in many respects (Evrard 1990b). Subsequent simulations have developed this theme further, generally with qualitatively similar conclusions (e.g. Thomas \\& Couchman 1992; Kang \\etal 1994; Cen \\& Ostriker 1994, Bryan \\etal 1994; Navarro, Frenk \\& White 1995, Owen \\& Villumsen 1997). Yet, it has been clear for some time, that the simulations (at least in an $\\Omega=1$ CDM cosmology) do not reproduce important systematic trends of the observed cluster {\\it population} such as the slope of the relation between X-ray temperature and luminosity. This has led several authors to argue that effects not included in the simulations, such as cooling or preheating of the gas, must have played a role in the evolution of the cluster population (Kaiser 1986, Evrard \\& Henry 1991, Navarro, Frenk \\& White 1995). In particular, Navarro, Frenk \\& White (1995) showed that moderate preheating at high redshift leads to an acceptable luminosity-temperature relation without spoiling the overall agreement with the observed structure of the X-ray gas in individual clusters. Useful insights into the evolution and systematic properties of the cluster population may also be obtained by studying scaling relations, an approach developed by Kaiser (1986; see also White \\& Rees 1978, and White 1982). These authors recognized that since gravity has no preferred scales, cluster properties determined primarily by gravity (or by other scale-free processes such as pressure gradients or hydrodynamical shocks) should obey simple scaling relations. Kaiser derived these for a population of clusters formed by hierarchical clustering from power-law initial density fluctuations in an Einstein-de Sitter universe. He concluded that, for most power spectra of interest, the cluster X-ray luminosity function evolves with redshift in the opposite sense to that indicated by the data available at the time (Edge \\etal 1990, Gioia \\etal 1990). More recent data, however, appear to be consistent with little or no evolution in the cluster X-ray luminosity function out to $z\\simeq 0.3$ (Nichol \\etal 1997, Rosati \\etal 1998). In this paper, we carry out a detailed investigation of the evolution of clusters in a low-density, $\\Omega_0=0.3$, CDM universe. We impose the flat geometry required by inflation by setting the cosmological constant $\\Lambda_0=0.7$ \\footnote{ Throughout this paper we write the cosmological constant $\\Lambda$ in units of $3 H^2$, so that a universe with $\\Omega+\\Lambda=1$ has a flat geometry. The present value of Hubble's constant, $H_0=H(z=0)$, is parameterized by $H_0=100 \\, h$ km s$^{-1}$ Mpc$^{-1}$.}. We perform a set of N-body-hydrodynamical simulations of cluster formation in this cosmology. We also generalize Kaiser's scaling laws to the case of an arbitrary cosmology and generic density fluctuation spectra. Cluster evolution in low-density universes has been explored numerically in a few previous papers (Cen \\& Ostriker 1994; Evrard, Metzler \\& Navarro 1996), but none has yet addressed in detail the evolutionary properties of the X-ray emission from individual clusters. Our extension of Kaiser's scaling laws in based on recent numerical results by Navarro, Frenk \\& White (1995, 1996, 1997, hereafter NFW95, NFW96, and NFW97, respectively), who found that virialized systems formed by hierarchical clustering exhibit a remarkable structural similarity. Throughout this paper we make the simplifying assumption that only gravity, pressure gradients and hydrodynamical shocks are important in the evolution of clusters. The plan of this paper is as follows. In \\S 2 we derive generalized scaling laws describing correlations between various cluster properties and their redshift evolution. In \\S 3, we describe our numerical methods and provide details of the ten clusters which we have resimulated at high resolution. These span a range of formation histories and dynamical states. Our main numerical results are presented in \\S 4 where we investigate the structure of the dark matter and gas in our clusters, the accuracy of cluster mass estimates, the evolution of their baryon fraction and the origin of possible deviations from a universal mean baryon fraction. In this section we also carry out a comparison with the generalized scaling laws derived in \\S 2. In \\S 5 we compare our results with previous numerical work and with observations. A summary of our main conclusions is given in \\S 6. ", + "conclusions": "\\label{sec:betaconc} We have used N-body/gasdynamical simulations to study the structure and evolution of X-ray clusters formed in a low-density CDM universe ($\\Omega_0=0.3$, $\\Lambda_0=0.7$, $h=0.7$, $\\sigma_8=1.05$). The simulations include gravity, pressure gradients and hydrodynamical shocks, but neglect the effects of radiative cooling or of galaxy formation. A summary of our main conclusions follows. \\noindent (1) The density profiles of clusters of different mass identified at various redshifts are described accurately by the fitting formula proposed by Navarro, Frenk \\& White (eq.~3; see also NFW95 and NFW96). The parameters of the fit are in good agreement with the analytical model proposed by these authors (NFW97). This formula provides an adequate description of the mass profile over approximately two decades in radius, out to the `virial' radius beyond which infall dominates. The extent of this virialized region is consistent with a definition of `virial radius' based on the spherical top-hat collapse model. \\noindent (2) The structure of cluster dark matter halos is in excellent agreement with the distribution and dynamics of galaxies in the clusters analyzed by the CNOC project. This is consistent with the idea that galaxies and dark matter in clusters are not spatially segregated or dynamically biased to a significant degree. \\noindent (3) The gas density profiles of simulated clusters differ significantly from the dark matter profiles and are better described using the $\\beta$-model (eq.17). However, the gas and dark matter density profiles remain proportional to each other regardless of cluster mass and redshift. Numerical estimates of the X-ray luminosity converge quickly when the scale radius of the dark matter and the core radius of the gas are resolved numerically. For an SPH simulation, this typically requires $\\gsim 3 \\times 10^3$ particles per cluster and an effective spatial resolution better than about one percent of the virial radius. \\noindent (4) The structural similarity between the dark and gas components implies that simple scaling laws relate the mass, velocity dispersion, temperature, and X-ray luminosity of galaxy clusters. These scaling laws can be derived using the fact that clusters of a given mass are described by a single free parameter: their characteristic density or concentration. These laws extend the scale-free relations of Kaiser (1986) to universes with $\\Omega \\neq 1$ and perturbation spectra different from power-laws. The predictions of these scaling laws, as a function of cluster mass and redshift, are in remarkable agreement with the results of the simulations. This provides an impressive validation of the Smooth Particle Hydrodynamics technique. \\noindent (5) The X-ray luminosity in simulated clusters scales approximately as the square of the temperature, roughly as predicted by the scaling laws. This is a shallower dependence of $L_X$ on $T$ than is observed for X-ray clusters. We interpret this disagreement as requiring additional physical processes not included in these simulations (eg. radiative cooling or preheating) to account for the X-ray properties of clusters, particularly of low-temperature ($kT < 5$ keV) systems. This relation is expected to evolve only weakly with redshift. At a given temperature clusters are not expected to brighten by more than $30 \\%$ at $z \\sim 0.3$, consistent with published measurements. \\noindent (6) The average baryon fraction within the virial radius is $85$-$90 \\%$ of the universal mean, $\\Omega_b/\\Omega_0$, and is lower in the inner regions. This result calls for caution when interpreting the baryon fraction measured in clusters in terms of the universal mean. The inclusion of physical processes neglected here, such as radiative cooling, may affect the cluster baryon fraction although such effects are likely to be small. \\noindent (7) X-ray emission-weighted temperatures can be used to estimate reliably the total mass and velocity dispersion of clusters (eqs.~18 and 19). These estimators are essentially unbiased and have small scatter, $\\sim 20 \\%$ for the mass-temperature, and $\\sim 15 \\%$ for the mass-velocity relations. A consequence of this is that semianalytical techniques and N-body simulations can be used to predict the statistical properties of X-ray clusters in different cosmological models without the need for expensive hydrodynamical simulations. Physical processes not included in our simulations, such as radiative cooling, galaxy formation, or non-gravitational heating may all have a significant effect on the temperature of the intracluster medium and on the X-ray luminosity of galaxy clusters. We have chosen to neglect them in this study, and to concentrate instead on the simpler `adiabatic' evolution of gas within an evolving population of dark matter halos. The failure of `adiabatic' clusters to reproduce the observed luminosity-temperature relation indicates that additional physics must be included in the numerical modeling in order to develop a full understanding of the origin and evolution of the X-ray properties of galaxy clusters. We are currently working on these issues." + }, + "9708/astro-ph9708136_arXiv.txt": { + "abstract": "We study the geometry and topology of the large--scale structure traced by galaxy clusters in numerical simulations of a box of side 320 $h^{-1}$ Mpc, and compare them with available data on real clusters. The simulations we use are generated by the Zel'dovich approximation, using the same methods as we have used in the first three papers in this series. We consider the following models to see if there are measurable differences in the topology and geometry of the superclustering they produce: (i) the standard CDM model (SCDM); (ii) a CDM model with $\\Omega_0=0.2$ (OCDM); (iii) a CDM model with a `tilted' power spectrum having $n=0.7$ (TCDM); (iv) a CDM model with a very low Hubble constant, $h=0.3$ (LOWH); (v) a model with mixed CDM and HDM (CHDM); (vi) a flat low--density CDM model with $\\Omega_0=0.2$ and a non-zero cosmological $\\Lambda$ term ($\\Lambda$CDM). We analyse these models using a variety of statistical tests based on the analysis of: (i) the Euler--Poincar\\'{e} characteristic; (ii) percolation properties; (iii) the Minimal Spanning Tree construction. Taking all these tests together we find that the best fitting model is $\\Lambda$CDM and, indeed, the others do not appear to be consistent with the data. Our results demonstrate that despite their biased and extremely sparse sampling of the cosmological density field, it is possible to use clusters to probe subtle statistical diagnostics of models which go far beyond the low-order correlation functions usually applied to study superclustering. ", + "introduction": "The study of the distribution of matter on the largest scales amenable to observation can provide important constraints on models of the formation of cosmological structures. In particular, it has now become well established that a very accurate and efficient way of describing very large scale structure in the galaxy distribution is obtained by not looking at galaxies themselves but at rich clusters of galaxies. If the `standard' model of structure formation -- the gravitational instability picture -- is correct, the expected displacements of galaxy clusters from their primordial positions are much smaller than the typical separation of these objects. In principle, therefore, clusters of galaxies can yield clues about the primordial spectrum of perturbations that gave rise to them, without such clues being trampled on by the effects of non--linear evolution. Moreover, because clusters represent highly overdense regions in the cosmological density field, these objects display an enhanced clustering signal relative to that of galaxies on the same scale, an effect usually known as biasing (Kaiser 1984). This is the reason why so much effort has been devoted to compiling deep cluster surveys, starting with the pioneering work of Abell (1958), Zwicky et al. (1968) and Abell, Corwin \\& Olowin (1989), and leading up to extended redshift surveys both in the optical (e.g. Postman, Huchra \\& Geller 1992; Dalton et al. 1994; Collins et al. 1994, and references therein) and in the X--ray (e.g. Nichol, Briel \\& Henry 1994; Romer et al. 1994; Ebeling et al. 1996) regions of the spectrum. The properties of galaxy clusters may help to resolve some of the issues that have led to the present relative stagnation in the theory of structure formation. Since the demise of the standard model of the 1980s -- the standard Cold Dark Matter model (SCDM) -- a number of contending theories have been proposed which are in better agreement with the observations than SCDM but between which it is difficult to discriminate using present observations of galaxy clustering and the cosmic microwave background; for a review, see Coles (1996). It is therefore important to try to find statistical diagnostics of clustering that may reveal differences between these models and the data to see if they do indeed explain the details of the observed clustering phenomenon, as well as between the models themselves so one can understand how the various extra ingredients involved in these models alter specific characteristics of the clustering pattern. Simple two--point statistical descriptions of superclustering (i.e. the clustering of galaxy clusters) have already yielded important clues about the shape of the matter power spectrum on large scales (e.g. Peacock \\& Dodds 1994; Borgani et al. 1997) and, more recently, this has been extended to simple properties of the higher--order moments (e.g. Plionis \\& Valdarnini 1995; Plionis et al. 1995; Borgani et al. 1995; Gaztanaga, Croft \\& Dalton 1995). However, the complete statistical characterisation of the clustering requires knowledge of all the higher order moments or, equivalently, knowledge of the complete set of $n$--point correlation functions (Peebles 1980). Such a description is extremely laborious to construct, tends to be swamped by discreteness effects and sampling errors even at quite small $n$ and is in any case rather difficult to interpret geometrically. For these reasons it is useful to seek a description of clustering which by-passes this more orthodox approach and looks for intrinsically geometrical or topological signatures. One can hope that such approaches might lead to robust quantitative descriptions of the void-filament network which is visually apparent in the distribution of galaxies, and to relate this visual appearance to the interaction of non--linear gravitational dynamics on an initial density field with some assumed power spectrum. The hope is therefore to pick out differences between models which are hard to discern in measures such as the power spectrum. Various approaches to this question have been suggested and some of them have been more successful than others in their application to the data. One particular problem such descriptors face when they are applied to superclustering, for example, is that these objects are extremely rare and there are strong shot-noise effects which have to be compensated for in some way. In this paper, we aim to investigate a particular set of topological or geometrical descriptors of the pattern present in simulated cluster distributions and, where possible, to compare the results from simulations with the analogous results from the Abell/ACO cluster catalogue. We should stress at the outset that this is an exploratory work and there are reasons to suspect that the task of discriminating between these models and the data might be extremely difficult. First there is the problem of shot-noise we alluded to above. Secondly, the available cluster sample is quite small and may suffer from unknown selection effects. One can hope, however, that better controlled cluster samples may emerge fairly soon from ongoing galaxy redshift surveys. Third, it is extremely difficult to construct sufficiently large $N$-body simulations of galaxy clustering and select the appropriate clusters within them in the same way that clusters are selected observationally (e.g. Bahcall \\& Cen 1992; Croft \\& Efstathiou 1994; Eke et al. 1996). And finally, there is the ubiquitous problem of understanding how the objects one sees relate to the distribution of matter one calculates, a difficulty generically known by the name of biasing and which was first discussed in the context of rich clusters by Kaiser (1984). In the spirit of exploration, therefore, we shall use simplified models of superclustering, generated by using a method based on the Zel'dovich approximation. This method has been used in a number of previous studies of the distribution of clusters in both position and velocity space (Borgani, Coles \\& Moscardini 1994; Plionis et al. 1995; Borgani et al. 1995; Tini Brunozzi et al. 1995; Moscardini et al. 1996; Borgani et al. 1997) and is known to be accurate in comparison with the full $N$--body approach, provided the degree of non--linear evolution at the scale of individual clusters is not too strong. The outline of the paper is as follows. In Sections 2 and 3 we briefly describe our simulation method and the observed Abell/ACO cluster sample, respectively. We then go on to discuss the various clustering descriptors we use to analyse these data sets. First, in Section 4, we discuss the topological properties of the isodensity regions in the distribution traced by clusters, using a method described in detail by Coles, Davies \\& Pearson (1996) and which is similar (but not identical) to the well--known {\\em genus} statistic (reviewed by Melott 1990) and which has recently been applied to cluster data by Rhoads, Gott \\& Postman (1994). We next, in Section 5, discuss an analysis based on percolation theory. The last of our three approaches, presented in Section 6, is based on properties of a graph-theoretical construction known as the minimal spanning tree, in conjunction with a set of mathematical quantities intended to describe the shapes of pieces of the trees obtained (Pearson \\& Coles 1995). Each of the three analyses we attempt is expected to perform better in some situations than others, so in Section 7 we present an analysis of the statistical {\\em power} of these tests at discriminating between different models and between the models and the observed data. We also discuss the virtues of combining the various tests and show the statistical significance of the results we obtain by combining the different analyses into a composite test. We present our conclusions in Section 8. ", + "conclusions": "In the Introduction to this paper, we stressed that this analysis was to be treated as exploratory because there were reasonable grounds to doubt the quality of present clustering data and that looking for geometrical signatures of the pattern of superclustering was in any case difficult because of the extreme rareness of rich clusters and the consequent sparse sampling and shot-noise this implies. Nevertheless, as a guide to the results one might expect from larger and better controlled cluster samples the results we have obtained are extremely encouraging, at least for some of the tests we have used. Although this optimism is largely based on results from simulations which may be reasonably argued to be much `cleaner' than real data are likely to be, our results show at least that there are perceptible differences between these models on large--scales and that these do in principle allow one to discriminate between them using shape- and topology-based descriptors. For our topological analysis, based on the EPC, clear differences emerge between the models. One has to be a little careful here, however, because the form of the statistic we use actually contains information about the one--point distribution function of the objects, because of the choice of threshold parameter $\\nu$. Remember also that the amplitude of the EPC curve is related to the coherence length of the density field and that this is simply derived from the power spectrum. Comparing the trends we see in the EPC analysis with the trends of the one--point distribution found in an analysis of the same models by Borgani et al. (1995) together with the coherence lengths of the initial power spectra, shows that the behaviour of the EPC for different simulations can, roughly speaking, be `explained' in terms of these other descriptions. Although differences therefore show up between the models, they are largely the same as the differences one finds in non--topological descriptors. One would be justified therefore in saying that this descriptor does not add very much: it just provides a different way of seeing differences in one and two--point information. Nevertheless, folding such information in with the topology (which is in any case very easy to measure) does seem to provide a simple methodology for discriminating between models which does not require the computation of power-spectra and distribution functions and may in any case incorporate at least some extra information than these quantities do. On the other hand, the topology of the Abell/ACO data does not display the same kind of EPC graph that one would expect by looking at the results of Plionis et al. (1995) and Borgani et al. (1995) and assuming it follows the same trends as our models. This may be telling us that the Abell/ACO is essentially different to all of the models we have looked at in this paper, which in turn may mean that either all the models are incorrect or that there is something suspicious about the catalogues or the way we have interpreted them. In particular, the effects of redshift selection, galactic extinction and the differences in number density between the Abell and ACO catalogues introduce some uncertainty into our conclusions. The one model that does have a topological description in reasonable accord with the Abell/ACO data is the $\\Lambda$CDM model, a result which agrees with the results of Kerscher et al. (1997) (although the model they used had a rather smaller value of $\\Omega_\\Lambda=0.65$ than the model we have used here). This model also survives the tests described in Borgani et al. (1995), but there was uncertainty attached to that analysis because of the possibility of that model being too strongly clustered to be adequately described by the Zel'dovich approximations. We have shown that this extra evolution does not influence the behaviour of the EPC to any significant extent and the claim that this model can reproduce the behaviour of Abell/ACO in terms of topology and low-order moments therefore stands up to scrutiny. This, of course, still admits the possibility that this is telling us more about problems with the catalogue than about the real distribution of overdensities. The performance of our percolation test depends strongly on the kind of statistic one extracts from the percolated set. If one looks only at the statistic $\\mu_\\infty$ then the power of discrimination is mediocre, but this rises strongly if one uses $\\mu^2$ instead or together with $\\mu_\\infty$. The one disappointment of this analysis is the performance of the MST/shape functions we introduced in Pearson \\& Coles (1995). Although they do perform well for relatively well--sampled distributions, we were unable to get useful results for any of the simulated samples of clusters. The application of this statistic, at least in the form we have used it here, is not recommended for extremely sparsely-sampled distributions like those of Abell clusters. Our final conclusion, however, is that topological and geometrical descriptors (of which we have studied only three) are at least in principle capable of diagnosing differences between very sparsely-sampled distributions in a fashion which is quite independent of the one- and two-point statistics which are more familiar in the cosmological community. With the arrival of larger and better controlled samples of galaxy redshifts and the cluster catalogues which will accompany them, clustering data will not only be more amenable to this type of analysis, they will also {\\em require} such an approach if one is to extract as much information as possible." + }, + "9708/astro-ph9708252_arXiv.txt": { + "abstract": "We used {\\it HST} to obtain surface brightness fluctuation (SBF) observations of four nearby brightest cluster galaxies (BCG) to calibrate the BCG Hubble diagram of \\markcite{lp92} Lauer \\& Postman (1992). This BCG Hubble diagram contains 114 galaxies covering the full celestial sphere and is volume limited to 15,000 km s$^{-1},$ providing excellent sampling of the far-field Hubble flow. The SBF zero point is based on the Cepheid calibration of the ground $I_{KC}$ method \\markcite{t97}(Tonry et al. 1997) as extended to the WFPC2 F814W filter by \\markcite{a97} Ajhar et al. (1997). The BCG globular cluster luminosity functions give distances essentially identical to the SBF results. Using the velocities and SBF distances of the four BCG alone gives $H_0=82\\pm8{\\rm ~km~s^{-1}~Mpc^{-1}}$ in the CMB frame, valid on $\\sim$4,500 km s$^{-1}$ scales. Use of BCG as photometric redshift estimators allows the BCG Hubble diagram to be calibrated independently of recession velocities, yielding a far-field $H_0=89\\pm10{\\rm ~km~s^{-1}~Mpc^{-1}}$ with an effective depth of $\\sim$11,000 km s$^{-1}$. The error in this case is dominated by the photometric cosmic scatter in using BCG as distance estimators. The concordance of the present results with other recent $H_0$ determinations, and a review of theoretical treatments on perturbations in the near-field Hubble flow, argue that going to the far-field removes an important source of uncertainty, but that there is not a large systematic error to be corrected for to begin with. Further improvements in $H_0$ depend more on understanding nearby calibrators than on improved sampling of the distant flow. ", + "introduction": "A key part of measuring the Hubble constant, $H_0,$ is to look out far enough so that the bulk velocities of galaxies are trivial compared to the Hubble flow itself. Due to the Virgo cluster infall pattern, observation of the unbiased Hubble flow can only be contemplated at distances in excess of $\\sim3000$ km s$^{-1}$. Furthermore, bulk flows on even larger scales, such as those associated with the Great Attractor, may bias measurement of $H_0.$ \\markcite{tco92}Turner, Cen, \\& Ostriker (1992) and \\markcite{shi}Shi, Widrow, \\& Dursi (1996), for example, show that under standard theories of structure formation, measurements of $H_0$ can depart significantly from its true ``global'' value due to the inhomogeneous distribution of matter in the universe, unless care is taken to sample deeply with large angular coverage. Indeed, a common concern with many recent $H_0$ determinations is that they are not truly sampling the distant Hubble flow \\markcite{bar}(Bartlett et al. 1995). Characterizing the far-field requires observing large numbers of objects at large distances so that the Hubble diagrams are insensitive to random peculiar velocities or bulk flows. Hubble diagrams at present are largely based on the Tully-Fisher or $D_n-\\sigma$ relationships, the luminosities of supernovae (SN Ia or SN II), and brightest cluster galaxies (BCG). Tully-Fisher distances are available out to $\\sim$9,000 km s$^{-1},$ and have been recently used to measure a far-field $H_0$ \\markcite{gio97}(Giovanelli et al, 1997), while $D_n-\\sigma$ have full-sky coverage out to only $\\sim$6,000 km s$^{-1}.$ Only a few SN II have been observed in sufficient detail at large distances \\markcite{schm}(Schmidt et al. 1994), but the SNIa Hubble diagram is becoming richer with time and provides some sampling of the Hubble flow out to $\\sim$30,000 km s$^{-1}$ \\markcite{rpk}(Riess, Press, \\& Kirshner 1996; \\markcite{ham} Hamuy et al. 1996). At present, however, calibration of the SNIa distance scale remains controversial (see \\markcite{sand96}Sandage et al. 1996), and the SN Ia diagrams remain relatively sparse at large distances. In this work we focus on calibrating the BCG Hubble diagram, which is based on a recent characterization of BCG as relative distance estimators \\markcite{pl}(Postman \\& Lauer 1995). In the classic work of \\markcite{sand72} Sandage (1972) and \\markcite{sand73} Sandage \\& Hardy (1973), BCG were used to show that the Hubble flow was linear over a large range in redshift. \\markcite{lp} Lauer \\& Postman (1994) observed BCG to define a frame for measuring the peculiar velocity of the Local Group, but as this work was in progress they realized that they could test for $H_0$ variations with distance with greater precision than was previously available in response to the concerns of \\markcite{tco92}Turner, Cen, \\& Ostriker (1992). \\markcite{lp92} Lauer \\& Postman (1992) presented a Hubble diagram based on the 114 BCG that defined the volume-limited full-sky sample of Abell clusters within 15,000 km s$^{-1},$ which is shown again here in Figure \\ref{bcg_hub}. In brief, the absolute magnitudes of BCG, $L_m,$ measured in apertures of fixed metric size, $r_m,$ can be predicted from $\\alpha\\equiv d\\log L_m/d\\log r|_{r_m}$ \\markcite{hoe}(Hoessel 1980). Figure \\ref{bcg_hub} shows the metric luminosities as apparent fluxes, corrected by the $L_m-\\alpha$ relationship to a standard value of $\\alpha=0.5.$ \\begin{figure}[tb] \\plotone{figure1.ps} \\caption{The BCG Hubble diagram. R-band metric luminosities of the BCG, corrected by the $L_m-\\alpha$ relationship, are plotted as a function of velocity in the Local Group frame. The line is the mean Hubble relation fitted.} \\label{bcg_hub} \\end{figure} The BCG Hubble diagram slope is $0.996\\pm0.030$ of the expected value, consistent with a uniform Hubble flow over $0.01\\leq z\\leq0.05.$ \\markcite{lp92} Lauer \\& Postman (1992) limit any variation in the {\\it apparent} or local $H_0$ (the Hubble constant measured over a limited depth) as compared to $H_0$ measured globally over the entire volume, to $\\delta_H\\equiv\\Delta H_0/H_0<0.07.$ The SNIa Hubble diagram also shows no evidence for $H_0$ variations with distance; \\markcite{rpk}Riess, Press, \\& Kirshner (1996) show its slope (relative to Euclidean) to be $1.005\\pm0.018.$ The full-sky coverage of the Abell cluster sample is crucial, as any dipole pattern caused by large bulk flows (such at that advanced by \\markcite{lp}Lauer \\& Postman 1994) will integrate out of the Hubble diagram to first order. The linearity of the BCG Hubble diagram shows that an excellent estimate of the far-field $H_0$ can be obtained once the zero point of the diagram is calibrated. We note that BCG presently provide the only volume-limited sample that explores the Hubble flow at these distances. A Hubble constant can be obtained from the BCG Hubble diagram once an absolute distance is known to a subset of the galaxies. In essence, one transfers the full sample to a common distance, and finds the average absolute luminosity of the BCG on the assumption that the calibrating set is typical. Random velocities and bulk flows of the BCG contribute to the ``cosmic scatter'' in their luminosity distribution, but cause no systematic offset (with the caveats discussed in $\\S$\\ref{far_enough}). We contrast this approach to others that use the apparent distance ratio between the Virgo and Coma clusters, or any other near and far aggregate of galaxies, to reach the far-field. Instead, we are using the BCG as complete probes of the Hubble flow over a large volume. We chose to calibrate the BCG Hubble diagram with surface brightness fluctuation (SBF) distance estimates to four of the nearest BCG. The SBF method \\markcite{ts88} (Tonry \\& Schneider 1988) uses the ratio of the second to first moments of the stellar luminosity function within early-type stellar systems as a distance estimator. The ratio of moments corresponds to an apparent magnitude, \\mbar, that in the near-IR corresponds to the brightness of a typical red giant star. When the images are deep enough such that a star of apparent luminosity \\mbar~contributes more than a single photon to an observation, the random spatial point-to-point surface brightness fluctuations in a galaxy image are dominated by the finite number of stars it comprises, rather than photon shot noise. A power spectrum of the SBF pattern provides \\mbar. Use of the SBF method on galaxies with distances known from other methods \\markcite{jetal92}(see Jacoby et al. 1992 for additional details) provides the zero point \\Mbar, allowing absolute distances to be computed from \\mbar. The most recent calibration of the SBF method is presented by \\markcite{t97} Tonry et al. (1997). Major components of this work are: 1) understanding how \\Mbar\\ varies with stellar population, 2) determining the zero point of the method, and 3) establishing the universality of the calibration. Tonry et al. observe in the $I_{KC}$ band, which minimizes variations in \\Mibar with stellar population {\\it ab initio.} They also show that variations in \\Mibar~are fully characterized by the ($V-I$) colors of the stellar systems. Based on 149 nearby galaxies they find \\begin{equation} \\label{eqMIbar} \\Mibar = (-1.74 \\pm 0.07) + (4.5 \\pm 0.25) [\\viz - 1.15]. \\end{equation} This relationship has scatter of only 0.05 mag and agrees well with the theoretical calculations of \\markcite{wora} Worth\\-ey (1993a, \\markcite{worb} 1993b) both in slope {\\it and} zero point. \\markcite{ta92} Tammann (1992) was concerned that an earlier SBF calibration based on $(V-I)$ was incomplete and that \\Mibar~additionally depended on the galaxies' ${\\rm Mg_2}$ indices. In response, Tonry et al. use their extensive sample to show that there is no correlation between the residuals of equation (\\ref{eqMIbar}) and ${\\rm Mg_2.}$ The zero point of equation (\\ref{eqMIbar}) is based on Cepheid distances to seven spiral galaxies with bulge SBF observations. Tonry et al. present numerous comparisons of SBF to PNLF, Tully-Fisher, $D_N-\\sigma$, SNIa, and SN II distances, finding no evidence for any systematic offset between SBF bulge and elliptical galaxy measurements, nor any other systematic effect that challenges the calibration. Although the nearest BCG are too far away for the SBF method to work from the ground, the high spatial resolution of {\\it HST} allows SBF to be used beyond the 15,000 km s$^{-1}$ depth of the \\markcite{lp} Lauer \\& Postman (1994) sample. An important caveat is that there is no direct match to the $I_{KC}$ filter among the WFPC2 filter set. The F814W filter is a close analogue to $I_{KC}$ (see \\markcite{holtza}Holtzman et al. 1995a), but requires additional calibration to tie it to the \\markcite{t97}Tonry et al. (1997) zero point. \\markcite{a97} Ajhar et al. (1997) accomplished this task in preparation for the present work, by comparing {\\it HST} F814W SBF observations to the $I_{KC}$ results for 16 galaxies in the Tonry et al. sample. For the WFPC2 CCDs and F814W filter, Ajhar et al. find \\begin{equation} \\label{eqM8bar} \\Mf8bar = (-1.73 \\pm 0.07) + (6.5 \\pm 0.7) [\\viz - 1.15], \\end{equation} with scatter similar to that about equation (\\ref{eqMIbar}). A key difference between equation (\\ref{eqMIbar}) and (\\ref{eqM8bar}) is the steeper relationship between \\Mf8bar and $(V-I),$ which Ajhar et al. show is consistent with the differences between the F814W and $I_{KC}$ filters. Calibration of {\\it HST} for SBF work is thus crucial for the present work. ", + "conclusions": "We have used {\\it HST} to obtain SBF distances to four BCG beyond 4,000 km s$^{-1}$ to calibrate the \\markcite{lp92}Lauer \\& Postman (1992) BCG Hubble diagram, producing an estimate of the global value of $H_0$ valid on $\\sim$11,000 km s$^{-1}$ scales. This method gives $H_0=89\\pm10{\\rm ~km~s^{-1}~Mpc^{-1}},$ and is based on the full \\markcite{lp}Lauer \\& Postman (1994) 15,000 km s$^{-1}$ volume limited BCG sample. As such, the result is independent of Virgo or Coma cluster distances and membership issues, as well as the recession velocities of the four BCG studied. The large error reflects the photometric scatter about the $L_m-\\alpha$ ridge line, which was used to transfer the BCG Hubble diagram to the SBF distance scale. As more BCG are observed with {\\it HST,} the formal errors in this far-field $H_0$ should decrease. Our review of the present understanding of the formation of large scale structure argues that we are likely to have fairly sampled the far-field. Even theories with enough power on large spatial scales to generate bulk flows as large as those observed by \\markcite{lp}Lauer \\& Postman (1994) are unlikely to have deviations outside of $|\\delta_H|\\lesssim0.05$ for the volume sampled by the BCG Hubble diagram. In contrast, the compatibility of our results with those based on more nearby objects argues that there is little effect on $H_0$ and the depth of the measurements. Going to the far-field most likely removes a source of uncertainty, rather than correcting for a systematic error. Indeed we find $H_0=82\\pm8{\\rm ~km~s^{-1}~Mpc^{-1}}$ just from Hubble ratios based on the SBF distances and observed recession velocities to the four SBF-calibrated BCG at $\\sim$4,000 km s$^{-1}$ alone, a result consistent with our far-field result. The present $H_0$ rests on calibration of the SBF method and an understanding of its systematic effects. At the fundamental level, we are tied to the nearby Cepheid calibrators. Changes in the Cepheid scale will propagate to the present results through the Tonry et al. and Ajhar et al. calibrations. As noted in the introduction, \\markcite{t97}Tonry et al. (1997) SBF calibration is tied to seven spiral galaxies with Cepheid distances. Further, Tonry et al. have observed enough galaxies to perform an exhaustive series of tests, finding no systematic offsets between SBF observations of bulges and elliptical galaxies. A weaker link is transferring the ground $I_{KC}$ method to the WFPC2 F814W filter, a task accomplished by \\markcite{a97}Ajhar et al. (1997); we will attempt to refine this calibration as more nearby systems are observed with {\\it HST.} We conclude that the major uncertainties in the distance scale are those close to home rather than far away." + }, + "9708/astro-ph9708228_arXiv.txt": { + "abstract": "We review the main observational characteristics of AM Herculis stars (polars) at X-ray, EUV, UV, IR and optical wavelengths. Particular emphasis is given to multi-epoch, multi-wavelength observations of the eclipsing polar HU Aqr (RX\\,J2107.9-0518). In AM Herculis stars the broad-band spectral energy distribution from X-rays to the IR is governed by only very small structures: the hot accretion regions on the footpoints af accreting field lines. The most extended structures in the binary systems on the other hand, the mass-donating secondary stars and the accretion streams, distinctly appear only as Doppler-shifted emission or absorption lines. They can best be studied by investigating selected narrow spectral features in the optical, ultraviolet or the near infrared. In this contribution both aspects will be highlighted, the structure of the accretion regions as inferred from multi-wavelength observations with low or no spectral resolution, as well as the structure of the secondary stars and the accretion streams as inferred from high-resolution spectral observations and Doppler mapping. ", + "introduction": "The broad-band spectral energy distribution of polars is governed by the processes in the small accretion regions on the footpoints of field lines, which channels the originally free-falling accretion stream down to the white dwarf. The release of gravitational energy is manifested primarily as bremsstrahlung at hard X-rays, quasi-blackbody radiation, which is prominent at EUV/soft X-ray wavelengths, and as cyclotron radiation, which is detected from the IR to the near-UV regime. Details about the relevant processes acting in the accretion region and the influence of the main parameters: the specific mass flow rate $\\dot{m}$ (in g cm$^{-2}$ s$^{-1}$), the magnetic field strength $B$ and the mass of the white dwarf $M_{\\rm wd}$, plus compilations of the observational data related to those (X-ray spectra, low-resolution optical spectra) have been given in recent reviews by e.g.~Beuermann (1997), Beuermann \\& Burwitz (1995) and Schwope (1996). In the present paper we concentrate on the shape of light curves mainly in the X-ray region. A different perspective of these systems is possible when viewed through ultraviolet ``glasses'', which reveal both the heated and unheated parts of the photosphere of the white dwarf plus reprocessed stream emission. Although observations in the ultraviolet have been performed for decades now using IUE, the data quality has improved with the advent of HST observation of polars. We describe here some preliminary results of low spectral resolution UV observations, with full phase-coverage of the eclipsing polar HU Aqr. The existence of accretion streams in polars is a well-established observational fact. It derives e.g.~from the broad emission lines in the optical/IR/UV with high radial velocities (up to $\\sim$2000 km s$^{-1}$) varying quasi-sinusoidally, or the absorption dips seen preferentially in the X-ray light curves or, more indirectly, from the existence of hot plasma in small regions at the white dwarf magnetic poles. However, more direct information about the size of the streams and the distribution of (luminous) matter in the magnetosphere, has only been revealed recently by Doppler imaging of a few systems (Diaz \\& Steiner 1994, Shafter et al.~1995, Schwope et al.~1997). Methods which allow us to infer the distribution of luminous matter in the magnetosphere are ideal complements to those which allow the study of distributions of 'dark' (i.e.~photoabsorbing) matter, which shape the EUV/X-ray light curves. We discuss here two examples. Finally, the donor stars are addressed briefly. We show an example of trailed spectra of photospheric absorption lines. The Doppler image shows a ``half-star'' (i.e.~just one hemisphere) only, demonstrating the large effects of X-ray illumination on the photospheric structure of the secondary star. Preliminary results of a broad multi-wavelength, multi-epoch observational campaign of the eclipsing polar HU Aqr (RX\\,J2107-0518, Hakala et al.~1993, Schwope et al.~1993) provides the backbone of this paper. Data of similar systems will be shown and discussed at appropriate places. A full analysis of the data will appear in a series of papers in the near future. ", + "conclusions": "Using some selected systems we have demonstrated how multi-epoch, multi-wavelength observations have allowed new and detailed insights into the accretion phenomena in AM Herculis stars. In the foreseeable future mapping of even smaller substructures will become possible. XMM will probably allow to measure the sizes of hard and soft X-ray emissions separately, 8m-class optical telescopes will help to explore the accretion stream in the magnetosphere in even greater detail by constructing eclipse maps in velocity bins." + }, + "9708/astro-ph9708191_arXiv.txt": { + "abstract": "We present analysis of the evolution of dark matter halos in dense environments of groups and clusters in dissipationless cosmological simulations. The premature destruction of halos in such environments, known as the {\\it overmerging}, reduces the predictive power of $N$-body simulations and makes difficult any comparison between models and observations. We analyze the possible processes that cause the overmerging and assess the extent to which this problem can be cured with current computer resources and codes. Using both analytic estimates and high resolution numerical simulations, we argue that the overmerging is mainly due to the lack of numerical resolution. We find that the force and mass resolution required for a simulated halo to survive in galaxy groups and clusters is extremely high and was almost never reached before: $\\sim 1-3$ kpc and $10^8-10^9\\Msun$, respectively. We use the high-resolution Adaptive Refinement Tree (ART) $N$-body code to run cosmological simulations with the particle mass of $\\approx 2\\times 10^8h^{-1}{\\ }{\\rm M_{\\odot}}$ and the spatial resolution of $\\approx 1-2h^{-1}{\\ }{\\rm kpc}$, and show that in these simulations the halos do survive in regions that would appear overmerged with lower force resolution. Nevertheless, the halo identification in very dense environments remains a challenge even with the resolution this high. We present two new halo finding algorithms developed to identify both isolated and satellite halos that are stable (existed at previous moments) and gravitationally bound. To illustrate the use of the satellite halos that survive the overmerging, we present a series of halo statistics, that can be compared with those of observed galaxies. Particularly, we find that, on average, halos in groups have the same velocity dispersion as the dark matter particles, i.e. do not exhibit significant velocity bias. The small-scale (100~kpc -- 1~Mpc) halo correlation function in both models is well described by the power law $\\xi\\propto r^{-1.7}$ and is in good agreement with observations. It is slightly antibiased ($b\\approx 0.7-0.9$) relative to the dark matter. To test other galaxy statistics, we use the maximum of halo rotation velocity and the Tully-Fisher relation to assign the luminosity to the halos. For two cosmological models, a flat model with the cosmological constant and $\\Omega_0=1-\\Omega_{\\Lambda}=0.3, h=0.7$, and a model with a mixture of cold and hot dark matter and $\\Omega_0=1.0, \\Omega_{\\nu}=0.2, h=0.5$, we construct the luminosity functions and evaluate mass-to-light ratios in groups. Both models produce luminosity functions and the mass-to-light ratios ($\\sim 200-400$) that are in a reasonable agreement with observations. The latter implies that the mass-to-light ratio in galaxy groups (at least for $M_{vir}\\lesssim 3\\times 10^{13}h^{-1}{\\ }{\\rm M_{\\odot}}$ analyzed here) is not a good indicator of $\\Omega_0$. ", + "introduction": "It is generally believed that the dark matter (DM) constitutes a large fraction of the mass in the Universe and thus significantly affects, and on some scales dominates, the process of galaxy formation. Observational evidence for large fractions of DM in galaxies, groups, and galaxy clusters ranges from flat rotational curves of spiral (e.g., Faber \\& Gallagher, 1979; Rubin et al. 1985; Persic, Salucci, \\& Stel 1996; Courteau \\& Rix, 1997) and X-ray emission and mass-to-light ratios of elliptical (Forman et al.1985; Rix 1997; Brighenti \\& Mathews 1997) galaxies to the baryon fractions in clusters of galaxies (White et al. 1993; Evrard 1997). For galaxies, the extent of the DM halos estimated using satellite dynamics, is $\\sim 0.2-0.5h^{-1}$ Mpc\\footnote{Throughout this paper we assume that the present-day Hubble constant is $H_0=100h{\\ }{\\rm km s^{-1} Mpc^{-1}}$.} (Zaritsky \\& White 1994; Carignan et al.1997). A convincing evidence for substantial amounts of dark matter even in the very inner regions of galaxies comes from the recent HI studies of the dwarf and low surface brightness (LSB) galaxies. The observed amounts of stars and gas in some of these galaxies can account for less than $10\\%$ of the observed rotational velocities at the last measured point of the rotation curves. (e.g., Carignan \\& Freeman 1988; Martimbeau, Carignan, \\& Roy 1994; de Blok \\& McGaugh 1997). If the observed galaxies have large DM halos, then $N$-body simulations can, in principle, be used to predict distribution of the dark matter component, to associate the simulated DM halos with galaxies, and to predict the bulk properties of these galaxies such as position, mass, and size. One should be able then to make predictions about spatial distribution and motion of these simulated galaxies and compare these predictions with corresponding observations. Unfortunately, the dissipationless numerical simulations have been consistently failing to produce galaxy-size dark matter halos in dense environments typical for galaxy groups and clusters (e.g., White 1976; van Kampen 1995; Summers, Davis, \\& Evrard 1995; Moore, Katz, \\& Lake 1996). This apparent absence of substructure in the virialized objects, known as {\\it the overmerging problem}, reflects the fact that simulated galaxies seem to merge much more efficiently in comparison with real galaxies in groups and clusters. In the central regions of a cluster ($\\sim 500$ kpc ), the ``overmerging'' erases not only large-scale substructure, but also any trace of small halos that could be associated with ``galaxies''\\footnote{The term ``galaxy'' traditionally refers to luminous observed objects (i.e. to clumps of stars and gas, not the DM), which, possibly, are embedded in a considerably larger DM ``halo''. The term ``halo'', however, is rather general. We can use this term to indicate a galaxy cluster, group, or a galaxy-size halo. In some cases, we want to make a clear distinction between these. We will thus use terms ``simulated galaxies'' or ``galaxy-size halo'' to indicate the DM halos formed in the simulations which could be associated with places where luminous baryons could reside.}, leaving a smooth giant lump of dark matter. The overmerging problem was traditionally explained by the lack of dissipation in $N$-body simulations (e.g., Katz, Hernquist \\& Weinberg 1992; Summers et al. 1995). Indeed, the DM halos are much larger than baryonic extent of the galaxies due to the dissipational nature of the latter. The radiative cooling, for example, allows baryonic component to sink into the center of the DM halo where it forms a compact, tightly bound object. In dense environments the large DM halo can be easily stripped by the tidal field of a galaxy cluster or group, whereas the more compact and denser gas clump may survive (Summers et al. 1995). Although it is clear that to produce a realistic galaxy we need to include the energy dissipation by baryons, it is not clear whether the dissipation is vital for the halo survival in a cluster. Two arguments can be presented against the traditional explanation for the overmerging. First, if the dissipation helps galaxies to survive in clusters, then galaxies should be dominated by baryons at all scales within their visible extent. Most of the observed galaxies, however, appear to have a substantial fraction of DM inside their optical radius (e.g., Persic et al. 1996). The survival of a galaxy dominated at its optical radius by the dark matter will depend mostly on the dark matter, not on the baryons. The DM dominated dwarfs must have been tidally disrupted, but dwarf galaxies are observed in clusters (e.g. Smith, Driver, \\& Phillipps 1997; Lopez-Cruz et al. 1997). Second, as we argue in this paper, even in the absence of the baryons DM halos are dense enough to survive inside clusters and to be identified, provided that simulations have sufficient resolution {\\em in both} mass and force (similar conclusion was reached by Moore et al. 1996). The main goal of this paper is to demonstrate that with a sufficient computational effort the overmerging problem can be tamed (at least to some extent) even in the purely dissipationless simulations. The computational costs are higher than cost of an average cosmological $N$-body simulation. However, they may be considerably lower than computational expense of the corresponding $N$-body$+$hydro simulations. This especially true in the case of large-volume ($\\sim 50-100h^{-1}{\\ }{\\rm Mpc}$) simulations required to quantify the statistical properties of the ``galactic'' populations such as correlation function, pairwise velocity dispersions etc. To understand how the problem should be dealt with, it is important to understand what processes lead to the overmerging. Several recent studies have addressed this question from different viewpoints and using different numerical and analytical techniques. Thus, for example, van Kampen (1995) studied formation of galaxy clusters in purely dissipationless simulations and concluded that two-body relaxation and tidal disruption are primarily responsible for the overmerging. He found, however, that the two-body effects are important only for the smallest halos ($\\lesssim 30$ DM particles), in quantitative agreement with experiments of Moore et al. (1996). The latter study addressed also effects of particle-halo and halo-halo heating on the survival of DM halos. The authors concluded that particle-halo heating does not pose a problem as long as DM particle mass is $\\lesssim 10^{10} M_{\\odot}$, but halo-halo heating may be important if force resolution is not adequate ($\\gtrsim 10 {\\rm kpc}$). The general conclusion of Moore et al. is that the overmerging problem is due mainly to tidal heating by the cluster and halo-halo heating, both effects being enhanced by poor force resolution. They note also that these effects depend crucially on the density structure of DM halos. The density profile of dark matter halos in the CDM models is now known reasonably well. Navarro, Frenk \\& White (1995, 1996, 1997, hereafter NFW) gave an analytical fit that describes with a reasonable accuracy the density profiles of DM halos formed in the standard cold dark matter scenario over large range of scales and masses. The analytical form of the density profile advocated by NFW, allows one to estimate tidal effects and effects of dynamical friction analytically. We present these estimates in \\S 2. Two recent numerical studies (Tormen, Diaferio \\& Syer 1998; and Ghigna et al. 1998) present evidence that higher resolution significantly aleviates the overmerging problem in dissipationless simulations. Both studies conclude that many halos survive for a long time after falling onto a large halo, although overmerging problem persists to a certain degree in the central dense region of the large cluster-size halo. Tormen et al. also point out that there are real physical processes which would lead to the erase of substructure even in the ideal very high-resolution simulation. Namely, the dynamical friction drives massive satellites to the cluster core where they get tidally stripped and quickly disrupted. This effect is, of course, real and should be distinguished from artificial overmerging caused by insufficient numerical resolution. We will address these issues in \\S 2. Two goals of the study presented in this paper are 1) to make an approximate estimates of the effects leading to the overmerging for the halos with the NFW density distribution; and 2) to demonstrate that dissipationless simulations with sufficiently high resolution in force and mass are affected by the overmerging to a considerably lesser degree. In other words, we present an attempt to estimate to which extent the overmerging problem can be solved with numerical resolution that can be achieved with current codes and computational resources. The goal is worthwhile. If the problem can be minimized at a computational cost that is not prohibitive, the dissipationless simulations can be used for a direct study of the statistical properties of ``galaxies''. This may include studies of their spatial distribution, velocity field, environmental effects etc. This may allow us to make a step towards solution of the long-standing and particularly important issue of the galactic bias. The overall plan of the paper is as follows. In \\S 2 we discuss the numerical and physical effects which may lead to the erasure of substructure inside the dense massive halos. Specifically, we present analytical estimates for tidal disruption of dark matter halos and effects of dynamical friction assuming that halos are described by the NFW density profile. We use these calculations to make a rough estimate of what resolution would be required to minimize the overmerging. Numerical simulations and cosmological models are discussed in \\S 3. In \\S 4 we discuss difficulties associated with identification of dark matter halos in very dense regions. We describe two new halo finding algorithms developed to handle the halos in such environments. In \\S 5 we present results of high-resolution dissipationless simulations, illustrate the performance of the new halo finding algorithms and discuss the degree to which the simulations are affected by the overmerging. Our conclusions are illustrated using well-known statistics such as two-point correlation function, velocity bias, luminosity function and $(M/L)$ ratio. The conclusions are summarized in \\S 6. ", + "conclusions": "We have presented arguments that the overmerging problem is mostly due to inability of a numerical code to provide a sufficient numerical resolution to prevent tidal destruction of galaxy-size halos by the tidal forces of group or cluster and have estimated the resolution needed to prevent such disruption. We argue that although energy dissipation by the baryonic component helps galaxies to survive in clusters, at distances $\\gtrsim (50-70)\\kpch$ from the cluster center the gravity of the dark matter alone is enough to keep them alive. The main result of this work is estimate of the numerical resolution needed to overcome the overmerging problem. The results of our analytic estimates and numerical experiments show that although it is feasible to overcome overmerging in pure $N$-body simulations, resolution required to avoid artificial destruction of galaxy-size halos (mass $\\gtrsim 10^{11}h^{-1} {\\rm M_{\\odot}}$) is quite high. For viable CDM models and realistic halo profiles this resolution is $\\lesssim 2h^{-1} {\\rm kpc}$ in force and $\\lesssim 10^{9}h^{-1} {\\rm M_{\\odot}}$ in mass. This requires simulations of $>10^7$ particles with dynamic range of $10^5$ in spatial resolution for statistically significant cosmological volumes ($\\sim 100 {\\rm Mpc}$), which remains challenging with the current computers and numerical codes. {\\pspicture(0.5,-1.5)(13.0,12.) \\rput[tl]{0}(-0.75,12.){\\epsfxsize=10.5cm \\epsffile{fig18.ps}} \\rput[tl]{0}(0.5,1.25){ \\begin{minipage}{8.7cm} \\small\\parindent=3.5mm {\\sc Fig.}~18.--- Dependence of the galaxy correlation function on mass in the $\\Lambda$CDM simulation of $30\\Mpch$ box. The correlation function increases with the rotational velocity, but all curves show the same tendency: positive bias on small scales, slight antibias on $(100-1000)\\kpch$ scales, and no bias on larger scales. Absolute values of the correlation functions are affected by the finite box size. \\end{minipage} } \\endpspicture} This makes simulations focused on the individual clusters of the type presented in Ghigna et al. (1998) a viable alternative. Unfortunately, even in the case of sufficient resolution, when halos do survive, the identification of the DM halos in the cores of galaxy groups and clusters in purely dissipationless simulations remains a challenge. In the environments that dense, most of halo's dark matter will be be tidally stripped, which makes it difficult to identify the leftover on the very dense, smooth background of high-velocity dark matter particles streaming around and through the halo. We have presented two new halo finding algorithms designed to identify satellite halos located inside the virial radius of a more massive host halo: the hierarchical friends-of-friends and bound density maxima algorithms. Both of our algorithms find practically the same halos, which are stable (existed at previous moments) and gravitationally bound. To exploit the fact that overmerging is (at least to a certain degree) ``beaten'' in our simulation, we consider several statistics of galaxy-size halos in our simulations and compare them to the corresponding observed statistics of galaxies. We use a simple scheme, based on the empirical Tully-Fisher relation, to assign a luminosity to the DM halos. The luminosity function of ``galaxies'' (i.e., galaxy-size halos assigned a luminosity) in the $\\Lambda$CDM model reproduces the luminosity function of the CfA catalog (Marzke et al. 1994) reasonably well. Both the simulations and the CfA catalog have an upturn in the number of faint galaxies ($m_B>-17$). However, magnitudes of faint ``galaxies'' in the simulations rely on a highly uncertain extrapolation of the Tully-Fisher relation and on uncertain assumption about the fraction of elliptical galaxies at these magnitudes. The number of ``galaxies'' predicted by the CHDM simulation is significantly higher than in the case of the $\\Lambda$CDM simulation with the same initial random numbers. We failed to produce as nice a fit to the observational data as for the $\\Lambda$CDM simulation. At this stage it is difficult to judge if this is a significant problem for the model or not. Due to the small size of our simulation boxes, one may argue that simulations with a large box will tend to produce lower luminosity function keeping at the same time the M/L of galaxy groups intact. Larger simulations are needed to clarify the situation. The mass-to-light ratios of galaxy groups in the simulations $\\sim (200-400)h^{-1}$ match those observed reasonably well. It was argued (e.g., Bahcall, Lubin, \\& Dorman 1995) that dynamics of galaxy groups favors the low-$\\Omega$ Universe. Our results show that mass-to-light ratios of groups of mass $\\lesssim 3\\times 10^{13}h^{-1}{\\ }{\\rm M_{\\odot}}$ is insensitive to the matter density. The halos in the CHDM model are clustered more strongly than the dark matter and one cannot save the argument for a low-$\\Omega$ Universe by assuming that groups in the CHDM model have too large fraction of galaxies. It seems that groups occupy too small fraction of the volume and thus their M/L ratios are not representative for the Universe as a whole. Comparison of the halo and matter correlation functions indicates that halos are antibiased on 100~kpc -- 1~Mpc scales. The antibias of the magnitude 0.7--0.9 found in the simulations is needed for the $\\Lambda$CDM model to be compatible with observational data on the power spectrum in the range of wavenumbers $k=(0.1-1)h\\Mpc^{-1}$ (Klypin et al. 1996; Smith et al. 1998). We attribute the antibias to the dynamical friction in groups of galaxies. The friction tends to drag some galaxies to the very central part of groups where they merge the central galaxy and disappear (see Kravtsov \\& Klypin 1999 for a more detailed analysis). Results of this paper can be used in design of the future numerical simulations. We have shown that efficient halo finding algorithms can be developed to identify gravitationally bound satellite halos inside the virial radius of the other halos. Our analytical estimates and numerical experiments show that the numerical resolution required to overcome the overmerging, although quite high, can be achieved with current numerical codes and computer hardware. The main challenge is thus purely computational. This is also true for the simulations that include dissipative hydrodynamics; while alleviating or obliterating some of the problems of dissipationless simulations, they are computationally more intensive. Both numerical approaches have a number of caveats and potential biases, which could only be avoided with inclusion of more realistic physics. The latter appears to be unavoidable, because we cannot reliably predict observed galaxy properties (and hence mimick the selection criteria of the observational catalogs) without realistic physics. Fast increase in computational capability of modern computers and recent developments of new efficient numerical algorithms make the perspective for advances in this direction look good." + }, + "9708/astro-ph9708103_arXiv.txt": { + "abstract": "Diffraction-limited spectroscopy with adaptive optics (AO) has several advantages over traditional seeing-limited spectroscopy. First, high resolution can be achieved without a large loss of light at the entrance slit of the spectrograph. Second, the small AO image width allows the cross-dispersed orders to be spaced closer together on the detector, allowing a large wavelength coverage. Third, AO spectrograph optics are slow and small, costing much less than for a traditional spectrograph. Fourth, small AO images provide high spatial resolution. Fifth, scattered light is less problematic. And last, the small entrance slit of the spectrograph can get rid of much of the sky background to obtain spectra of faint objects. We have done theoretical calculations and simulations for infrared spectroscopy at the MMT 6.5 m with laser guide star AO, which provides almost full sky coverage. The results show we can expect 40-60\\% of the photons from a unresolved source within 0.2 arcsec diameter circle for J, H, K, L and M bands under typical atmospheric seeing condition at 2.2 micron (r$_0$ = 1.0 m, t$_0$ = 21 ms, $\\theta_0$ = 15 arcsec and d$_0$ = 25 m). Therefore, the spectrograph entrance slit size should match the 0.2 arcsec image to obtain high throughput. Higher resolution can be achieved by narrowing down the slit size to match the diffraction-limited image core size of about 0.1 arcsec in the infrared. However, the throughput will be correspondingly reduced by a factor of two. Due to the limited atmospheric isoplanatic angle in the J, H and K bands, the encircled photon percentage within 0.2 arcsec diameter drops from 40-60\\% when the object is at the laser pointing direction to 20-40\\% when the object is about 30 arcsec away from the laser direction. Therefore, the useful field of view for AO multiple object spectroscopy is about 60 arcsec. Further studies of IR background (sky and thermal) and IR detector performance show that spectral resolution of R = 2,000 can take full advantage of AO images without much penalty due to the dark current of the IR detector and IR OH sky emission lines. We have also studied natural guide star AO spectroscopy. Though sky coverage for this kind of spectroscopy at the MMT 6.5 m is very limited, a bright star provides much better performance than the laser guide star AO spectroscopy. About 40-70\\% photons are concentrated within 0.1 arcsec diameter for guide stars brighter than 13 magnitude. Therefore, higher resolution and high throughput can be obtained simultaneously, given a bright enough natural guide object. The field-of-view for multiobject spectroscopy using a natural guide star is similar to that for laser guiding. ", + "introduction": "Adaptive optics (AO) promises revolutionary advances in imaging power for ground-based optical and infrared astronomical telescopes by eliminating the wave-front distortion caused by atmospheric turbulence. The AO corrected images will be nearly diffraction-limited, which is about a factor of ten times smaller than that limited by the atmospheric seeing for current 4 m class telescopes. For the largest of the new generation of telescopes, the most dramatic gain is possible, permitting an imaging performance of almost two orders of magnitude (100 times). Though adaptive optics has a big impact on improving ground-based telescope image quality, it cannot provide ideal diffraction-limited images in principle, due to the limited photon flux available from the reference source, finite response time and subaperture size of the AO systems (Sandler et al. 1994). The AO corrected images therefore consist of two components: a diffraction-limited core and a broad seeing-limited halo (Beckers 1993), which make the design of AO instruments different from that of seeing limited instruments. The much sharpened AO images have two main applications in astronomy research; imaging and spectroscopy. The two are closely related but not the same. The main focus of direct high resolution imaging is to sharpen the diffraction limited image core, to maintain stable uniform point spread function (PSF) in both spatial and temporal domains. On the other hand, the most concern of the AO spectroscopy is the flux concentration. The different demands for these two different applications determine different instrument design parameters. The application of adaptive optics in astronomy is still in its early phase and the design of AO optimized instruments, especially the spectrographs, is a new territory being opened for exploration. In seeing limited domain, the best resolution of a spectrograph is coupled to the telescope diameter, the larger the aperture size, the lower the spectral resolution for the normal available grating size. This coupling limited the best spectral resolution of traditional spectrographs to R $\\sim$ 50,000 for 4 m class telescopes (Vogt \\& Schroeder 1987). In order to obtain higher resolution, all kind of tricks such as image slicers, pupil slicers, grating mosaics etc have been applied and resulted in very large and expensive spectrographs at the Nasmyth or Coude focuses (e.g. Diego et al. 1995; Vogt et al. 1994; Tull et al. 1994). However, in the AO diffraction limited domain, because the AO corrected image size (i.e. diffraction limited core size) decreases proportionally with telescope aperture size, the coupling of spectrograph size with telescope aperture is removed. Very high resolution spectrographs can be made from normal size gratings. As the results, the next generation AO optimized high resolution spectrographs will have smaller overall scale, higher efficiency and also cost much less. As the first demonstration of next generation AO spectrographs, we have built a prototype AO cross-dispersed echelle spectrograph with a 125x250 mm$^2$ R2 echelle grating at Steward Observatory and tested at Starfire Optical Range 1.5 m AO telescope. The spectrograph can provide spectral resolution up to R = 700,000 (Ge et al. 1996). Because of the much smaller image size, a large amount of cross dispersed orders can be packed and recorded on the detector, and thus a factor of 100 times larger wavelength coverage over similar resolution traditional spectrographs was achieved (e.g. Diego et al. 1995; Lambert et al. 1990; Ge et al. 1996). And because of the much smaller image size, AO spectrographs can record astronomical phenomena of much smaller scale structure (Bacon et al. 1995). Further, smaller entrance slit used in the AO spectrographs can help to block most sky background, especially in the IR where the sky background is about 100 times brighter than in the visible, so much fainter objects can be observed. In this paper, we will first set out the types of error that arise in the adaptive optics systems and how they together affect overall performance of AO spectroscopy. Then we will use the MMT 6.5 m AO system under construction as an example to introduce the results from the semi-empirical analytical calculations and direct Monte Carlo simulations and relate these computational results to the design of AO spectrographs. ", + "conclusions": "" + }, + "9708/astro-ph9708273_arXiv.txt": { + "abstract": "The presence of low mass, degenerate secondaries in millisecond pulsar binaries offers the opportunity to determine an age for the binary system independent of the rotational properties of the pulsar. To this end, we present here a detailed calculation of the evolution of a grid of low mass ($< 0.5 \\, \\rm M_{\\odot}$) helium core white dwarfs. We investigate the effects of different Hydrogen layer masses and provide results for well-known optical band-passes. We supplement the OPAL opacity calculations with our own calculations for low effective temperatures ($ T_{\\rm eff} < 6000 \\, K$) and also provide fitting formulae for the gravity as a function of mass and effective temperature. In paper~II we shall apply these results to individual cases. ", + "introduction": "Studies of binary pulsars (see Phinney \\& Kulkarni 1994 for references) and close double degenerate systems (Marsh, Dhillon \\& Duck 1995) have discovered a number of low mass objects. The masses of these optically faint stars are estimated to be $\\sim$ 0.1 - 0.5 $\\rm M_{\\odot}$. Thus, it is believed that these are helium core white dwarfs, which are not massive enough to ignite core helium burning and burn to carbon (Sweigart \\& Gross 1978; Mazzitelli 1989). Such stars arise in binaries because the progenitor is disrupted by Roche lobe overflow at some point in its natural evolution, losing its envelope and leaving behind a low-mass, degenerate helium star (Kippenhahn, Kohl \\& Weigart 1967). Many millisecond pulsar binaries contain such dwarfs and, as such, their cooling ages offer an estimate of the age of the system, independent of the pulsar spin-down age and thus an interesting check of the traditional assumptions made about the ages of millisecond pulsars. Our aim in this paper is to present an extensive grid of cooling models covering the parameter space of low-mass helium core white dwarfs, much in the spirit of previous studies of the more ubiquitous carbon/oxygen core white dwarfs ( Wood~1992; D'Antona \\& Mazzitelli 1990 and references therein). In paper~II we apply these models to the observations of the low-mass binary pulsar population. In section~\\ref{Cooling} we briefly review the physical mechanisms which contribute to the white dwarf cooling process and discuss their relative importance. Section~\\ref{LCS} describes our calculations of the low-temperature opacities necessary to obtain accurate cooling sequences for the oldest systems. In section~\\ref{Code} we describe our numerical model and the tests of the code against other models from the literature. Finally, in section~\\ref{ResultsI} we present our cooling sequences and describe the details of the cooling models. ", + "conclusions": "We have presented a set of cooling sequences for low mass helium white dwarfs of different masses and with different masses of surface hydrogen. We provide blackbody absolute magnitudes and surface gravity - effective temperature relations as an aid to the analysis of future observations. In paper~II we shall apply these models to the optical observations of the companions to millisecond pulsars in order to derive cooling ages. The authors would like to thank Marten van Kerkwijk and Yanqin Wu for lengthy discussions about white dwarf physics and Glenn Soberman for helping with the initial conditions for our models. The generosity of messrs D. Saumon, G. Fontaine, I. Mazzitelli, F. Rogers and C. Iglesias in providing their microphysical results is also appreciated. We would also like to thank the referee, Dr. D'Antona, for some insightful comments on the atmospheric physics. \\begin{appendix}" + }, + "9708/astro-ph9708209_arXiv.txt": { + "abstract": "We have searched the Fourth Catalogue of Nearby Stars for halo stars and identified 15 subdwarfs and a high velocity white dwarf in the solar neighbourhood. This search was motivated mainly by the recent determinations of MACHO masses of about 0.5$\\cal M_\\odot$, which are typical for halo dwarfs. The local mass density of these stars is 1.5$\\cdot$10$^{-4}$ $\\cal M_\\odot$ pc$^{-3}$, which is only 3\\% of the current estimate of the local mass density of the MACHO population. We compare the local density of subdwarfs with constraints set by HST observations of distant red dwarfs. Using models of the stellar halo with density laws that fall off like $r^{-\\alpha}$, $\\alpha$ = 3.5 to 4, we find that the HST constraints can only be matched, if we assume that the stellar halo is flattened with an axial ratio of about 0.6. The non--detection of the analogues of MACHOs in the solar neighbourhood allows to set an upper limit to the luminosity of MACHOs of $M_{\\rm B} >$ 21 magnitudes. ", + "introduction": "The stellar halo of the Galaxy has been studied with renewed interest since the MACHO (Alcock et al.~1996) and EROS (Ansari et al.~1996) collaborations have reported results of their campaigns of observing micro-lensing events towards the LMC. These indicate that MACHOS have masses typically of half a solar mass and contribute about one half to the mass budget of the halo of the Milky Way. Such masses are typical for red and white dwarfs, implying that MACHOs might possibly be members of the stellar halo instead of the dark halo of the Galaxy, as was previously thought. Bahcall and collaborators (Bahcall et al.~1994, Flynn, Gould \\& Bahcall 1996, Gould, Bahcall \\& Flynn 1997) have made by very deep star counts based on HST data {\\em in situ} measurements of the space density of such stars in regions, where the micro-lenses are expected to be physically located. Their conclusion is that halo dwarfs contribute only a few percent to the mass of the halo and are thus unlikely MACHO candidates. Graff \\& Freese (1996) infer from the same data that the local mass density of halo dwarfs is even less than one percent of the combined local densities of the dark and stellar halos. On the other hand, halo dwarfs are directly observed in the solar neighbourhood. Liebert (1995) discusses the results of the USNO parallax programme for red dwarfs and finds that the number of halo dwarfs in that sample is consistent with the HST data. We have followed a complementary approach. Recently the Third Catalogue of Nearby Stars (CNS3) has been completed at the Astronomisches Rechen-Institut (Jahrei{\\ss} \\& Gliese 1997) and, after HIPPARCOS (ESA 1997) parallaxes and proper motions have become available, is presently updated to the Fourth Catalogue of Nearby Stars (CNS4; Jahrei{\\ss} \\& Wielen 1997). \\begin{figure*}[t] \\begin{center} \\leavevmode \\epsffile{vihip.eps} \\hfill \\parbox[b]{5cm}{\\caption{Colour--Magnitude diagram of nearby stars with $\\sigma_{\\rm M_{\\rm V}} \\leq$ 0.3 mag.}}% \\label{CMD}% \\end{center} \\end{figure*} This provides the now most complete inventory of the solar neighbourhood up to a distance of 25 pc from the Sun and allows a new determination of the local density of halo stars. Preliminary results based on the CNS3 have been reported by Fuchs \\& Jahrei{\\ss} (1997). \\begin{table*}[htb] \\caption{Subdwarfs in the CNS4.} \\vspace{0.4cm} \\begin{center} \\begin{tabular*}{12.0cm}{|l@{\\extracolsep\\fill}@{\\hspace{0.3cm}}r @{\\hspace{0.mm}}lrrrrrrr@{\\hspace{0.mm}}l|} \\hline \\rule[0mm]{0mm}{4mm} & \\multicolumn{2}{c}{d} & \\multicolumn{1}{c}{$M_V$} & \\multicolumn{1}{c}{ $V-I_c$} & U & \\multicolumn{1}{c}{V} & W & $[$Fe/H$]$ & \\multicolumn{2}{c|}{$\\cal M$}\\\\[0.5ex] &\\multicolumn{2}{c}{[pc]} & [mag] & \\multicolumn{1}{c}{[mag]} & & [km/s] & & {[dex]} & \\multicolumn{2}{c|}{ [$\\cal M_\\odot$]}\\\\[0.5ex] \\hline Gl\\,191 & 3.&9$^{\\rm H}$& 10.89 & 1.98 & 20 & --288 & --53 &$<-$1 & 0.&2\\\\ [0.5ex] Gl\\,299 & 6.&8 & 13.65 & 2.92 & 107 & --126 & --49 &$<-$1 & 0.&1\\\\[0.5ex] Gl\\,53A & 7.&6$^{\\rm H}$ & 5.78 & 0.79 & --43 & --158 & --35 & --0.6 & 0.&73 \\\\[0.5ex] Gl\\,53B & 7.&6$^{\\rm H}$ & 11.00 & & & & & & 0.&2 \\\\[0.5ex] Gl\\,451A & 9.&2$^{\\rm H}$ & 6.64 & 0.85 & 280 & --159 & --14 & $<-$1 & 0.&6 \\\\[0.5ex] Gl\\,699.1 & 15.&6 & 13.34 & 0.50 & --149 & --294 & --42 & & 0.&6 \\\\[0.5ex] GJ\\,1062 & 16.&0 & 12.00 & 2.21 & 74 & --219 & -5 & $<-$1 & 0.&15\\\\[0.5ex] Gl\\,781 & 16.&5 & 10.91 & 2.03 & 106 & --47 & 35 & $<-$1 & 0.&2 \\\\[0.5ex] Gl\\,158 & 18.&5$^{\\rm H}$ & 7.17 & 0.93 & --41 & --190 & 21 & $<-$1 & 0.&6 \\\\[0.5ex] LHS\\,3409 & 20.&3 & 13.59 & 2.76 & --31 & --101 & 43 & low & 0.&1\\\\[0.5ex] WO\\,9371 & 22.&6$^{\\rm H}$ & 10.43 & 1.87 & 135 & --309 & 124 & low & 0.&26\\\\[0.5ex] WO\\,9722 & 23.&9 & 11.39 & 2.05 & 264 & --215 & --93 & low & 0.&17\\\\[0.5ex] LHS\\,375 & 24.&0 & 13.78 & 2.27 & 30 & --183 & 153 & $<-$1 & 0.&1 \\\\[0.5ex] GJ\\,1064A & 24.&5$^{\\rm H}$ & 6.22 & 0.86 & --96 & --115 & --76 & --1 & 0.&70 \\\\[0.5ex] GJ\\,1064B & 24.&5$^{\\rm H}$ & 6.82 & 1.02 & --96 & --115 & --76 & $<-$1 & 0.&63 \\\\[0.5ex] LHS\\,2815 & 25.&0$^{\\rm H}$ & 5.99 & 0.72 & --29 & --47 & --37 & $<-$1 & 0.&72 \\\\[0.5ex] \\hline \\multicolumn{11}{c}{$^{\\rm H}$HIPPARCOS parallax and proper motions}\\\\[0.5ex] \\end{tabular*} \\end{center} \\end{table*} ", + "conclusions": "\\subsection{Local density} The local mass density of halo stars can be estimated in various ways. In Fig.~3 we show density estimates, which have been calculated from the cumulative distribution of the stars in our sample, i.e.~by adding up the masses of stars within given distances and dividing by the corresponding spherical volumes. The stars within 10 pc give rather high density estimates of the order of 10$^{-3}$ $\\cal M_\\odot$ pc$^{-3}$. Wielen \\& Jahrei{\\ss} (Wielen 1976) found similar values, when evaluating the second edition of the Gliese catalogue. A careful analysis of the spatial distibution of the stars shows, however, that the high density peaks are of purely statistical nature and certainly not realistic. We find a plateau in the density run, when we consider the stars within the 20 pc sphere (cf.~Fig.~3). Beyond that the density drops off, probably because the catalogue becomes incomplete. The most likely value of the local mass density lies according to our determination in the range 1.5 to 1$\\cdot 10^{-4}$ $\\cal M_{\\odot}$ pc$^{-3}$. We note that just outside the nominal boundary of the CNS4 there are two further subdwarfs, LHS\\,173 and GL\\,788.2, at distances 25.5 pc and 25.6 pc, respectively. If they are included, the mass density is 1$\\cdot 10^{-4} \\cal M_\\odot$ pc$^{-3}$, which, in our view, is a firm lower limit to the local mass density of halo stars.. The local density of the dark halo (Bahcall, Schmidt \\& Soneira 1983, Gates et al.~1996) is about 9$\\cdot 10^{-3}$ $\\cal M_{\\odot}$ pc$^{-3}$, so that the local density of halo stars is about 1.7\\% of the halo density or 3\\% of the local MACHO density as determined by the MACHO collaboration. \\begin{figure}[htb] \\begin{center} \\leavevmode \\epsffile{halo.eps} \\caption{Mass density of nearby halo dwarfs as function of the distance from the Sun. Within 10 pc for some of the stars only the positions are indicated. The errorbars indicate the statistical uncerntainties.} \\label{dens} \\end{center} \\end{figure} We have broken down our sample of subdwarfs with respect to absolute magnitudes. The resulting `luminosity function' is compared with the subdwarf luminosity function obtained by Dahn et al.~(1995) for stars with $M_{\\rm V} >$ 10 mag in Table 2. Both agree well within statistical uncertainties. \\begin{table}[htb] \\caption{Luminosity function of red subdwarfs.} \\vspace{0.4cm} \\begin{center} \\begin{tabular*}{7cm}{|c@{\\extracolsep\\fill}ccc|} \\hline \\rule[0mm]{0mm}{4mm} $M_{\\rm V}$ & N$_{\\rm CNS4}$ & $\\nu_{\\rm CNS4}$ & $\\nu_{\\rm Dahn}$\\\\[0.5ex] [mag] & & \\multicolumn{2}{c|}{[$10^{-5}$ mag$^{-1}$ pc$^{-3}$]} \\\\[0.5ex] \\hline 10 & 0 & 0 & 6 \\\\[0.5ex] 11 & 3 & 16 & 10 \\\\[0.5ex] 12 & 1 & 5 & 6 \\\\[0.5ex] 13 & 0 & 0 & 4 \\\\[0.5ex] 14 & 1 & 5 & 1 \\\\[0.5ex] 14.5 & 0 & 0 & 0.6 \\\\[0.5ex] \\hline \\end{tabular*} \\end{center} \\end{table} \\subsection{Matching the constraints by HST data.} It is interesting to confront the density of halo stars derived here with density estimates based on the HST data. Graff \\& Freese (1996) find in a detailed analysis of earlier HST data (Bahcall et al.~1994) a value about five times less than our estimate, which is partly due to their restriction to red dwarfs redder than $V-I \\approx$ 2 mag. Whereas these authors tried to infer the local density of halo stars from the star counts, we reverse the approach and make predictions of the expected number of stars in the star count field. Recently Flynn et al.~(1996) have reported an analysis of the Hubble Deep Field (HDF). The HDF has a size of 4.4 square arcminutes and is located at $l$ = 126$^\\circ$, $b$ = 55$^\\circ$. Flynn et al.~(1996) were able to identify stars down to a limiting magnitude of 26.3 mag in the $I$--band. In the following we assume for the population of halo stars a density law of the form \\begin{equation} {\\nu_{\\rm h}} = {\\nu}_{{\\rm h}_\\odot} \\left( \\frac{r_{\\rm c}^2+r_\\odot^2} {r_{\\rm c}^2+r^2} \\right)^\\alpha \\end{equation} with an arbitrarily chosen core radius of $r_{\\rm c}=$ 1 kpc. Such density laws are typical for other tracers of the halo population, such as RR~Lyrae stars (Wetterer \\& McGraw 1983), horizontal branch stars (Preston et al.~1983), in the Milky Way or external galaxies (Pritchet \\& van den Bergh 1994). The exponent $\\alpha$ lies typically in the range 3.5 to 4. Note that the dark halo is usually described by a density law of the form as in equation (1) with $\\alpha$ = 2. If the subdwarfs are observed locally at a density $\\nu_{{\\rm h}_\\odot}$, one predicts by integrating along the line of sight \\begin{eqnarray} N & = &{\\nu}_{{\\rm h}_\\odot} \\\\ & & \\cdot \\Omega \\int_{\\rm d_{min}}^{\\rm d_{max}}s^2 \\left( \\frac{r_{\\rm c}^2+r_\\odot^2} {r_{\\rm c}^2+r_\\odot^2+s^2-2sr_\\odot\\cos{l}\\cos{b}} \\right)^\\alpha ds \\nonumber \\end{eqnarray} stars in the star count field. $\\Omega$ is the angular area of the field and the minimum and maximum distances are determined by the limiting magnitudes. We concentrate first on the colour range $V-I$ = 1.8 to 3.5 mag. We have 4 stars in that range within 16.5 pc, where we believe our sample to be reasonably complete. If we project these into the cone towards the HDF, we obtain the star numbers summarized in Table 3. The limiting distances have been calculated for each star individually, \\begin{equation} d_{\\min,\\max} = 10^{0.2(I_{\\min,\\max}+(V-I)-M_{\\rm V}+5)} , \\end{equation} where, in order to avoid confusion with disk stars in the HDF, we adopt a lower limit of $I_{\\rm min}$ = 24.6 mag. As can be seen from Table 3 we predict 7 to 11 stars in the HDF, whereas Flynn et al.~(1996) have actually detected no star. \\begin{table}[htb] \\caption{Predicted Number of Stars in the HDF and the GS.} \\vspace{0.4cm} \\begin{center} \\begin{tabular*}{8cm}{|c@{\\extracolsep\\fill}cc|cc|} \\hline \\rule[0mm]{0mm}{4mm} & \\multicolumn{2}{c|}{$V-I >$ 1.8} & \\multicolumn{2}{c|}{$V-I <$ 1.8} \\\\[0.5ex] $\\alpha$ & c/a &n$_{\\rm HDF}$ & n$_{\\rm HDF}$ & n$_{\\rm GS}$ \\\\[0.5ex]\\hline 3.5 & 1 & 11 & 17 & 258 \\\\[0.5ex] 4 & 1 & 7 & 7 & 140 \\\\[0.5ex] 3.5 & 0.6 & 4 & 5 & 67 \\\\[0.5ex] 4 & 0.6 & 2 & 2 & 31 \\\\[0.5ex]\\hline \\end{tabular*} \\end{center} \\end{table} Despite the low number of stars involved, this discrepancy seems to be statistically significant. Several explanations might account for this discrepancy. First, the stellar halo might be much more irregular and lumpy than previously thought. Second, the stellar halo is almost certainly flattened (Wetterer \\& McGraw 1983, Preston, Shectman \\& Beers 1983, Pritchet \\& van den Bergh 1994) with an axial ratio around c/a $\\approx$ 0.6. If we take such a flattening into account in the halo model, \\begin{equation} {\\nu_{\\rm h}} = {\\nu}_{{\\rm h}\\odot} \\left( \\frac{r_{\\rm c}^2+R_\\odot^2} {r_{\\rm c}^2+R^2+z^2/(c/a)^2} \\right)^\\alpha\\,\\,, \\end{equation} where $R, z$ denote cylindric coordinates, we predict 2 to 4 stars in the HDF, i.~e.~star numbers which are statistically consistent with no star seen by Flynn et al.~(1996). The star counts in the HDF in the colour range $V-I <$ 1.8 mag can be interpreted in a similar way. Disk stars in this colour range are so bright that they would have to lie several kpc above the midplane to appear fainter than $I$ = 22 mag. Thus in order to avoid confusion with disk stars in the HDF, we consider only stars fainter than $I$ = 22 mag. In our sample there are 4 stars within 20 pc in this colour range. Their $V-I$ colours cluster around 0.9 and the white dwarf has $V-I$ = 0.5. Thus we define a colour range 0.5 $< V-I <$ 1.8, in which we compare the predictions with actually observed numbers of stars. If the stars of our sample are projected into the cone towards the HDF using spherical halo models, we predict 7 to 17 stars depending on the model, whereas Flynn at al.~(1996) have observed 4 stars. If we assume again a flattening of c/a = 0.6, the number of expected stars is consistent with the observed number. Gould et al.~(1997) have very recently investigated another field observed by HST. This is the so-called Groth strip (GS), which has an angular size of 114 square arcminutes and is located at $l$ = 96$^\\circ$ and $b$ = 60$^\\circ$. Its limiting magnitude is $I$ = 23.9 mag. Gould et al.~(1997) have very kindly made their data available to us, so that we were able to make star counts in this field. Due to the limiting magnitude of $I$ = 23.9 mag even the faintest stars in the GS with colours redder than $V-I$ = 1.8 could be disk stars. Thus we focus on the colour range $V-I <$ 1.8 and adopt the same halo zone as in the HDF. The GS has 66 stars in this zone. There is a slight chance that this zone might be contaminated by faint white dwarfs in the galactic disk. It can be shown, however, that, if the local density of disk white dwarfs as determined by Jahrei{\\ss} (1987, 1997) and conventional estimates of the vertical scale height of the galactic disk are adopted, one would expect about one white dwarf in this zone of the GS. If we project the local subdwarfs into the cone towards the GS, we obtain the numbers of predicted stars summarized in the last column of Table 3. Again a flattened halo model is a better fit to the data. \\subsection{Constraints on the absolute magnitude of MACHOs.} Like previous authors we have {\\em not} detected the analogues to MACHOs in the solar neighbourhood. This raises the question why. All subdwarfs in our sample appear in the LHS high-proper motion star catalogue. Being halo objects, MACHOs in the solar neighbourhood must have similar proper motions. The fact that they have not been found in high-proper motion surveys allows to set an upper limit to their luminosity (Fuchs \\& Jahrei{\\ss} 1997). The LHS catalogue (Luyten \\& La Bonte 1973, Dawson 1986) is claimed to be about 90 \\% complete for stars down to apparent magnitude $B$ = 21 mag and with proper motions 2.\\hspace{-0.2em}$^{''}$5/a $>\\mu>$ 0.\\hspace{-0.2em}$^{''}$5/a. For stars brighter than $B$ = 10 mag it is claimed that all stars with proper motions $\\mu>$ 0.\\hspace{-0.2em}$^{''}$3/a are contained in Luyten's NLTT catalogue. We assume that the MACHOs are homogeneously distributed around the Sun and have a gaussian velocity distribution, \\begin{equation} dN = \\frac{\\nu_{{\\rm M}_\\odot}}{(2\\pi)^{3/2}\\sigma_{\\rm M}^3} exp - \\frac{1}{2 \\sigma_{\\rm M}^2} (U^2+(V-\\overline{V})^2+W^2) d^3v d^3r\\,\\, , \\end{equation} centered on $\\overline{V}$ = --220 km/s and with a velocity dispersion typical for halo objects, $\\sigma_{\\rm M} = \\overline{V}/\\sqrt{2}$. Integrating now over all radial velocities and the proper motions according to the specifications of the LHS we can determine the radius $r_{\\rm 2}$ of the sphere out of which one would expect say 2 MACHOs in the LHS, \\begin{eqnarray} 2 & = &4.74^2 \\frac{\\nu_{{\\rm M}_\\odot}}{\\sigma_{\\rm M}} \\int_0^{r_2}dr r^4 \\int_{\\mu_l}^{\\mu_ h}d\\mu \\mu \\int_0^{2\\pi}dl \\int_{-\\pi}^{+\\pi}db \\cos b \\nonumber \\\\ & & \\cdot I_0\\left(\\frac{\\overline{V}}{\\sigma_{\\rm M}^2} 4.74 \\mu r \\sqrt{1-\\sin^2l\\cos^2b}\\,\\right) \\nonumber \\\\ & & \\cdot exp -\\frac{1}{2\\sigma_{\\rm M}^2}\\left(\\overline{V}^2(1-\\sin^2l\\cos^2b)+(4.74 \\mu r)^2\\right)\\,\\, , \\end{eqnarray} where $I_0$ denotes the Bessel function. According to Poisson statistics two expected MACHOs would be still consistent with no MACHO actually seen in the LHS survey. In Table 4, we give results of numerical integrations of equation (6) assuming a local density of MACHOs of $\\nu_{\\rm M_{\\odot}}$ = 0.5 $\\cdot$ 0.01 ${\\cal M}_\\odot$ pc$^{-3}$ / 0.5 ${\\cal M}_\\odot$ as determined by the MACHO collaboration. A lower limit of the absolute magnitude of the MACHOs is then given by $M_{\\rm B} = 21 -5\\log{r_2} +5$, because MACHOs more distant than $r_{\\rm 2}$ will not show up in the LHS, since they are fainter than the apparent magnitude limit of the LHS, $B$ = 21 mag. From Table 4 we conclude that MACHOs are fainter than $M_{\\rm B}$ = 21.2 mag. \\begin{table}[htb] \\caption{Estimated absolute magnitudes of MACHOs.} \\vspace{0.4cm} \\begin{center} \\begin{tabular*}{7cm}{|c@{\\extracolsep\\fill}cc|} \\hline \\rule[0mm]{0mm}{4mm} $r_{\\rm N}$ & N & $M_{\\rm B}$\\\\[0.5ex] [pc] & & [mag]\\\\[0.5ex] \\hline 6 & 0.31 & 22.1 \\\\[0.5ex] 9 & 2.3 & 21.2\\\\[0.5ex] 12 & 9.4 & 20.6\\\\[0.5ex] 15 & 27.3 & 20.1\\\\[0.5ex] 18 & 64.0 & 19.7\\\\[0.5ex] \\hline \\end{tabular*} \\end{center} \\end{table} The southern hemisphere ($\\delta < -30^\\circ$) and the galactic belt are not sampled by the LHS survey. Thus one could argue that only the region $b \\ge 30^\\circ$ is properly sampled. This would reduce the numbers in the second column of Table 4 by a factor of 0.25 and shift the lower limit of the absolute magnitude of the MACHOs to $M_{\\rm B}$ = 20.6 mag. If the distribution of MACHOs is irregular and lumpy, the micro-lensing results could mimic a too high mean density of MACHOs. If we assume a local MACHO density of only 10\\% of the value determined by the MACHO collaboration, the lower limit of the absolute magnitude of MACHOs would be $M_{\\rm B}$ = 20.1 mag. If the MACHOs were identified with faint red dwarfs, their masses would have to be less than 0.1 ${\\cal M}_\\odot$ in contradiction to the estimate by the MACHO collaboration. As an alternative very faint white dwarfs have been discussed in the literature as MACHO candidates (Charlot \\& Silk 1995). Graff et al.~1997 discuss various aspects of this scenario and point in particular to the problem that such faint white dwarfs must have been born due to the long cooling times very early in the evolution of the universe. What would be the requirements to detect such faint objects in future proper motion surveys? Assuming, for instance, a limiting magnitude of $B =$ 22.5 mag as for UKST plates, an object with an absolute magnitude of $M_{\\rm B}$ = 21 mag could be detected up to a distance of 20 pc from the Sun. Its typical proper motion would be {3.\\hspace{-0.2em}$^{''}$5/a}. Assuming again a volume density of $10^{-2}$pc$^{-3}$, one would expect 333 objects distributed evenly over the sky." + }, + "9708/astro-ph9708179_arXiv.txt": { + "abstract": "Using recently calculated analytic models for the thermal structure of ultramagnetized neutron stars, we estimate the thermal fluxes from young ($t\\sim 1000$ yr) ultramagnetized ($B \\sim 10^{15}$ G) cooling neutron stars. We find that the pulsed X-ray emission from objects such as 1E 1841-045 and 1E 2259+586 as well as many soft-gamma repeaters can be explained by photon cooling if the neutron star possesses a thin insulating envelope of matter of low atomic weight at densities $\\rho < 10^{7}-10^{8}$ g/cm$^3$. The total mass of this insulating layer is $M \\sim 10^{-11}-10^{-8} M_\\odot$. ", + "introduction": "In recent years, several ``breaking'' (\\cite{Mere95}) or ``anomalous'' (\\cite{vanP95}) x-ray pulsars have been discovered (\\cite{Vasi97b,Corb95}). These objects typically have pulsed X-ray emission with steadily increasing periods $\\sim$ 10 s, X-ray luminosities $\\sim 10^{35}-10^{36}$ erg/s, soft spectra, and no detected companions or accretion disks. Furthermore, they are typically observed through hydrogen column densities $\\sim 10^{22}$ cm$^{-2}$ indicating that they are not common. \\jcite{Vasi97b} describe several of these sources and present observations of 1E 1841-045, whose properties are characteristic of this class of objects. Specifically, \\jcite{Vasi97b} use archival observations of the supernova remnant Kes 73 obtained with the ASCA and ROSAT satellites. They find that the x-ray source in the center of the SNR, 1E 1841-045, had a period of 11.766684 s from the ASCA data taken in 1993 October. The ROSAT data of 1992 March is best fitted with a period of 11.7645 s, yielding a period derivative of ${\\dot P} \\simeq 4.73 \\times 10^{-11} \\rmmat{s s}^{-1}$ and a characteristic spin-down age of 4,000 yr -- close to the estimated age for Kes 73 of 2,000 yr. Using these values and assuming that magnetic dipole radiation dominates the spin-down, they estimate the dipolar field strength of the neutron star to be $\\sim 10^{15}$ G, well above the quantum critical field, $B_{cr} \\approx 4.4\\times 10^{13}$ G. Other anomalous x-ray pulsars (AXPs) generally have small ages and long periods, leading one to derive similar field strengths. With x-ray luminosities $L_X \\sim 10^{35}-10^{36}$ erg/s, AXPs are underluminous relative to accretion powered X-ray pulsars and are generally isolated. For 1E 1841-045, \\jcite{Vasi97b} estimate a spin-down power of $10^{33}$ erg/s which falls short of the observed luminosity. They also argue that although 1E 1841-045 has a period near the equilibrium spin period for a young pulsar with $B \\sim 10^{12}$ G and $L_X \\sim 10^{35}$ erg/s, only an unlikely evolutionary process could spin down the neutron star to this rate within the 2,000 year age of Kes 73. They suggest that 1E 1841-045 may be powered by magnetic field decay in a dipolar field of strength $B \\sim 10^{15}$ G (\\cite{Thom96,Gold92}). In this {\\it Letter}, we propose a natural explanation for the observed X-ray emission from AXPs. Neutron stars with ages $\\sim$ 1,000 yr and magnetic fields $B\\gtrsim 10^{15}$ G have thermal emission in the X-ray-band with total luminosities $\\sim 10^{35}$ erg/s, if their surface layers consist of light-weight material, such as hydrogen and helium. In previous papers, we have developed an analytic model for ultramagnetized neutron star envelopes (\\cite{Heyl97a}) and calculated the emission through iron envelopes and showed how a strong magnetic field affects neutron-star cooling (\\cite{Heyl97b}). Here, we will examine the properties of ultramagnetized hydrogen and helium envelopes and draw parallels with the observed properties of AXPs. ", + "conclusions": "We find that young ($t\\sim 1,000$ yr) neutron stars with strong magnetic fields ($B \\sim 10^{15}$ G) and hydrogen or possibly helium envelopes have photon luminosities similar to those observed from anomalous X-ray pulsars and soft-gamma repeaters in their quiescent state. The total mass of the insulating layer is $10^{-11} - 10^{-8} M_\\odot$. A strongly magnetized neutron star could accrete enough material from the ISM within 1,000 years only if the ISM is sufficiently dense, $n \\gtrsim 10^4 \\rmmat{cm}^{-3}$. However, sufficient material is expected to fall back onto the neutron star surface following the explosion of massive stars." + }, + "9708/astro-ph9708098_arXiv.txt": { + "abstract": "The detection of microlensing events from stars in the Large Magellanic Cloud and in the Galactic bulge raise important constraints on the distribution of dark matter and on galactic structure, although some events may be due to a new type of intrinsic variability. When lenses are relatively close to the sources, we predict that chromatic and spectroscopic effects are likely to appear for a significant fraction of the microlensing events. These effects are due to the differential amplification of the limb and the centre of the stellar disc, and present a systematic dependence with wavelength and time that provide an unambiguous signature of a microlensing event (as opposed to a new type of intrinsic stellar variability). In addition, their measure would provide a direct constraint on stellar atmospheres, allowing a 3-dimensional reconstruction or imaging of its structure, a unique tool to test the current models of stellar atmospheres. ", + "introduction": "Following the pioneering idea by Paczi\\'nsky (1986), the detections of several microlensing events from stars in the Large Magellanic Cloud (Alcock \\etal\\ 1993, 1996; Aubourg \\etal\\ 1993) and of more than 100 from stars in the Galactic bulge (Udalski \\etal\\ 1993, 1994; Alcock \\etal\\ 1995, 1997) have raised important constraints on the structure of the different components of our Galaxy. The associated databases containing millions of light curves are also a goldmine for studies in stellar variability (see Ferlet \\etal\\ 1997, and references therein). Yet doubts may reasonably be cast on the detections of at least some microlensing events. The two EROS candidates are variable stars : EROS 1 is a Be star (Beaulieu \\etal\\ 1995), while EROS 2 is an eclipsing binary (Ansari \\etal\\ 1995). Although spectra of other candidates taken after the event (Della Valle 1994) or during the event itself (Benetti \\etal\\ 1995; Lennon \\etal\\ 1996) show no apparent signs of anomalies, the possibility remains that these events are just detections of a new type of variable star. The detection of ``bumpers'' (Alcock \\etal\\ 1996), a hitherto unknown type of variable star, and the possibility that dwarf novae and cataclismic variables mimick some events (Della Valle \\& Livio 1996) are particularly worrying, even though the distribution of the events in the HR diagram indicates that not all the events can be explained by such phenomena. In addition to continuous and follow-up monitoring, it seems important to find unambiguous proofs that the events detected can indeed be associated with microlensing events. In this context it is interesting to note that the large rate observed towards the bulge may be accounted for if about 50 to 90\\% of the events are due to lenses within the bulge itself (Kiraga \\& Paczy\\'nski 1994), particularly if the bar is oriented close to the line of sight. The same applies if many of the lenses are within the LMC (Wu 1994; Sahu 1994) or in the halo and bulge of M31. If this is the case, a significant number of microlensing events should present characteristic chromatic (Valls-Gabaud 1994; Witt 1995; Bogdanov \\& Cherepaschuk 1995; Gould \\& Welch 1996) and spectroscopic signatures (Valls-Gabaud 1994; Loeb \\& Sasselov 1995), and we present in the {\\it Letter} detailed predictions for both effects. These effects arise from the differential amplification between the limb and the inner regions of the stellar disc, where the emerging intensities are different due to the temperature gradient in the atmosphere. The differential amplification also produces a non-zero degree of polarisation during the events (Simmons \\etal\\ 1995ab), and a spectroscopic line shift (Maoz \\& Gould 1994). ", + "conclusions": "The gravitational imaging of stars is then possible, where not only the brightness distribution on the stellar disc is measured (via the solution of the integral equation in Eq. \\ref{eqn:mag}), but also the source function can be recovered, solving Eq. \\ref{eqn:source}, hence giving a 3-dimensional picture of the stellar atmosphere. This is unprecedented in astrophysics, where only the Sun and very few nearby stars are spatially resolved, either directly (e.g., Gilliland \\& Dupree 1996) or by interferometry (e.g., Baldwin \\etal\\ 1996). Actual techniques for the inversion of Eqs. \\ref{eqn:mag} and \\ref{eqn:source} cannot fit within the scope of this {\\it Letter}, and will be dealt with elsewhere. We simply note here the important implications of stellar imaging. The predictions presented here are based on model atmospheres, in LTE or NLTE, and the expected limb darkening coefficients have seldom been compared with observations others than the solar ones. The reason is that the only way to measure these coefficients (besides the solar case) accurately is using eclipsing binaries, which must have circular orbits, well-detached components, and produce total eclipses. Even then, the derived coefficients correlate with the assumed sizes of the stars (e.g., Tabachnik 1969), so the measure of limb darkening coefficients in microlensing events is a unique opportunity to test stellar atmosphere models. Note however that the inclusion of non-linear terms produces insignificant differences in the resulting magnification, yet even the measure of the linear coefficients is important, for they provide an estimation of the contribution function of the continuum (basically the integrand of Eq. \\ref{eqn:source}). This in turn gives a direct estimation of the temperature profile, something that is known in very few stars (e.g. Matthews \\etal\\ 1996) and that provides unprecedented constraints on the opacity sources. This technique, applied to the lines (ideally the profiles, but the EWs are still extremely important) may lead to the element abundance imaging of the star. In general, atmospheric diagnostics will be possible, by selecting lines sensitive to pressure, or temperature, etc. Unlike Doppler or Zeeman-Doppler imaging, the gravitational imaging is unbiased towards bright stars with large spots, although we note that the presence of spots will considerably complicate the inversion unless the periods are much longer than the time scale of the event. Complications may also arise from blending (Kamionkowsky 1995) and extinction. The ongoing real-time follow-up teams, PLANET (Albrow \\etal\\ 1997) and GMAN (Abe \\etal\\ 1997), could adapt their observing strategies to measure these signatures, by increasing the time sampling, widening the wavelength range covered and obtaining the highest resolution spectra possible. In conclusion, the chromatic and spectroscopic signatures of microlensing events allow us not only to deduce properties of the lenses, but also, and maybe more importantly, to open an entire new window on stellar atmospheres." + }, + "9708/astro-ph9708051_arXiv.txt": { + "abstract": "A detailed study of the absorbing molecular clouds towards the radio source PKS1413+135, at a redshift z$=$0.247, is reported. Physical conditions (density, temperature, filling factors) are derived for the molecular gas, in the frame work of two models: a homogeneous multicomponent model with equal filling factors, and a two--phase medium consisting of dense clumps embedded in a more diffuse component. It is shown that our absorption data are consistent with the presence of a diffuse gas component, dominating the observed opacity, and a dense component, accounting for most of the mass. We also show that without knowledge of the small scale structure of the absorbing molecular gas, we can only derive lower limits to the column density. Given the very narrow absorption spectrum, the size of the overall absorbing cloud along the line of sight must be quite small, of the order of 1\\,pc. The variability of the absorption spectrum has been studied over a time range of more than 2 years. We find that the opacity ratio between two absorbing components in CO has varied by a factor $2.3 \\pm 0.3$. The variations are interpreted as a change of the line of sight due to structural changes in the background source. We discuss what can be derived from molecular absorption line observations concerning the invariance of physical constants, such as the mass of the molecules. We discuss molecular line data from redshifts 0.27--0.89, corresponding to look--back times of 30\\% to 60\\% of the age of the universe. ", + "introduction": "Studies of absorption lines towards high redshift QSOs give detailed information about the interstellar medium (ISM) in distant objects at a linear resolution limited only by the size of the background continuum source. Absorption line observations also have a superior sensitivity compared to the corresponding emission lines and can successfully be used as probes of the composition as well as the physical and chemical conditions in the ISM. Such observations have for the most part been done at optical wavelengths, but the last few years have seen an increase in the number of high redshift absorption line systems observed at radio wavelengths. The neutral atomic gas component has been observed through absorption of the 21\\,cm\\ line of atomic hydrogen. About a dozen 21\\,cm HI absorbers at redshifts $z \\ga 0.2$ have been detected (Carilli 1995), with typical column densities $N(\\hI)=5 \\times 10^{18}\\,(T_{\\rm sp}/f_{\\rm HI})$\\,\\cmsq\\ (e.g. Briggs 1988; Carilli 1995), where $T_{\\rm sp}$ is the spin temperature in K and $f_{\\rm HI}$ is the covering factor of atomic gas across the extent of the background continuum source. Compared to the ionized and atomic parts of the ISM, the molecular gas component is characterized by high densities and low temperatures. Stars are formed directly from the molecular gas, and its physical and chemical status gives information about the stellar formation history as well as giving the initial conditions for the on--going star formation. We have recently detected absorption of molecular rotational transitions in four galaxies at redshifts 0.25--0.89 (Wiklind \\& Combes 1994, 1995, 1996a,b; Combes \\& Wiklind 1996). The background source is in all cases a highly obscured radio--loud QSO, located either in the background or in the galaxy where the absorption takes place. Absorptions in the millimeter wave band provide a velocity resolution more than two orders of magnitude higher than in optical spectroscopy. Since the amount of molecular gas needed to produce a detectable absorption line is less than 1\\,\\mo, the sensitivity of molecular absorption line observations in terms of total mass is $\\sim 10^{12}$ times higher than for the corresponding emission lines (cf. Wiklind \\& Combes 1994). This permits us to derive physical properties of the molecular component of the ISM at high redshifts which would otherwise be impossible. \\medskip In this paper we present new molecular absorption line data on the absorption system at z$=$0.24671 towards PKS1413+135. In Sect.\\,2 we give a general background to the radio source PKS1413+135 and the absorbing galaxy. The observations and data reduction are presented in Sect.\\,3. Our results are given in Sect.\\,4: in particular, we derive the physical properties of the molecular gas and we discuss possible time variations in the absorption lines. In Sect.\\,5 we compare the chemistry of the redshifted molecular gas with Galactic and nearby extragalactic systems and we discuss constraints to variations in physical constants derived from our absorption line data. Finally in Sect.\\,6 we summarize our results. ", + "conclusions": "\\subsection{Small scale structure in the molecular gas} In the Milky Way, several lines of evidence suggest that the molecular ISM is clumpy on all scales from 20\\,AU to 100\\,pc and characterized by a fractal structure (cf. Falgarone et al. 1992). Time variations in Galactic H$_2$CO absorption towards extragalactic point sources (Marscher et al 1993) even suggest structures in the molecular ISM on scales of 10\\,AU. Through numerical simulations Marscher \\& Stone (1994) were able to derive constraints on the fractal structure of molecular clouds from the time variability detections. The mean number of small clumps along the line of sight should be larger than previously thought, i.e. the size spectrum of clumps should be a steeper power--law, constraining the fractal dimension. Detailed multilevel studies of CO and CS near the edge of a nearby molecular cloud complex have revealed that the molecular ISM is confined to small structures (down to 35\\,AU) which are dense ($n_{\\rm H_2} \\sim 10^{4}$\\,cm$^{-2}$), cold ($T_{\\rm K} \\sim 10-15$\\,K) (Falgarone \\& Phillips 1996) and characterized by a fractal structure (Falgarone et al. 1992). In fact, Falgarone \\& Phillips (1996) find no evidence for a pervasive diffuse molecular ISM, although this could be a peculiarity for the specific region studied. Although emission studies are not sensitive to the pervasive diffuse medium, the absorption studies of Lucas \\& Liszt (1996) and Liszt \\& Lucas (1996) show a higher rate of molecular absorption in the Galaxy than was originally expected, and precisely detect the low--opacity, large--covering factor medium. Small scale structures are also present in the atomic ISM. Multi--epoch observations of 21\\,cm absorption against high velocity pulsars have detected opacity variations of the ISM on a range of scales from 5\\,AU to 100\\,AU (Frail et al. 1994). VLBI observations at 21\\,cm in front of 3C radio sources have shown clumps on a scale of 25\\,AU (Diamond et al. 1989, Davis et al. 1996). In the case of 21\\,cm absorption against pulsars, where small--scale opacity structures are detected towards {\\it all} line of sights, the mean opacities are between 0.1 and 2.5, corresponding to N(HI) as low as 10$^{19}-10^{20}$\\,cm$^{-2}$ (Frail et al. 1994). The opacity variations dectected are quite high (half of them have $\\delta\\tau > 0.1$, Frail et al. 1994), so they cannot be accounted for by mild density fluctuations. Clumps with density larger than 10$^5-10^6$ cm$^{-3}$ are implied (Moore \\& Marscher 1995). These clumps could represent 10--20\\% of the total column density. Small dense clumps in the atomic gas has also been inferred in high redshift radio galaxies (van Oijk et al. 1997), where absorption within the Ly--$\\alpha$ emission regions suggest the presence of neutral atomic clumps of size $\\sim$0.03\\,pc. \\medskip Is the molecular gas along the line--of--sight towards PKS1413+135 similar to that of our Milky Way concerning small scale structure? Combining our molecular absorption line data with X--ray data gives three arguments which suggest the presence of a molecular ISM consisting of small dense clumps embedded in a more diffuse medium: \\begin{itemize} \\item The excitation temperatures for three transitions of HCO$^+$ are consistent with a multicomponent molecular gas. If two or more components coexist along the line--of--sight, the assumption of a single gas component underestimates the true column density. \\item The total column density averaged over the extent of the background source and along the line--of--sight (atomic and molecular) is much lower than that derived from the deficiency of low energy X--ray photons. Since the covering factor of gas appears to be high, this implies that we indeed underestimate the total column of gas. \\item Possible time variations in the depth of the molecular absorption lines indirectly suggests the presence of small scale structure. \\end{itemize} The first item has been presented in Sect.\\,4.3 and 4.4. Below we discuss the second and third item in more detail. \\subsubsection{The extinction towards PKS1413+135} The minimum column density of molecular hydrogen towards PKS1413+135, inferred directly from the CO(0$\\rightarrow$1) absorption, is $4 \\times 10^{20}$\\,cm$^{-2}$ (Sect.\\,4.1). Together with the estimated HI column density of $1 \\times 10^{21}$\\,cm$^{-2}$ (where we have assumed a covering factor of unity and a spin temperature of 100\\,K), the total column of hydrogen, $N(H) = 2N(H_2)+N(HI)$, is $\\sim 2 \\times 10^{21}$\\,cm$^{-2}$. This is more than 10 times lower than the column inferred from the deficiency of low energy X--ray photons which indicates $A_v > 30$\\,mag (Stocke et al. 1992). Is this discrepancy caused by an imprecise X--ray measurement, by the particular distribution of the obscuring gas relative to the radio core, or by an underestimate of the molecular and/or atomic column density? McHardy et al. (1994) analyzed ROSAT observations of PKS1413+135 in the 0.4--2.4\\,keV range. Although they did not obtain enough data for a complete spectrum, the very low flux of low--energy X--ray photons measured by ROSAT is consistent with the deficiency observed by the Einstein satellite (Stocke et al. 1992). It is thus not likely that the X--ray results are erroneous. Also, the absence of any lines from both the narrow-- and broad emission line regions, even when viewed in the near--IR (McHardy et al. 1994; Stocke et al. 1992; Perlman et al. 1996), indicates that the extinction is very high. An obscuring dusty torus around the AGN can cause a large extinction. In Cyg\\,A, for instance, the optical extinction, originating in a torus near the center, is estimated to be $\\sim$170\\,mag, or N(H) $= 3.75 \\times 10^{23}$\\,cm$^{-2}$ (Ueno et al. 1994). Although 21\\,cm HI absorption is detected (Conway \\& Blanco 1995), no molecular absorption is seen despite this high column density (Drinkwater et al. 1995). The non--detection of molecular absorption could be due to either a lack of molecular gas as such, or that the gas is dense enough to render the excitation temperature very high, depopulating the lower rotational levels of existing molecules, or that the molecular rotational levels are radiatively coupled with the ambient radiation field (e.g. Maloney et al. 1994). The presence of a dense torus in PKS1413+135, similar to Cyg\\,A, is, however, not likely in view of the absence of a near--infrared excess emission, caused by dust grains heated to high temperatures. This implies that the column density of the molecular gas in PKS1413+135, which should be situated at a relatively large distance from the AGN in order to explain the lack of near--infrared emission, and to be compatible with the very narrow line--width, is severly underestimated, strongly suggesting the presence of a dense clumpy medium as emphasized in previous sections. \\subsubsection{Time variability} Changes in the absorption lines towards background continuum sources can be caused either by motions of small scale structure in the absorbing gas across the angular extent of the radio core, or by spatial changes in the radio core itself. Such variations can give important information about both the size of the radio core and of the scale spectrum of the molecular interstellar medium. The radio source PKS1413+135 consists of a compact radio core and several emission components on a parsec--scale, while no radio emission is seen at scales of kiloparsecs (Perlman et al. 1994, 1996). The parsec--scale components all have a steep--spectrum and do not contribute to the flux at millimeter wavelengths. The core has a spectral index $\\alpha = +1.7$ between $\\lambda$18cm and $\\lambda$3cm (Perlman et al. 1996). The core remains unresolved at an angular resolution of 2.3 mas, corresponding to $\\sim$7\\,pc. It is, however, likely to be significantly smaller. Perlman et al. (1996) derive a lower limit to the size of the core from the variability time--scale of $\\sim$10\\,$\\mu$arcsec, corresponding to $\\sim$0.03\\,pc. The galaxy associated with PKS1413+135 is seen edge--on (McHardy et al. 1994) and we can therefore expect that the transverse velocity is equal to the rotational velocity. The latter is unknown but likely to be $\\sim$250\\,km\\,s$^{-1}$. This corresponds to a transverse shift of 50\\,AU\\,yr$^{-1}$. The time scale for a significant change of the gas properties due to a shift of the obscuring molecular gas is therefore at least 100 years, possibly much longer. The only possibility to explain much shorter variations is the existence of high velocity ($\\ga$ 25 000 km\\,s$^{-1}$) shocks propagating outwards from the radio core. In Table\\,4 we also give the continuum flux relative to that of May 1994 (which was the highest during the extent of these observations). The deepest absorption (largest ratio of the first and second component) occurs when the continuum flux is at maximum, suggesting that the changes are due to structural changes in the background radio source rather than transverse motion of small scale structure in the molecular gas. Knowing the extent of the changes of the radio core during an outburst would give us a limit to the size of the molecular gas structures. Future mm--VLBI observations may provide such data. Also, the change in component 1 is likely to be recurrent with the next outburst, if this takes place in the same region of the radio core. \\begin{figure*} \\psfig{figure=6227-fig12.ps,bbllx=72mm,bblly=25mm,bburx=165mm,bbury=245mm,width=17.5cm,angle=-90} \\caption[]{The HCN/HNC ratio versus the column density of HCO$^+$, HCN and HNC. The trend of decreasing HCN/HNC abundance ratio with increasing column density of HCN indicated in Fig.\\,11 is clearly visible. Designations are as in Fig.\\,10.} \\end{figure*} \\subsection{Molecular cloud properties} The inferences to be made from the analysis of multicomponent molecular gas along the line of sight to PKS1413+135 (Sect.\\,4.3 \\& 4.4) is that it is not possible to derive a correct column density without knowledge about the structure of the molecular ISM. However, if we can assume homogeneity for the physical and chemical conditions over the extent of the background source and along the line of sigth, we can quite confidently derive abundance ratios\\footnote{Abundance is used in synonym with column density while discussing column density {\\em ratios}.}. Furthermore, assuming similarities between the structure of the molecular ISM in PKS1413+135 and the Milky Way, we can compare our column densities with those derived for our Galaxy through similar observing methods (i.e. absorption line measurements). These assumptions are justified by the small amount of molecular gas that actually contributes to the observed absorption. The angular extent of the background continuum source is of the order 10--30\\,$\\mu$arcsec, as implied by time variability (Perlman et al. 1996) and by inferences from mm--wave VLBI of a similar radio source (Lerner et al. 1993). The amount of molecular hydrogen probed by the line of sight to PKS1413+135 is thus only $(0.5-2) \\times 10^{-2}$\\,\\mo, where we have used N(H$_2)=4 \\times 10^{20}$\\,cm$^{-2}$ (cf. Sect.\\,4.1). The size--linewidth relation found for molecular clouds in the Milky Way (Larson 1981; Solomon et al. 1987) implies thay the molecular gas is extended $\\sim$1\\,pc along the line of sight. Assuming that most of the mass is contained in small clumps with $n(H_2) = 10^4$\\,cm$^{-3}$ (Sect.\\,4.4), the total H$_2$ mass is $f_V(0.4-4)$\\,\\mo, where $f_V$ is the volume filling factor of the dense clumps. $f_V$ is likely to be $<<$1. In either case we are dealing with a very small portion of the molecular ISM in the galaxy at z$=$0.247. In Table\\,5 we present abundance ratios for PKS1413+135. The column densities are taken from Table\\,2. The upper and lower limits represents 1$\\sigma$ limits. Below we discuss in some more detail various results that can be derived from these abundance ratios. \\subsubsection{Molecular oxygen O$_2$} One of the more interesting questions in interstellar chemistry is the abundance of O$_2$, which is supposed to bind most of the free oxygen atoms and to be one of the most important coolants (e.g. Goldsmith \\& Langer 1978). This molecule has yet to be observed in the ISM. The redshift of PKS1413+135 shifts the main N(J)$=1_0\\rightarrow 1_1$ line of O$_2$, which is usually not observable from the ground, to an easily accesible window (95.25\\,GHz) However, the rather low average opacity means that our limit to O$_2$/CO$<0.3 \\pm 0.09$ is 20 times higher than the upper limit derived towards B0218+357 at z$=$0.68 (Combes \\& Wiklind 1995). Nevertheless, this limit is of interest since it concerns the ground state and thus is sensitive to excitationally very cold gas. The transitions observed in B0218+357 were the N(J)$=1_2\\rightarrow 3_2$ and the N(J)$=1_1\\rightarrow 3_2$ lines. The $1_1$ level can not be directly populated from the ground state through collisions, but has to be populated through collisional excitation from N(J)$=1_0$ to $1_2$ and $3_2$ (Bergman 1995). The latter energy level corresponds to a temperature $\\sim$23\\,K. Radiative transitions from the ground state to the N(J)$=1_1$ level is allowed and probed by the observations presented here. This transition has now also been observed in B0218+357, giving an improved upper limit to the O$_2$/CO ratio for this system (Combes \\& Wiklind 1997). \\begin{table} \\begin{flushleft} \\caption[]{Abundance ratios} \\small \\begin{tabular}{crclc} \\hline & & & & \\\\ ${{\\rm O}_2 \\over {\\rm CO}}$ & $<$0.3 & $\\pm$ & 0.1 & \\\\ & & & & \\\\ ${{\\rm O}_2 \\over {\\rm HCN}}$ & $<$0.9 & $\\pm$ & 0.3 & $10^3$ \\\\ & & & & \\\\ ${^{12}{\\rm CO} \\over ^{13}{\\rm CO}}$ & $>$74 & $\\pm$ & 53 & \\\\ & & & & \\\\ ${{\\rm H}^{12}{\\rm CO}^+ \\over {\\rm H}^{13}{\\rm CO}^+}$ & $>$111 & $\\pm$ & 6 & \\\\ & & & & \\\\ ${{\\rm HCN} \\over {\\rm CS}}$ & $>$1.1 & $\\pm$ & 0.3 & \\\\ & & & & \\\\ ${{\\rm HCN} \\over {\\rm H}_2{\\rm CO}}$ & $>$1.5 & $\\pm$ & 0.6 & \\\\ & & & & \\\\ ${{\\rm HCN} \\over {\\rm N}_2{\\rm H}^+}$ & $>$10 & $\\pm$ & 3 & \\\\ & & & & \\\\ \\hline & & & & \\\\ ${{\\rm HCO}^+ \\over {\\rm CO}}$ & 1.3 & $\\pm$ & 0.4 & $10^{-3}$ \\\\ & & & & \\\\ ${{\\rm HCN} \\over {\\rm HCO}^+}$ & 0.25 & $\\pm$ & 0.07 & \\\\ & & & & \\\\ ${{\\rm HCN} \\over {\\rm CN}}$ & 0.4 & $\\pm$ & 0.2 & \\\\ & & & & \\\\ ${{\\rm HCN} \\over {\\rm HNC}}$ & 1.4 & $\\pm$ & 1.0 & \\\\ & & & & \\\\ \\hline \\end{tabular} \\end{flushleft} \\end{table} \\subsubsection{Isotopic ratios} The isotopic abundace ratios for CO and HCO$^+$ presented in Table\\,5 suggests relatively low abundances of the $^{13}$C isotopic variants. The large formal error for $^{12}$CO/$^{13}$CO comes from the assumed range of excitation temperatures (cf. Table\\,2). Since the abundance ratio is largely independent of $T_{\\rm x}$, the true uncertainty of this ratio is much smaller. In the ISM of the Milky Way the $^{12}$C/$^{13}$C isotope ratio varies from $\\sim$20 in the Galactic center region to $\\sim$70 in the local ISM (Wilson \\& Matteucci 1992). The correlation between the $^{12}$C/$^{13}$C ratio and that of the substituted isotopic molecules depends on chemical fractionation (enhancing the isotopic variants) and selective photoinization (decreasing the isotopic species due to less self--shielding). In Galactic molecular clouds as well as in nearby galaxies, the final result is an increase in the $^{12}$CO/$^{13}$CO ratio (and likewise for HCO$^+$) compared to the actual $^{12}$C/$^{13}$C ratio. The high lower limits found in PKS1413+135 therefore suggests a low $^{13}$C abundance relative to $^{12}$C. Whereas $^{12}$C is produced in both low and high mass stars, as a primary product of hydrostatic burning, $^{13}$C is only produced through incomplete proton burning in the red giant stage of low and intermediate mass stars (e.g. Wilson \\& Matteucci and references therein). A low $^{13}$C abundance may therefore indicate a young chemistry, where the low and intermediate mass stars have not yet had time to reach the red giant stage. $^{14}$N is produced and ejected into the ISM in much the same way as $^{13}$C. The high lower limit HCO$^+$/N$_2$H$^+ > 41$ found in PKS1413+135 is therefore in aggrement with the interpretation that the low and intermediate mass stars have not yet reached their red giant stage. In PKS1830-211 at z$=$0.89, we found a HCO$^+$/N$_2$H$^+$ ratio of $\\sim$1.4 (Wiklind \\& Combes 1996a). In this molecular absorption line system we also detect several $^{13}$C isotopic species. Although it is not entirely clear whether this is due to an extremely high opacity, PKS1830--211 seems to be more chemically evolved than PKS1413+135, despite the much higher redshift. \\subsubsection{Dark cloud chemistry} The equilibrium abundances in dark clouds (i.e. molecular gas where the main ionization source is cosmic rays) can be divided into a High and Low Ionization Phase (HIP and LIP) (Flower et al. 1994; Le Bourlot et al. 1995; Gerin et al. 1997). The LIP is characterized by a chemistry driven by ion--molecule reactions where proton transfer involving H$_{3}^{+}$ plays a major role. The C/CO ratio is low and most of the gas phase carbon is locked up in CO. This phase has a high abundance of O$_2$ and molecular ions such as HCO$^+$ and N$_2$H$^+$. Other carbon bearing molecules, such as CN, have low abundances. The HIP has a chemistry driven by charge transfer reactions involving H$^+$ and is characterized by a large C/CO ratio. This leads to low abundances of O$_2$ and molecular ions. Abundances of HCN and HNC are only marginally affected by the ionization state of the gas. The abundance ratios presented in Table\\,5 mainly uses HCN as a `reference'. Are our abundance ratios for PKS1413+135 consistent with either of these two phases for dark cloud chemistry? Comparing our abundance ratios with the model results of Le Bourlot et al. (1995) we find that our O$_2$/CO ratio does not constrain the ionization state of the gas, although our limit is of the same order as the expected O$_2$/CO ratio for the low ionization phase. Our limit to the O$_2$/HCN ratio is only compatible with a high ionization phase, while our abundance ratios for HCN/CO, HCN/HNC and the lower limit to HCN/CS only are compatible with a low ionization phase. Hence, the molecular absorption line data for PKS1413+135 does not present a clear--cut case for either of the two ionization phases. Also, the HCO$^+$/CO ratio is too large, by more than an order of magnitude, to be compatible with either the high or low ionization phase. However, as discussed by Hogerheijde et al. (1995), turbulence might influence the chemistry, producing enhanced amounts of HCO$^+$. If this is the case for the high abundance of HCO$^+$ this would indicate a diffuse gas, since the proposed formation pathway is quenched at $n(H_2) \\ga 10^{3}$\\,cm$^{-3}$. Dark cloud chemistry shows bi--stability and can switch from one phase to the other on a short time scale (Le Bourlot et al. 1992, 1995), making it possible to have a mixture of ionization phases over small spatial scales. Improved limits on the O$_2$ and CS abundances may enable us to put limits to the amount of gas along the line of sight to PKS1413+135 that can exist in HIP and LIP states. \\subsubsection{Comparison of column densitites} In Fig\\,10 we plot the column density of HCO$^+$ versus the column densities of HCN, HNC and CS. The plot contains data from Galactic absorption measurements in low density gas (Lucas \\& Liszt 1993, 1994, 1996) and in higher density gas seen towards Sgr\\,B2 (Greaves \\& Nyman 1996), as well as our recent results from absorption line measurements towards Cen A (Wiklind \\& Combes 1997). In addition, we include results from two additional high--z molecular absorption line systems 1504+377 at z$=$0.67 (Wiklind \\& Combes 1996b) and PKS1830--211 at z$=$0.89 (Wiklind \\& Combes 1996a). The dotted line is a one--to--one correspondance between the column densities and not a fit to the data. In Fig.\\,11 we plot the column densities of HNC and CS versus HCN in the same way as in Fig.\\,10. Despite the presence of a considerable scatter, the most striking impression from the figures is the remarkable correlation, over more than 3 orders of magnitude, in column density. This suggests that the chemistry is similar for these simple molecules in environments characterized by very different column densities. The high--z systems do not show any peculiarites compared to the local values, suggesting that the molecular ISM offers similar conditions for star formation at earlier epochs as it does in the present one. The only trend in abundance ratios is found between HCN and HNC (Fig.\\,11), where the abundance of HNC seems to decreases relative to HCN with decreasing HCN abundance. The solid line is a least--square fit to the HNC vs. HCN abundances: \\begin{equation} N(HNC) = 2.3 \\times 10^{-3} \\left[N(HCN)\\right]^{1.16}\\ . \\end{equation} In Fig.\\,12 we plot the HCN/HNC abundance ratio as a function of the HCO$^+$, HCN and HNC column densities. The HCN/HNC ratio decreases with increasing column density in all three cases. Isomers like HCN and HNC have similar chemistry, only differing in a few key reactions which determine the abundance ratio. Through chemical modelling of the HCN and HNC abundances Schilke et al. (1992) found that the HCN/HNC ratio increases with increasing kinetic temperature (cf. Irvine et al. 1987) as well as with increasing molecular hydrogen density, $n_{\\rm H_2}$. The reason for this is primarely the key reactions \\begin{eqnarray} HNC & + & O\\ \\longrightarrow\\ NH\\ + CO \\nonumber \\\\ \\nonumber \\\\ HNC & + & C\\ \\longrightarrow\\ HCN + H \\nonumber \\end{eqnarray} which are effective at relatively high temperatures and preferentially destroy HNC relative to HCN. A large column density does not necessarily imply a large volume density of H$_2$ and, since the abundance of HCN is less dependent on $T_{\\rm k}$ for low H$_2$ abundances (Schilke et al. 1992), the correlations between the HCN/HNC and column densitites seen in Fig.\\,12 is likely to be a temperature effect, affecting primarily the HNC abundance. The lowest column densities being associated with the kinetically warmest gas. However, the dependence of the HCN/HNC abundance ratio on {\\em both} temperature and density makes it less valuable as a temperature probe then previously thought. The high--z systems all tend to have HCN/HNC ratios $\\sim$2, among the lowest of the line--of--sights presented in Figs.\\,10--12. It is unclear whether this is a selection bias towards both kinetically and excitationally cold gas, or an effect of low metallicity, making the HNC depletion reactions less effective. \\begin{table} \\begin{flushleft} \\caption[]{Redshifts and variations of molecular mass} \\scriptsize \\begin{tabular}{lcc} \\hline & & \\\\ \\multicolumn{1}{c}{Transition} & \\multicolumn{1}{c}{z$_a$} & \\multicolumn{1}{c}{$\\left[{\\Delta{\\rm z} \\over {\\rm (1+z)}}\\right]^{a)}$} \\\\ & & \\\\ \\hline & & \\\\ {\\bf PKS1413+135} \\\\ & & \\\\ HI & $0.24671 \\pm 1 \\times 10^{-5}$ & --- \\\\ & & \\\\ CO(0$\\rightarrow$1) & $0.2467091 \\pm 3 \\times 10^{-7}$ & $(+0.7 \\pm 8.0) \\times 10^{-6}$ \\\\ & & \\\\ HCN(1$\\rightarrow$2) & $0.2467112 \\pm 1 \\times 10^{-6}$ & $(-0.9 \\pm 8.1) \\times 10^{-6}$ \\\\ HCN(2$\\rightarrow$3) & $0.2467105 \\pm 3 \\times 10^{-7}$ & $(-0.4 \\pm 8.0) \\times 10^{-6}$ \\\\ & & \\\\ HCO$^+$(1$\\rightarrow$2) & $0.2467102 \\pm 3 \\times 10^{-7}$ & $(-0.2 \\pm 8.0) \\times 10^{-6}$ \\\\ HCO$^+$(2$\\rightarrow$3) & $0.2467097 \\pm 3 \\times 10^{-7}$ & $(-0.2 \\pm 8.0) \\times 10^{-6}$ \\\\ & & \\\\ HNC(1$\\rightarrow$2) & $0.2467114 \\pm 3 \\times 10^{-7}$ & $(-1.1 \\pm 8.0) \\times 10^{-6}$ \\\\ & & \\\\ \\hline & & \\\\ {\\bf B0218+357} \\\\ & & \\\\ HI & $0.68466 \\pm 4 \\times 10^{-5}$ & --- \\\\ & & \\\\ $^{13}$CO(1$\\rightarrow$2) & $0.684693 \\pm 1 \\times 10^{-6}$ & $(+2.0 \\pm 2.4) \\times 10^{-5}$ \\\\ & & \\\\ \\hline & & \\\\ {\\bf B3\\,1504+377} \\\\ & & \\\\ HI & $0.67324 \\pm 1 \\times 10^{-5}$ & --- \\\\ & $0.67343 \\pm 1 \\times 10^{-5}$ & --- \\\\ & & \\\\ HCO$^+$(1$\\rightarrow$2) & $0.673184 \\pm 1 \\times 10^{-6}$ & $(-3.4 \\pm 0.6) \\times 10^{-5}$ \\\\ & $0.673327 \\pm 1 \\times 10^{-6}$ & $(-6.2 \\pm 0.6) \\times 10^{-5}$ \\\\ & & \\\\ \\hline & & \\\\ {\\bf PKS1830} \\\\ & & \\\\ HI & --- & --- \\\\ & & \\\\ HCO$^+$(1$\\rightarrow$2) & $0.8858261 \\pm 3 \\times 10^{-7}$ & \\\\ H$^{13}$CO$^+$(1$\\rightarrow$2) & $0.8858647 \\pm 6 \\times 10^{-7}$ & \\\\ HCO$^+$(2$\\rightarrow$3) & $0.8858298 \\pm 3 \\times 10^{-7}$ & \\\\ N$_2$H$^+$(1$\\rightarrow$2) & $0.8858262 \\pm 4 \\times 10^{-7}$ & \\\\ H$_2$CO(1$_{10}\\rightarrow 2_{11}$) & $0.8858263 \\pm 4 \\times 10^{-7}$ & \\\\ CS(2$\\rightarrow$3) & $0.8858295 \\pm 3 \\times 10^{-7}$ & \\\\ & & \\\\ \\hline \\end{tabular} \\ \\\\ a)\\ $\\Delta z/(1+z) = (z_{HI} - z_{mol})/(1+z_{HI})$ \\\\ HI references: PKS1413+135 (Carilli et al. 1992), B0218+357 (Carilli et al. 1993), B3\\,1504+377 (Carilli et al. 1996). \\\\ Molecular line references: PKS1413+135 (this paper), B0218+357 (Combes \\& Wiklind 1995), B3\\,1507+377 (Wiklind \\& Combes 1996b), PKS1830--211 (Wiklind \\& Combes 1996a,1997). \\end{flushleft} \\end{table} \\begin{figure} \\psfig{figure=6227-fig13.ps,bbllx=72mm,bblly=85mm,bburx=165mm,bbury=200mm,width=8.5cm,angle=-90} \\caption[]{A comparison of the CO(0$\\rightarrow$1) and 21\\,cm HI absorption towards PKS1413+135. In both cases the optical depth is plotted. Notice that Fig.\\,3 in Wiklind \\& Combes 1994 suffers from an erroneous velocity scale, causing the CO and HI lines to be separated by $\\sim$11\\,km\\,s$^{-1}$.} \\end{figure} \\subsection{Invariance of physical constants} A potential diagnostic application of molecular absorption lines at high redshift is to check the invariance of fundamental physical parameters. (cf. Thompson 1975). The energy difference between two adjacent rotational levels in a molecule consisting of two atoms is proportional to $(\\mu r^2)^{-1}$, where $r$ is the bond length and $\\mu$ is the reduced mass of the atoms. A vibrational transition of the same molecule has, to a first order approximation, a $\\mu^{1/2}$ dependence on the reduced mass. Comparison of a vibrational--rotational spectrum from an earlier epoch with a present one can therefore be used to set limits on the changes of the proton and neutron masses. This has been done for molecular hydrogen seen in absorption in a damped Lyman--$\\alpha$ system at z$=$2.81 (Foltz et al. 1988, Varshalovich \\& Potekhin 1996), giving a limit of $2 \\times 10^{-4}$ to $\\Delta m_p/m_p$. For pure rotational transitions one can compare the line frequencies (or similarly, the measured redshifts) with the frequencies of electronic transitions. This gives a measure of the invariance of the electron to proton mass. It is also possible to compare the 21\\,cm HI line with a molecular rotational line. The formers energy difference to first order depending on $m_{e}^2$, again giving a measure of the invariance of the $m_e/m_p$ ratio (see also Tubbs \\& Wolfe 1980 for a comparison of the 21\\,cm line with optical resonance transitions). In an Appendix we express the energy levels of molecular rotational lines in fundamental physical parameters. It is shown that for a differential measurement, i.e. measuring the redshifts of two different molecular species, $\\Delta z/(1+z)$ is directly proportional to $\\Delta \\mu/\\mu$. Strictly speaking, this is only valid when the molecule is approximated as a rigid rotator, but this is a good approximation for low lying $J$ states. In Table\\,6 we list the observed redshifts of the molecular absorption lines in PKS1413+135, together with some lines from B\\,0218+357, B3\\,1504+377 and PKS1830--211. We also list the redshifts derived from 21\\,cm HI absorption when available. The redshifts have been obtained using Gaussian fits. We can compare both the redshifts derived for different molecules and the redshifts of the HI and molecular lines. If changes in the electron mass are neglible, the $\\Delta z/(1+z)$ measures $\\Delta \\mu/\\mu$, to first order, in both cases. The values given in Table\\,6 are the difference between the 21\\,cm HI and molecular redshifts. In Figure\\,13 we show the 21\\,cm HI absorption (Carilli et al. 1992) together with our CO(0$\\rightarrow$1) absorption towards PKS1413+135. It is quite clear that the HI spectra limit the accuracy with which $\\Delta z$ can be measured. Except for B3\\,1504+377, all the values in Table\\,6 are compatible with $\\Delta z = 0$. For the 21\\,cm HI absorption obtained towards B3\\,1504+377 (Carilli et al. 1997), the discrepancy between HI and molecular velocities is of the order of 30\\,km\\,s$^{-1}$, i.e. $\\Delta z \\approx 10^{-4}$, an order of magnitude larger than the HI accuracy. The HI result has been confirmed to an accuracy better than 1\\,km\\,s$^{-1}$ with the Westerbork Synthesis Radio Telescope (Carilli et al 1997, preprint). Hence we reach here the basic limitation of the method, consisting of observing different atoms or molecules. It cannot be expected that all lines peak at the same velocity, since their local association is not perfect, and their excitation might be different. Supporting this, note that the molecular component in absorption in B3\\,1504+377 at $z = 0.67150$ has no HI counterpart. To obtain more statistics, atomic and molecular absorptions in the Milky Way must be considered. Liszt \\& Lucas (1997) compared the absorption profiles of several different molecular lines with that of 21\\,cm HI and found that the velocity of the peak opacity often differs between the molecular and atomic lines. This shows that the absorptions do not arise in the same gas and any comparison is limited by the possible velocity difference of the two ISM components where the absorptions originates. It is also possible to use different transitions from the same molecule to estimate $\\Delta \\mu/\\mu$; While the rotational constant is proportional to $\\mu^{-1}$, the centrifugal distortion coefficient is proportional to $\\mu^{-2}$ (Townes \\& Schawlow 1975). In the Appendix we show that we actually have a dependence on the invariance of the electron mass and the fine structure constant in this case. Comparing the J$=$1$\\rightarrow$2 and J$=$2$\\rightarrow$3 lines of HCO$^+$ in PKS1413+135 we get at the 1$\\sigma$ level: $\\Delta z/(1+z) = (4.0 \\pm 3.4) \\times 10^{-7}$. Similarly for HCN: $\\Delta z/(1+z) = (5.6 \\pm 8.4) \\times 10^{-7}$. Comparing CO(0$\\rightarrow$1) and HCO$^+$(1$\\rightarrow$2) we get: $(-1.1 \\pm 0.3) \\times 10^{-6}$. This last result is marginally different from zero at a 3$\\sigma$ level. However, even when comparing molecular transitions one has to face the possibility that the lines do not come from the same location in phase space. HCO$^+$, for instance, has been shown to exist in regions where the CO molecular is no longer self--shielded (Lucas \\& Liszt 1996), and its absorption profile can therefore be weighted with velocity components not seen in the CO line. A further complication is that HCO$^+$ is an ion and susceptible to magnetic fields present in the ISM. This can give it a drift velocity different from neutral molecular species. It is possible to have absorption lines dominated by different velocity components even when considering transitions from the same molecular species. This is due to the nonlinear dependence of the opacity to the excitation temperature (Sect.\\,4.2). Different transitions can therefore be weigthed differently by gas along the line of sight. To summarize, measurements of redshift differences between molecular absorption lines and 21\\,cm HI are likely to be dominated by velocity differences between the two gas components along the line of sight. This is a severe concern when comparing different molecular species and even different transitions of the same molecule. Taken at face value, the HI accuracy limits the nucleon mass variation to typically $\\Delta \\mu/\\mu \\la 10^{-5}$ at a 3$\\sigma$ level. In one source (B3\\,1504+377) a discrepancy as large as $10^{-4}$ has been observed." + }, + "9708/astro-ph9708267_arXiv.txt": { + "abstract": "The low dispersion (400$\\AA$/mm) spectrum of the optical counterpart of a flat-spectrum radio source 87GB 080315.5+512613, which is one of two possible radio counterparts of 2EG J0809+5117, was obtained recently. The optical counterpart, which is $2.02''$ away from 87GB 080315.5+512613 and $19.3'$ away from 2EG J0809+5117, was identified as a quasar with redshift of 1.14. We noted that Mattox et al. (1997) suggested the other radio counterpart 87GB 080459.4+495915 (OJ 508), which is $87.1'$ away from 2EG J0809+5117, is the more potential identification, though it was previously suggested to be the identification (with low confidence) of another nearby EGRET source 2EG J0807+4849. Our observation suggests that it is quite possible that 87GB 080315.5+512613 is the identification of 2EG J0809+5117 rather than 87GB 080459.4+495915. But we still can not exclude the possibility of 87GB 080459.4+495915 at present. Moreover, in order to determine whether or not 87GB 080315.5+512613 is a blazar type quasar, the optical polarization and variability measures of its optical counterpart are strongly encouraged. ", + "introduction": "The Energetic Gamma Ray Experiment Telescope (EGRET) is the high-energy \\gm telescope on the {\\it Compton Gamma-Ray Observatory (CGRO)}. The telescope covers the energy range from about 30 MeV to over 20 GeV. From April 1991 to October 1994, the all-sky survey program of EGRET has completed three phases observations. Up to now, the published EGRET catalogs include 157 \\gm sources (Fichtel et al. 1994; Thompson et al. 1995; Thompson et al. 1996). Among them, 61 sources have been identified, including 43 AGN with high confidence, 11 AGN with lower confidence, 5 pulsars, one solar flare and Large Magellanic Cloud (LMC). Other 96 sources still remain unidentified. Except LMC, all previously identified EGRET sources with higher galactic latitude (e.g., $|b| >10^o$) are blazar type AGN. These AGN usually have strong, compact, flat-spectrum ($\\alpha \\geq -0.5$, where $S(\\nu)\\propto\\nu^{\\alpha}$) radio emission, strong optical polarization and significant optical variations on short time scales. The blazar class includes objects classified as BL Lacertae type objects, high polarization quasars (HPQ), and optically violently variable (OVV) quasars. Although the \\gm radiation of blazars has not been well understood, some observational properties of blazars are believed to result from a relativistic jet which is directed within $\\sim 10^o$ of the line of sight. By introducing Bayes' theorem to assess the reliability of the identification of EGRET sources with extragalactic radio sources, Mattox et al. (1997) recently demonstrated conclusively that EGRET is detecting the blazar class of AGN. They also indicated possible radio identifications of sources with $\\mid b \\mid >3^o$ in the second EGRET catalog and its supplement. Most of these radio sources have 5GHz radio flux larger than 50 mJy and spectra index larger than -0.5. In order to assure some of these radio sources are more probably the identifications of EGRET sources, optical identifications of these radio sources are necessary. Based on this idea, we are planning a program at Beijing Astronomical Observatory to do the optical spectroscopic studies of the possible optical counterparts of these flat-spectrum radio sources. ", + "conclusions": "We have identified an optical source, which is most probably the counterpart of a flat-spectrum radio source 87GB 080315.5+512613, as a quasar with redshift of 1.14. Although more optical observations are still needed to be done to see whether or not it is a blazar-type quasar, we think it is quite possible that 87GB 080315.5+512613 is the identification of 2EG J0809+5117 rather than another nearby flat-spectrum source 87GB 080459.4+495915 (OJ508). The latter one has been previously suggested to be the lower confidence identification of another EGRET source 2EG J0807+4849 (Thompson et al. 1993; Nolan et al. 1996). But Mattox et al. (1997) recently indicated that another steep spectrum OVV quasar 3C 196 is the more potential identification of 2EG J0807+4849 than 87GB 080459.4+495915 based on the probability analyses. However, in their analyses the priori probability is assumed independent on the EGRET exposure and radio spectra index, which might bring some errors to the estimated posteriori probability. Our observation still can not exclude the possibility that 87GB 080459.4+495915 is the counterpart of 2EG J0809+5117. If the future observations on the optical variability and polarization of the optical counterpart of 87GB 080315.5+512613 confirm it is a blazar-type quasar, we think it will enhance the possibility that 87GB 080315.5+512613 is the identification of 2EG J0809+5117. However, even at this stage we can not state conclusively that it is the true identification because there are still a lot of blazars not detected by EGRET. Therefor, in order to improve the present status of the EGRET source identifications, much more observations and analyses are still expected to be done." + }, + "9708/astro-ph9708117_arXiv.txt": { + "abstract": "In the last two years there have been major advances in our ability to identify and study normal star forming galaxies at high redshifts, when the universe was only 15\\% of its present age. We review the steps which have led to the discovery of a widespread population of objects at $z \\sim 3$ with many of the characteristics which we expect for primeval galaxies, and emphasize in particular the advantages of a colour selection technique which targets the Lyman discontinuity at 912~\\AA. Star forming galaxies at $z = 3$ resemble local starbursts, although they are typically more luminous by more than one order of magnitude. The ultraviolet continuum is dominated by the integrated light of O and early B type stars and shows prominent interstellar absorption lines which are often blueshifted relative to the systemic velocity of the galaxy, indicating highly energetic outflows in the interstellar medium. \\lya\\ emission is generally weak, probably as a result of resonant scattering. The spectral slope of the ultraviolet continuum and the strength of the \\Hb\\ emission line, which we have detected in a few cases with pilot observations in the infrared $K$ band, suggest that some interstellar dust is already present in these young galaxies and that it attenuates their UV luminosities by a factor of $\\sim 3$. The efficiency of our photometric selection technique has allowed us to establish that large scale concentrations of galaxies were already in place at $z = 3$; these structures may be the precursors of today's rich clusters of galaxies, at a time when they were beginning to decouple from the Hubble expansion. In the context of Cold Dark Matter models of structure formation, the galaxies we see must be associated with very large halos, of mass $M \\simgt 10^{12}$\\Ms, in order to have developed such strong clustering at $z = 3$. We conclude by pointing out the need for infrared space observatories, such as the proposed {\\it Next Generation Space Telescope}, for pushing the quest for the origin of galaxies beyond $z = 5$. ", + "introduction": "The quest for the origin of galaxies has been one of the main themes of observational cosmology for many years. A key aspect of this search is the identification of `primeval' galaxies---a somewhat lose concept given that we don't know how galaxies form, but generally taken to mean the progenitors of galaxies like the Milky Way at the time when they first assembled a significant fraction of their mass and began forming their first generations of stars. Until recently the search for primeval galaxies had been a highly frustrating affair with only a handful of objects, mostly discovered serendipitously or through gravitational lensing, as the meagre return for the investment of many months of observing time on large telescopes. This, we now realize, was due less to a lack of adequate instrumentation than to the fact that we did not know how to recognise what we were looking for. A typical deep CCD image obtained at the prime focus of a 4 m telescope shows some 3000 galaxies which are at distances stretching from the vicinity of the Milky Way to the most distant reaches of the universe, corresponding to look-back times of 90\\% of the age of the universe. Thus, such a CCD image in principle contains much of the information required to identify the origin of galaxies and follow their evolution over most of the Hubble time. The challenge until now has been to devise an efficient way to sort this multitude of galaxies according to their redshifts and ages and thereby identify the young counterparts of present day luminous galaxies. The situation has improved dramatically in the last two years and objects which conform closely to our ideas of a primeval galaxy are now being discovered routinely and in large numbers. In this conference contribution we give an account of these recent exciting developments. We first describe the method which has proved to be most profitable for identifying star forming galaxies at high redshifts. We then review some of the most significant properties of this population of objects, deduced from the analysis of their redshift distribution and their spectra. We end with some comments on future prospects, focussing in particular on the infrared spectral region and on the key role which the {\\it Next Generation Space Telescope} will play in pushing this field of research to the earliest epochs. Before proceeding just one word of clarification. When dealing with distant galaxies astronomers normally refer to their redshift, because this is the quantity which is directly measured from the spectra. Redshift, however, is only a proxy for look-back time---we are interested in how far in the past we observe a particular galaxy at redshift $z$. The mapping of redshift to look-back time is not yet as accurate as we would like, although the precision in the measurement of the Hubble constant has improved in the last few years. For reference, Table 1 gives the look-back times corresponding to values of redshift used most often in this review, adopting a Hubble constant $H_0 = 70$~km~s$^{-1}$~Mpc$^{-1}$ and a deceleration parameter for the expansion of the universe $q_0 = 0.1$ (unless otherwise stated, these values are assumed throughout this article). \\begin{table} \\caption{Look-Back Time as a Function of Redshift.}\\label{tbl-1} \\begin{tabular}{rrr} \\\\ $z$ & ~~$T$ (Gyr)\\tablenotemark{a} & ~~$T/T_{\\infty}$\\tablenotemark{a}\\\\ \\tableline \\\\ 0 & 0 & 0 \\\\ 0.5 & 4.6 & 0.39 \\\\ 1 & 6.8 & 0.57 \\\\ 2 & 8.8 & 0.74 \\\\ 3 & 9.8 & 0.83 \\\\ 4 & 10.3 & 0.87 \\\\ 10 & 11.4 & 0.96 \\\\ $\\infty$ & 11.9 & 1.00\\\\ \\end{tabular} \\tablenotetext{a}{$H_0 = 70$~km~s$^{-1}$~Mpc$^{-1}$; $q_0 = 0.1$} \\end{table} ", + "conclusions": "" + }, + "9708/astro-ph9708231_arXiv.txt": { + "abstract": "We discuss the spectral properties of a sample of type-2 Seyfert galaxies based upon the analysis of \\asca\\ data. In this paper we consider the sources for which the X-ray spectra appear to be dominated by the nuclear continuum, transmitted through a large column of absorbing material. We find that both Seyfert-2 galaxies and NELGs show iron K$\\alpha$ line profiles indicative of reprocessing of nuclear X-rays in a face-on accretion disk. Such line profiles are also observed in Seyfert-1 galaxies. This result is contrary to unification models, which would predict the inner regions of Seyfert-2 galaxies to be observed edge-on. This raises some questions as to the orientation of the circumnuclear absorber. If the observed differences between Seyfert type-1 and type-2 galaxies, and NELGs are not due to differences in the orientation of the absorbing material, then we suggest that differences in dust composition and grain size, and in the density of the circumnuclear gas could be of primary importance. ", + "introduction": "\\label{sec:intro} The classification of an Active Galactic Nucleus (AGN) depends upon the wavelength at which one observes. Historically, optical observations yielded the categories of Seyfert types 1 and 2, classifying most of the sources which are now ``famous''. Seyfert-2 galaxies differ from type-1 in that the former show only narrow emission lines in their optical spectra. It was soon realized that some Seyferts showed weak broad components along with the narrow emission lines and consequently the subclasses Seyfert-1.5, 1.8 and 1.9 were introduced to quantify the differences in strength of the broad-line components relative to the narrow lines. Narrow Emission Line Galaxies (NELGs) are bright and variable X-ray sources, discovered in early X-ray sky surveys (Marshall \\etal\\ 1979). The narrow optical emission lines often have weak, broad H$\\alpha$ and P$\\beta$ emission (Ward \\etal\\ 1978, Veron \\etal\\ 1980, Shuder, 1980) making the optical spectra similar to Seyfert-1.9 galaxies. Thus the NELG classification is indicative that the source was discovered in an X-ray survey, but many NELGs are otherwise indistinguishable from Seyfert-1.9 galaxies. Optical spectroscopy and spectropolarimetry, infrared spectroscopy, X-ray spectroscopy and temporal studies plus $\\gamma$-ray spectra are all important in the determination of the fundamental nature of obscured nuclei. Unification models for AGN postulate that large amounts of dense, molecular material exists between the broad-emission-line region (BLR) and the narrow-emission-line region (NLR), in some cases within parsecs of the active nucleus (see Antonucci 1993 for a review of Unified Models for AGN). The simplest geometry consistent with observations is a torus, and consequently it has been suggested that one of the primary factors in Seyfert classification is the orientation of the absorbing torus to our line-of-sight. This hypothesis is consistent with the existence of circumnuclear molecular gas suggested by absorption measurements in a number of wavebands (e.g. Braatz \\etal\\ 1993, Greenhill \\etal 1996). In unified models, sources observed within the opening angle of the torus correspond to those classified optically as type-1 AGN, while sources with lines-of-sight intersecting the torus correspond to type-2 AGN. In the latter case, the nuclear light can nevertheless be observed via scattering or transmission. Antonucci \\& Miller (1985) provided compelling support for this model when they detected broad, Seyfert-1 type emission lines in the polarized optical spectrum of the Seyfert-2 galaxy NGC 1068, and similar results were later obtained for a number of Seyfert-2 galaxies (Miller and Goodrich 1990; Tran, Miller and Kay 1992). Recently, Veilleux \\etal\\ (1997) provided further support for the model with their infrared observations of Pa$\\beta$, Br$\\gamma$ and Br$\\alpha$ lines in many Seyfert 2 galaxies, revealing hidden BLRs which were not always detectable in scattered optical light. Optically-thick torii would be expected to result in collimation of nuclear continuum radiation, and imaging of optical emission lines has shown preferential elongation of some narrow-line regions along the radio axis of the AGN (Haniff, Wilson \\& Ward 1988). These regions must be ionized by a more intense radiation field than is directly observed, again, supporting the Unified Model (e.g. Pogge 1988, Tadhunter and Tsvetanov 1989). Further support came from \\ginga\\ X-ray spectra, which showed large absorbing columns and iron K$\\alpha$ lines of high equivalent width (EW) (e.g. Awaki \\etal\\ 1991). The {\\it ASCA} satellite (Makishima et al. 1996) consists of four co-aligned grazing-incidence X-ray telescopes (XRTs; Serlemitsos et al. 1995). The focal-plane instruments are two solid-state imaging spectrometers (SISs), each consisting of four CCD chips, providing an effective bandpass $\\sim$0.4--10~keV (Burke et al. 1994), and two gas imaging spectrometers (GISs) at the focus of the other two XRTs, providing coverage over $\\sim$0.8--10~keV (Ohashi et al. 1996 and references therein). The sources presented here were systematically analyzed in the same way as the broad-line Seyfert galaxies presented in Nandra \\etal\\ 1997 (hereafter N97). The analysis method is also described in Paper~I. In a previous paper we presented the basic data-analysis results from a sample of {\\it ASCA} observations of type-2 Seyfert galaxies (Turner \\etal\\ 1997a, hereafter Paper I), i.e. those having predominantly narrow optical emission lines. The original sample of 26 observations of 25 narrow-line AGN comprised 17 Seyfert-2 galaxies and 8 NELGs drawn from the {\\it ASCA} public archive. In Paper~I we found the 0.6-10 keV \\asca\\ spectra of a sample of Seyfert-2 galaxies and NELGs to be complex, often containing a heavily-absorbed continuum component, a soft X-ray component and numerous X-ray emission lines. We found the 6-7 keV regime to be dominated by line flux from gas with ionization-state $<$ Fe {\\sc XVI}. Several sites are expected to produce significant X-ray line emission in AGN including the line-of-sight absorber, optically-thick material out of the line-of-sight (both the putative accretion disk and other larger-scale systems such as the torus), ionized (scattering) gas and regions of starburst emission. The absence of a strong 6.4 keV iron line component in starburst galaxies (Ptak \\etal\\ 1997), indicates that the presence of such a line is likely to be an indication of nuclear activity. Iron K$\\alpha$ yields the strongest observed X-ray emission line in Seyfert galaxies, and thus provides an important probe of conditions in the reprocessing material. A discussion of sources dominated by scattered and Compton-reflected X-rays was presented in a second paper (Turner \\etal\\ 1997b, hereafter Paper II). Fig.~1 (repeated from Paper~II, for ease of reference) shows the equivalent width of the iron K-shell line plotted against neutral X-ray absorbing column, $N_H$, for the sample sources. Equivalent widths were measured against the continuum component dominating the 6 - 8 keV range, and based upon the fit to a narrow Gaussian profile (see Paper~I for details). The dot-dashed line in Fig~1 denotes the equivalent width of iron K$\\alpha$ predicted to be produced by transmission through a uniform shell of neutral material (with solar abundances subtending 4$\\pi$ to a continuum source of photon index $\\Gamma=2.0$, Leahy \\& Creighton 1993), where the photon flux $N(E) \\propto E^{-\\Gamma}$. The dashed line shows the equivalent width predicted via Compton-reflection from optically-thick material, as a function of \\nh, assuming that only the power-law is absorbed. In this case the reflection is assumed to be produced from the accretion disk with an equivalent width $\\sim 230$ eV (N97), as typically observed in Seyfert-1 galaxies. Coincidently, 230 eV represents both the maximum equivalent width observed when iron K$\\alpha$ was parameterized as a narrow Gaussian line, and the mean equivalent width assuming a relativistic line profile (N97). Sources can lie significantly above both of these model lines if the direct continuum is hidden but the reprocessed emission is observed, as the line equivalent width is then measured against a suppressed continuum. Consideration of the iron K$\\alpha$, \\verb+[+O{\\sc iii}\\verb+]+ $\\lambda 5007$ line and X-ray variability together suggested that NGC~1068, NGC~4945, NGC~2992, Mrk~3, Mrk~463E and Mrk~273 are dominated by reprocessed X-rays (Paper~II). These sources were denoted ``group C'' (marked with squares on Fig.~1). Sources lying on the ``Leahy and Creighton line'' were denoted ``group A'' (marked as circles on Fig.~1). Thus group A is composed of NGC~1808, NGC~4507, NGC~5252, NGC~6240, ESO~103-G35, IC~5063, NGC~7172 and NGC~7582. (Table~1 shows the group designation of the sources based upon the original sample). Sources with iron K$\\alpha$ equivalent widths lying between that line and the 230 eV line are consistent with Seyfert-1 spectra transmitted through a high absorbing column and are denoted ``group B'' (marked as stars on Fig.~1). Group B is composed of NGC~526A, NGC~2110, MCG-5-23-16 and NGC~7314, which are all NELGs. This classification implies we see the nuclear component directly in group B, and this is supported by the observation of rapid X-ray variability in their flux (Paper~I and Hayashi \\etal\\ 1996) and rapid variability of the iron line profile in NGC~7314 (Yaqoob \\etal\\ 1996). This division into groups A, B and C leaves MCG-01-01-043, NGC~4968, NGC~6251, E~0449, NGC~5135 and NGC~1667 unclassified, i.e. those with the lowest signal-to-noise ratio in the \\asca\\ data. As we will demonstrate, this crude classification of sources yields a useful insight into the relative importance of the regions contributing to the X-ray spectra in several different cases. The distribution of sample sources is shown in Table~2. To summarize, in Paper~II we showed that in group-C sources the iron K$\\alpha$ complex contains significant contributions from neutral and high-ionization species of iron, thus Compton-reflection, hot gas and starburst emission all could make significant contributions to the observed X-ray spectra. Mrk~3 appeared to be the only source which had little contamination by starburst activity and in this case the {\\it ASCA} spectrum below 3 keV is dominated by gas with an X-ray ionization parameter $U_{X}\\sim 5$ (as defined by Netzer 1996) and effective column density $N_{H} \\sim 4 \\times 10^{23} {\\rm cm}^{-2}$. This material is more highly ionized than the zone of material comprising the warm absorber seen in Seyfert-1 galaxies (George \\etal\\ 1997, hereafter G97), but may contain a contribution from shock-heated gas associated with the jet. In this paper the X-ray properties of sources in groups A and B are examined in the context of unified models for AGN. ", + "conclusions": "Examination of the iron K$\\alpha$ line from a sample of NELGs and Seyfert-2 galaxies shows profiles similar to those observed in Seyfert-1 galaxies, and indicative of an origin in an accretion disk orientated face-on. This result is contrary to unification models, which predict the inner regions of Seyfert-2 galaxies to be observed edge-on. The preference for a face-on orientation of both the accretion disk and host galaxy of type-1 and type-2 Seyfert galaxies poses some questions as to the orientation and geometry of gas comprising the circumnuclear absorbing material. If the absorber is composed of clouds then differences in gas density, size of dust grains and dust composition and the distribution of the clouds seem a likely explanation of the observed differences in the X-ray and optical absorption properties of Seyfert-1 galaxies, Seyfert-2 galaxies and NELGs. These results indicate that further refinement is required for unified models for AGN." + }, + "9708/astro-ph9708188_arXiv.txt": { + "abstract": "I discuss population synthesis methods in the context of low-mass compact binaries. Two examples, both constraining the largely unknown strength of orbital angular momentum losses, illustrate the power and problems of such studies. For CVs, the ``standard'' disrupted magnetic braking model predicts that systems below the period gap are on average older than systems above the gap. The corresponding difference in the space velocity dispersion is testable by observations, independent of brightness--dependent selection effects. For LMXBs, the fraction of transients among short--period neutron star systems turns out to be an important diagnostic quantity constraining not only the angular momentum loss rate but also the kick velocity imparted to the neutron star at birth and the common envelope efficiency. Small kicks ($\\la 100$~km/s), low efficiencies and weak magnetic braking are strongly favoured. ", + "introduction": "A binary population synthesis study considers global properties of a certain binary class and tries to relate these to known (or assumed) global properties of the progenitor population, ideally ZAMS binaries, by following evolutionary channels forming binaries of this class. Evolutionary timescales are usually much too long to give rise to directly observable changes of binary parameters in a given system. The only way to test evolutionary theories against observations is to consider a large sample of binaries in a class at different evolutionary states and compare the observed properties of this sample with results from population synthesis calculations. So, not surprisingly, it has become increasingly popular to supplement the presentation of observational results with ``predictions'' from binary synthesis studies, often in ignorance of the significant uncertainties involved. A population synthesis necessarily deals with a large number of parameters, rendering predicted absolute numbers in many cases meaningless. The strength of population synthesis studies lies in the differential comparison between suggested models, allowing one to test their sensitivity to different parameters. A further severe problem is the fact that selection effects may distort the observed picture of a binary class so much that it bears no resemblance to the true intrinsic population. In the following I illustrate the power and problems of population synthesis models with two applications to low--mass compact binaries. These are either cataclysmic variables (CVs) with a white dwarf accretor, or low--mass X-ray binaries (LMXBs), with a neutron star or black hole accretor. The Roche--lobe--filling donor is a low--mass main--sequence or giant branch star. These binaries represent a long--lived, interacting species at the endpoint of a chain of complex progenitor phases, ideal for study with population synthesis methods. ", + "conclusions": "" + }, + "9708/astro-ph9708141_arXiv.txt": { + "abstract": "The galaxy density field as extracted from the \\iras~1.2~Jy redshift survey is compared to the mass density field as reconstructed by the POTENT method from the Mark III catalog of peculiar velocities. The reconstruction is done with Gaussian smoothing of radius $12\\hmpc$, and the comparison is carried out within volumes of effective radii $31-46\\hmpc$, containing $\\approx\\!10-26$ independent samples. Random and systematic errors are estimated from multiple realizations of mock catalogs drawn from a simulation that mimics the observed density field in the local universe. The relationship between the two density fields is found to be consistent with gravitational instability theory in the mildly nonlinear regime and a linear biasing relation between galaxies and mass. We measure $\\betai\\equiv\\Omega^{0.6}/\\bi=0.89\\pm0.12$ within a volume of effective radius 40 \\hmpc, where $\\bi$ is the \\iras\\ galaxy biasing parameter at $12\\hmpc$. This result is only weakly dependent on the comparison volume, suggesting that cosmic scatter is no greater than $\\pm 0.1$. These data are thus consistent with $\\Omega=1$ and $\\bi \\simeq 1$. If $\\bi > 0.75$, as theoretical models of biasing indicate, then $\\Omega > 0.33$ at 95\\% confidence. A comparison with other estimates of $\\betai$ suggests scale-dependence in the biasing relation for \\iras\\ galaxies. ", + "introduction": "\\label{sec:intro} A comparison of the galaxy density field derived from a redshift survey with the mass-density fluctuation field inferred from galaxy peculiar velocity data, allows one to test both gravitational instability theory (GI) and models for the relation between galaxies and mass. If the data are consistent with the assumed model, one can then estimate the value of the cosmological density parameter $\\Omega$. This can be illustrated, for simplicity, by the linear approximation to GI, for which the relation between the mass density fluctuation field $\\delta(\\vx) \\equiv[\\rho({\\vx})-\\bar\\rho]/\\bar\\rho$ and the peculiar velocity field $\\vv (\\vx)$ is \\begin{equation} \\divv = -f(\\Omega)\\, \\delta \\ , \\quad f(\\Omega )\\approx \\Omega^{0.6} \\ , \\quad \\vert \\delta \\vert \\ll 1 \\ , \\label{eq:del=divv} \\end{equation} with distances measured in \\kms\\ (\\ie, the Hubble constant is set to unity). Observations of galaxy peculiar velocities allow us to measure the quantity on the left-hand side, while galaxy redshift surveys provide a measure of the {\\it galaxy\\/} density fluctuation field, $\\delg(\\vx)$. The latter need not be identical to $\\delta (\\vx)$ (\\cf, Bardeen \\etal 1986; Dekel \\& Rees 1987); we adopt here the simplest toy model relating the two fields, linear biasing, \\begin{equation} \\delg({\\vx})=b\\, \\delta({\\vx})\\ , \\label{eq:lin_bias} \\end{equation} where $b$ refers to the galaxies in a specific redshift survey and at a fixed smoothing length. With this model for biasing, \\equ{del=divv} can be rewritten as a relation between the observable quantities, \\begin{equation} {\\bf \\nabla }\\cdot {\\vv}= -\\beta\\, \\delg \\ , \\quad \\beta \\equiv f(\\Omega)/ b \\ . \\label{eq:gi+b} \\end{equation} Thus, in the context of linear GI and linear biasing, the comparison of peculiar velocities and the galaxy distribution enables a measurement of $\\beta$. However, $\\beta$ provides only an indirect estimate of $\\Omega$, because it is a degenerate combination of $\\Omega$ and $b$. There is quite an extensive literature on the comparison of peculiar velocity and redshift survey data to measure $\\beta$ (for reviews see Dekel 1994; Strauss \\& Willick 1995; Strauss 1997b; Dekel 1997ab; Dekel, Burstein \\& White 1997 for a general review of $\\Omega$ measurements). The comparison is not straightforward: the peculiar velocity data are sparse, inhomogeneously distributed, limited to the radial component, and quite noisy, while the redshift data need to be corrected to real space and may trace the mass distribution in a non-trivial way. These difficulties give rise to a variety of statistical biases which depend on the details of the analysis carried out. There are two approaches to this problem, depending on whether the quantities that are actually compared are velocities or densities. Integrating both sides of \\equ{gi+b} yields a predicted velocity field given measurements of the galaxy density field (equation~\\ref{eq:v=int_over_x_lin} below). Comparison of these predictions to observed radial peculiar velocities allows a determination of $\\beta$ (Strauss 1989; Kaiser \\etal 1991; Hudson 1994). One can make the comparison more sophisticated by smoothing the two velocity fields before comparing them (Davis, Nusser \\& Willick 1997), or by using the predicted velocity field to minimize the scatter in the distance indicator relation from which the peculiar velocities are measured in the first place (Strauss 1989; Roth 1994; Schlegel 1995; Shaya, Peebles, \\& Tully 1995; Willick \\etal 1997b). For \\iras\\ galaxies, these velocity comparisons yield values of $\\betai$ ranging from $0.49 \\pm 0.07$ (Willick \\etal 1997b) to $0.86 \\pm 0.14$ (Kaiser \\etal 1991), depending on the details of the analysis, the smoothing scale, and the data used. Davis \\etal (1997) have claimed inconsistencies between the peculiar velocity and redshift survey data in the context of GI and linear biasing. Alternatively, one can use the POTENT method (\\S~\\ref{sec:potent}) to recover the density fluctuation field from the peculiar velocity data and use \\equ{gi+b}, or its nonlinear extension (see the discussion in \\S~\\ref{sec:potent_veldel}), to compare to the galaxy density field. Dekel \\etal (1993, hereafter PI93) carried out such an analysis, using the \\iras\\ 1.936 Jy redshift survey (Strauss \\etal 1992b), and the Mark II compilation of peculiar velocities (Burstein 1989). The advantages of the differential form include the facts that the direct comparison of densities is {\\it local\\/} (whereas the velocity field is sensitive to the mass distribution in a large volume, perhaps even outside the sampled volume), it is independent of reference frame, and it allows direct control over the smoothing of the fields. Monte-Carlo tests showed that the POTENT density field of PI93 was biased in a variety of ways, forcing the use of an elaborate maximum likelihood technique to quantify the consistency between data and model, and to measure $\\betai$. PI93 did find consistency, and concluded that $\\betai=1.28^{+0.75}_{-0.59}$ at 95\\% confidence. A similar comparison of POTENT densities with optically selected galaxies by Hudson \\etal (1995) yielded an acceptable fit, with $\\beta_{opt} = 0.74 \\pm 0.13$ (1$\\sigma$). The current paper, like PI93, follows the general approach of a {\\it density} comparison, but is a significant step forward due to a number of improvements in the quantity and quality of the data and the methods of analysis. In particular: \\begin{itemize} \\item The current analysis is based on the Mark III catalog of peculiar velocities (Willick \\etal 1995, 1996, 1997a). With $\\sim\\!3400$ galaxies, it is more than three times bigger than the Mark II data set, and has better space coverage. The data sets composing the Mark III catalog have been treated with more care, especially in self-consistently calibrating the Tully-Fisher (TF) relations, grouping, and correcting for inhomogeneous Malmquist bias (\\S~\\ref{sec:potent}). \\item The POTENT method for deriving the density fluctuation field $\\delp$ from peculiar velocity data has been much improved since PI93 (\\S~\\ref{sec:potent}; Dekel \\etal 1997). \\item The current analysis uses the \\iras\\ 1.2 Jy redshift survey (Fisher \\etal 1995), containing twice as many galaxies as in the 1.936 Jy survey. \\item The method for deriving a uniform galaxy density field $\\deli$ from the redshift data has been improved in several ways since PI93 (\\S~\\ref{sec:iras}). \\item The availability of much more realistic simulations of both the \\iras\\ and Mark III datasets (\\S~\\ref{sec:eval_mock}; Kolatt \\etal 1996) allows much better error analysis in both the peculiar velocity and density fields, and in the comparison. \\end{itemize} We use these simulations to assess biases in our determination of $\\betai$. Unlike PI93, these biases turn out to be negligible, allowing us to sidestep the rather elaborate likelihood analysis of that paper. Indeed, we will use the simulations themselves as a guide to whether our data are statistically consistent with the null hypothesis of GI and linear biasing. We also use them to quantify the statistical errors in our final results. The current analysis compares the density fields smoothed with a Gaussian window of radius $12\\hmpc$, where the fluctuations are of order unity and therefore require a mildly nonlinear treatment. The POTENT analysis computes the density field $\\delp$, a generalization of $-f(\\Omega)^{-1} \\divv$ that is a nonlinear function of $\\Omega$ and the spatial partial derivatives of the observed $\\vv(\\vx)$ (equation~\\ref{eq:delc+}). The \\iras\\ reconstruction, in turn, yields a mildly nonlinear galaxy density field, $\\deli$, that is a weak function of $\\Omega$ and $\\bi$ --- only via the corrections from redshift space to real space (equations~[\\ref{eq:cz-r}] and [\\ref{eq:v=int_over_x}]). Equation \\eq{gi+b} is thus replaced by \\begin{equation} \\deli = \\bi\\, \\delp \\ . % \\label{eq:deli=delp} \\end{equation} The $\\Omega$ dependence of equation~(\\ref{eq:gi+b}) is already included in $\\delp$. The density fields are recovered from the data for assumed values of $\\Omega$ and $\\bi$. Then, $\\bi$ is determined by \\equ{deli=delp}, and $\\betai$ is quoted. The analysis is carried out for several initial values of $\\Omega$ and $\\bi$ to confirm the robustness of the estimate of $\\betai$. The present work is less ambitious than PI93 in one respect. An attempt was made in PI93 to use the nonlinear effects to break the degeneracy between $\\bi$ and $\\Omega$, but with only limited success. The resulting constraints on each parameter separately were quite weak, indicating that the nonlinear effects associated with these data are not sufficient for this purpose. Even with the new data, the nonlinear effects are comparable to the errors that accompany the reconstructions. Furthermore, if we are to consider nonlinear gravity, we should also allow nonlinear extensions to the biasing relation, \\equ{lin_bias}, but in practice, it is not clear how to distinguish nonlinear gravity effects from nonlinear biasing effects. We therefore limit ourselves in this paper to determining the degenerate combination $\\betai$. The outline of this paper is as follows: In \\S~\\ref{sec:potent} we discuss the POTENT reconstruction of the mass-density field from peculiar velocities, and in \\S~\\ref{sec:iras} the reconstruction of the galaxy density field from the \\iras\\ redshift survey. In \\S~\\ref{sec:eval} we use mock catalogs to evaluate the random and systematic errors in the reconstructed density fields, and to minimize them if possible. In \\S~\\ref{sec:method} we describe our method of comparison of the two fields, and estimate the systematic and random errors in the measurement of $\\betai$. In \\S~\\ref{sec:results} we perform the comparison of the real data, evaluate goodness of fit, and determine the value of $\\betai$. In \\S~\\ref{sec:conc} we conclude with the implications of our results and compare them with other recent determinations of $\\betai$. ", + "conclusions": "\\label{sec:conc} We have compared the density of mass and light as reconstructed from the Mark III catalog of peculiar velocities and the \\iras\\ 1.2 Jy redshift survey, with Gaussian smoothing of $12\\hmpc$ and within volumes of effective radii $31-46\\hmpc$. Our two main conclusions are: \\begin{itemize} \\item{} The data are consistent with gravitational instability theory in the mildly nonlinear regime and a linear biasing relation between galaxies and mass on these large scales. \\item{} The value of the corresponding $\\beta$ parameter at this smoothing scale is $\\betai=0.89 \\pm 0.12$. The relative robustness of this value to changes in the comparison volume suggests that the cosmic scatter is no greater than $\\pm 0.1$. % \\end{itemize} This result is consistent with the simplest cosmological model of an Einstein -- de Sitter universe ($\\Omega=1$, $\\Lambda=0$) with unbiased \\iras\\ galaxies, and it also permits somewhat lower values of $\\Omega$ and $\\bi$. However, values as low as $\\Omega\\sim 0.3$ would require large-scale anti-biasing of $\\bi<0.75$ (with 95\\% confidence) --- a phenomenon that is not easily reproduced in theoretical simulations (e.g., Cen \\& Ostriker 1992; Evrard, Summers, \\& Davis 1994). Although our analysis uses an extension of the Zel'dovich approximation, we have not tried to directly measure nonlinear effects in these data to break the degeneracy between $\\Omega$ and $\\bi$: the effects are weak (PI93), and are of the same order as possible nonlinearities in the biasing relation. An interesting area for future work is to see if we can put an {\\it upper\\/} limit on the degree of nonlinearity; if we can, this puts a lower limit on $\\bi$, and therefore a lower limit on $\\Omega$. As explained in \\S~\\ref{sec:intro}, the current analysis is a significantly improved version of the PI93 analysis done with the earlier Mark II and \\iras\\ 1.936 Jy data, and of the Hudson \\etal (1995) comparison to optical galaxies. The result of PI93 was $\\betai=1.28^{+0.75}_{-0.59}$ at 95\\% confidence. The new result is lower, but only by about $1.3\\sigma$. The current analysis is superior in many respects. The improvements in the data include denser sampling of a larger volume, a careful procedure of selection, calibration and merging of the TF data sets, and a better correction for Malmquist bias. The methods of reconstruction were considerably improved, using new techniques and with extensive testing and error analysis using realistic mock catalogs drawn from simulations. The optimization of the comparison method using the mock catalogs led to an unbiased estimate of $\\beta$ with well-defined statistical errors. The quantity $\\beta$ has been measured using a variety of techniques from observational data sets (as discussed in the references in \\S~\\ref{sec:intro}); we here discuss other estimates of $\\betai$ from the Mark III and \\iras\\ 1.2 Jy data. Estimates of $\\betai$ from {\\it redshift distortions\\/} in redshift surveys are almost all within 2$\\sigma$ of our current result (Peacock \\& Dodds 1994, $\\betai = 1.0\\pm0.2$; Fisher \\etal 1994b, $\\betai = 0.45^{+0.3}_{-0.2}$; Fisher, Sharf, \\& Lahav 1994c, $\\betai = 1.0 \\pm 0.3$; Cole, Fisher, \\& Weinberg 1995, $\\betai = 0.5 \\pm 0.15$; Hamilton 1995, $\\betai = 0.7 \\pm 0.2$; Heavens \\& Taylor 1995, $\\betai = 1.1 \\pm 0.3$; Fisher \\& Nusser 1996, $\\betai = 0.6 \\pm 0.2$). The scatter between these results is mostly due to differences in method of analysis, but cosmic scatter, nonlinear effects, and complications in the biasing scheme probably also play a role. Our current analysis is a comparison at the density level (a $\\delta-\\delta$ comparison). Alternative analyses that performed the comparison at the {\\it velocity\\/} level ($v-v$ comparisons) typically yielded somewhat lower values for $\\betai$, ranging from $0.5$ to $0.9$: (Strauss 1989, $\\betai \\simeq 0.8$; Kaiser \\etal 1991, $\\betai = 0.86^{+0.2}_{-0.15}$; Roth 1994, $\\betai = 0.6\\pm0.2$; Nusser \\& Davis 1994, $\\betai = 0.6\\pm 0.2$; Davis \\etal 1997, $\\betai = 0.6\\pm 0.2$; Willick \\etal 1997b, $\\betai = 0.49 \\pm 0.07$). We focus here on the two most recent and sophisticated analyses of the $v-v$ type, termed ITF and VELMOD. The ITF analysis (Davis \\etal 1997) is a mode-by-mode comparison of velocity-field models, based on expansion in spherical harmonics and radial Bessel functions, that were fitted in redshift space to a processed version of the inverse TF data from the Mark III catalog and the \\iras\\ 1.2 Jy redshift survey. The ITF analysis finds a poor fit between the velocity fields and the model, in the form of a distance-dependent dipole at large distances, in apparent contrast with the good fit obtained here. One possible explanation for this difference is an error in one of the large-scale velocity fields, or both, that does not propagate to their local derivatives, the density fields. A differential offset in the zero point of the TF relation between data sets of the Mark III catalog, which generates an artificial coherent bulk motion, is one such possibility. An inaccurate definition of the Local Group frame in the \\iras\\ reconstruction, or missing density data beyond the sampled volume, are other possible sources of error in the velocities. Another source of error is the somewhat arbitrary adjustments that had to be made to the Mark III data sets when combining them into one system for the ITF analysis (cf., the discussion in Davis \\etal 1997). If one ignores the poor fit, the ITF analysis yields the formal estimate $\\betai = 0.6 \\pm 0.2$, $\\sim 1 \\sigma$ lower than our current result. The local nature of the density-density comparison, and the resulting better goodness of fit, argue that it may provide a more reliable estimate of $\\betai$. The VELMOD analysis (Willick \\etal 1997b) is a high-resolution $v-v$ comparison at a small smoothing scale of $3\\hmpc$, using about one quarter of the Mark III data within a volume of radius $< 30\\hmpc$, and applying a sophisticated likelihood analysis in the linear regime. The comparison revealed a good fit and an estimate of $\\betai=0.49\\pm0.07$. We can see several possible reasons for the difference in the estimates of $\\betai$. It is possible that the difference originates from systematic errors that were somehow overlooked despite the successful testing using the same mock catalogs. For example, the VELMOD analysis assumes pure linear theory, and may thus suffer from nonlinear effects that were not taken into account. In fact, it was never fully understood how the linear analysis of VELMOD did well at G3 smoothing where nonlinear effects are expected to be important. One suspicion is that the mock simulation is too smooth on small scales. Another source for the difference may be the partial data used in the VELMOD comparison, where the comparison is in fact dominated by data within $\\sim 20\\hmpc$. The difference in $\\betai$ can thus partly reflect cosmic scatter in $\\bi$ or in the local effective $\\Omega$. Table~1 shows that $\\betai$ increases by an insignificant $1\\sigma$ from the smallest to the largest comparison volume; the VELMOD analysis also found no statistically significant growth of $\\betai$ with scale. This allows us to put an upper limit on the cosmic scatter of order 0.1. Finally, a possible explanation for the difference in $\\betai$ is {\\it scale dependence} of the biasing relation between Gaussian smoothings of $3$ and $12\\hmpc$, which would be associated with non-linear biasing. Some support for such a trend is found by the SIMPOT analysis (Nusser \\& Dekel 1997), which fits a parametric model of velocity and density fields and $\\beta$ simultaneously to the peculiar velocity and redshift data (a $v-\\delta$ comparison). This analysis yields $\\betai =1.0\\pm0.15$ for G12 smoothing, and lower values of $\\betai$ for smaller smoothing scales, in qualitative agreement with the comparison of the results of the current paper and of VELMOD. Current theoretical simulations indicate possible scale dependence in the biasing relation between scales of one to a few megaparsecs (\\eg, Kauffman \\etal 1997; cf., Mo \\& White 1996), but it remains to be seen whether such a trend can arise on scales of $6$ to $12\\hmpc$. The biasing scheme, which could be nontrivial in several ways (see Dekel \\& Lahav 1997), is clearly a bottle-neck in the effort to measure the cosmological parameter $\\Omega$ via $\\beta$ --- the biasing is inevitably involved whenever galaxy-density data are used. The largest source of error in this analysis lies in the peculiar velocity data, and thus improved peculiar velocity datasets, with better control of systematic errors, denser sampling, and more complete sky coverage, will be of great importance for this work (cf., the reviews of Strauss 1997a; Giovanelli 1997). It will be interesting to make the comparison of the POTENT maps with the optical redshift data of Santiago \\etal (1995, 1996), although a proper calculation of the errors in the optical density field will be somewhat more difficult. Finally, more work lies ahead in testing the robustness of our results to the assumed value of $\\Omega$ (see the discussion in \\S~\\ref{sec:results_beta}), and carrying out the analysis at other smoothing lengths to look for nonlinear effects." + }, + "9708/astro-ph9708069_arXiv.txt": { + "abstract": "We examine the possibility that a significant component of the energy density of the universe has an equation-of-state different from that of matter, radiation or cosmological constant ($\\Lambda$). An example is a cosmic scalar field evolving in a potential, but our treatment is more general. Including this component alters cosmic evolution in a way that fits current observations well. Unlike $\\Lambda$, it evolves dynamically and develops fluctuations, leaving a distinctive imprint on the microwave background anisotropy and mass power spectrum. ", + "introduction": " ", + "conclusions": "" + }, + "9708/astro-ph9708155_arXiv.txt": { + "abstract": "The Space Interferometry Mission (SIM), with its launch date planned for 2005, has as its goal astrometry with $ \\sim 1 ~ \\mu $ arcsecond accuracy for stars as faint as 20th mag. If the SIM lives to expectations it can be used to measure astrometric displacements in the light centroid caused by gravitational microlensing in the events detected photometrically from the ground. The effect is typically $ \\sim 0.1 $ mas, i.e. two orders of magnitude larger than planned SIM's accuracy. Therefore, it will be possible to determine the mass, the distance, and the proper motion of almost any MACHO capable of inducing a photometric microlensing event towards the galactic bulge or the Magellanic Clouds, even though no light from the MACHO has to be detected. For strong microlensing events in which the source is photometrically resolved, like the recent MACHO 95-30 event, SIM's astrometry combined with accurate ground based photometry will allow the determination of the angular stellar radii, and therefore the effective temperature of the source. The effective astrometric cross sections for gravitational lensing by nearby high proper motion stars and brown dwarfs are $ \\sim (1'')^2 $ and the effective time scales are $ \\sim 1 $ year. SIM will provide the only practical way to measure masses of single nearby objects with $ \\sim 1\\% $ accuracy. The times of lensing events can be predicted years ahead of time. ", + "introduction": "The searches for gravitational microlensing events in the galactic bulge and in the Magellanic Clouds have matured, and frequent alerts of new events are provided in real time (cf. Paczy\\'nski 1996a for a review). Unfortunately, currently there is no way to firmly determine the distance to a lensing object and to measure its mass because a number of physical parameters combine into a single observable quantity: the time scale $ t_0 $. In some cases the degeneracy may be partly broken, as reviewed by Gould (1996). Also, at least one suggestion was made how to remove the degeneracy altogether in some rare cases (extreme microlensing: Gould 1997). H${\\rm \\o}$g, Novikov \\& Polnarev (1995), Miyamoto \\& Yoshi (1995), and Walker (1995) pointed out that accurate astrometry could permit the determination of the distance and the mass of a MACHO. It seems that the Space Interferometry Mission (SIM) is the first specific instrument which will have the capability adequate for this task. According to the mission specification, it will have an angular resolution of $ \\sim 10 $ mas, and positional accuracy down to $ \\sim 1 $ micro-arcsecond, for stars as faint as 20 mag. A complete description of the technical details as well as the scientific program were published by Allen, Shao, \\& Peterson (1997), and can also be found on the World Wide Web at: \\centerline{\\it http://huey/jpl.nasa.gov/sim/ } \\noindent Some of the topics covered in this letter (section 3) are presented in much more details in the paper written by three members of the SIM team: Boden, Shao, \\& Van Buren (1997). Recently, a somewhat related aspects of gravitational lensing astrometry were considered by Miralda-Escud\\'e (1996) and Paczy\\'nski (1996b), who in particular pointed out that astrometric cross section is much larger than photometric cross section. Therefore, the masses of single high proper motion stars and brown dwarfs can be determined accurately by measuring the astrometric lensing effects of the distant stars. The purpose of this paper is to present in one place various suggestions made in the past for the application of $ \\sim 1 $ micro-arcsecond astrometry to determine the masses of stars, brown dwarfs, and MACHOs, whatever the MACHOs might be. In addition, a new suggestion to use such astrometry to measure stellar radii and effective temperatures is also presented. ", + "conclusions": "A $ \\sim 1 ~ {\\rm \\mu s} $ astrometry of microlensing events can be used to determine the lens masses and the source radii (and hence their effective temperatures), but it requires Target of Opportunity mode of operation. In order to be useful the Space Interferometry Mission (Allen, Shao,\\& Peterson 1997, Boden, Shao \\& Van Buren 1997) would have to respond to ground based alerts within a few weeks. The only way to measure masses of single objects with an accuracy of $ \\sim 1\\% $ is with gravitational lensing. Lensing events caused by nearby high proper motion stars passing within $ 1''-10'' $ of distant background stars, can be predicted many years into the future, making it possible to prepare the SIM's observing program ahead of time, with no need for the TOO operation." + }, + "9708/astro-ph9708225_arXiv.txt": { + "abstract": "In recent years, the large angle COBE--DMR data have been used to place constraints on the size and shape of certain topologically compact models of the universe. Here we show that this approach does not work for generic compact models. In particular, we show that compact hyperbolic models do not suffer the same loss of large angle power seen in flat or spherical models. This follows from applying a topological theorem to show that generic hyperbolic three manifolds support long wavelength fluctuations, and by taking into account the dominant role played by the integrated Sachs-Wolfe effect in a hyperbolic universe. ", + "introduction": "The most obvious observational signature of a multiply connected universe would be repeated or ``ghost'' images of familiar objects such as galaxies or rich clusters\\cite{ghost}. However, searches for ghost images are hampered by evolution of the objects; our ability to recognise objects when viewed from different directions; and the difficulty in determining the distances to objects. Despite these problems, the consensus seems to be that there is no evidence for ghost images out to redshifts of $z\\sim 0.4$ -- the current depth of wide-field redshift surveys. It is interesting to note that this lack of ghost images is exactly what one expects for typical small compact hyperbolic models. According to Thurston\\cite{bill}, the expectation value for the length of the shortest closed geodesics in a typical small hyperbolic universe is roughly $(0.5 \\rightarrow 1.0)R_{0}$, where $R_{0}=H_0^{-1}/\\sqrt{1-\\Omega_0}$ is the comoving curvature radius. Here $H_0$ is the Hubble constant and $\\Omega_0$ is the matter density in units of the critical density. The first copies of the Milky Way galaxy or Coma cluster would not be seen before a conformal lookback time of $\\eta\\simeq 0.5 \\rightarrow 1.0$. Converting this to redshift space via the relation \\begin{equation} 1+z=\\frac{2(\\Omega^{-1}_0-1)}{\\cosh (\\eta _0-\\eta )-1} \\, , \\end{equation} where $\\eta_0={\\rm arccosh}(2/\\Omega_0 -1)$ is the present conformal time, we find that the first ghost images will be at a redshift of $z\\simeq 0.9 \\rightarrow 2.9$ in a universe with $\\Omega_0=0.3$. If the universe has $\\Omega_0$ closer to unity, the first ghost images will be even more distant. These numbers suggest that direct searches for ghost images of astrophysical objects will be unable to tell if we live in a compact hyperbolic universe. A more promising approach is to look for topological lensing of the last scattering surface by studying fluctuations in the cosmic microwave background radiation\\cite{css1,css2}. ", + "conclusions": "A hyperbolic drum produces a rich and complex sound. A compact hyperbolic universe is likewise infinitely more complex than its spherical or euclidean counterparts. The simple methods used to constrain flat models do not work when space is negatively curved. The eigenmodes in a compact hyperbolic space can only be calculated using sophisticated methods developed to treat quantum chaos. Moreover, hyperbolic models do not suffer the simple long wavelength cut-off used to exclude toroidal models. In addition to the issue of what fluctuations are supported on the last scattering surface, there is also the issue of what exactly it was that COBE measured. In a compact hyperbolic universe the curvature radius provides a natural length scale, $R_0=H_0^{-1}/\\sqrt{1-\\Omega_0}$. The curvature radius sets the length scale where we might hope to find the first evidence that we live in a multiply connected universe. The curvature radius also sets the angular scale beyond which fluctuations in the cosmic microwave background radiation no longer originate from the last scattering surface. This confluence of physical scales is very unfortunate for COBE since it means that the ISW effect takes over just when things get interesting. Fortunately the next generation of CMB satellites will be able to probe much smaller angular scales, so the ISW effect will not obscure their view of the large scale topology of the universe. The search for multi-connectedness in our universe is not over. It has barely begun." + }, + "9708/astro-ph9708013_arXiv.txt": { + "abstract": "We report the first detection of interstellar hydrogen fluoride. Using the Long Wavelength Spectrometer (LWS) of the Infrared Space Observatory (ISO), we have detected the 121.6973 $\\mu$m $J=2-1$ line of HF in absorption toward the far-infrared continuum source Sagittarius B2. The detection is statistically significant at the $13 \\,\\sigma$ level. On the basis of our model for the excitation of HF in Sgr B2, the observed line equivalent width of $1.0$ nm implies a hydrogen fluoride abundance of $\\sim 3 \\times 10^{-10}$ relative to H$_2$. If the elemental abundance of fluorine in Sgr B2 is the same as that in the solar system, then HF accounts for $\\sim 2 \\%$ of the total number of fluorine nuclei. We expect hydrogen fluoride to be the dominant reservoir of gas-phase fluorine in Sgr B2, because it is formed rapidly in exothermic reactions of atomic fluorine with either water or molecular hydrogen; thus the measured HF abundance suggests a substantial depletion of fluorine onto dust grains. Similar conclusions regarding depletion have previously been reached for the case of chlorine in dense interstellar clouds. We also find evidence at a lower level of statistical significance ($\\sim 5\\,\\sigma$) for an emission feature at the expected position of the $4_{32}-4_{23}$ 121.7219~$\\mu$m line of water. The emission line equivalent width of $0.5$~nm for the water feature is consistent with the water abundance of $5 \\times 10^{-6}$ relative to H$_2$ that has been inferred previously from observations of the hot core of Sgr B2. ", + "introduction": "Molecules are a ubiquitous component of the dense interstellar medium. To date, about one hundred distinct species have been detected in the interstellar gas (Ohishi 1997); they range from simple diatomic molecules to complex species containing as many as thirteen atoms. The wide variety of interstellar molecules demonstrates the thermodynamic tendency of most elements to form molecules under the conditions present in the dense interstellar medium. Although most of the interstellar molecules that have been detected previously contain elements\\footnote{viz.\\ hydrogen, oxygen, carbon, nitrogen, sulphur and silicon} of cosmic abundance greater than $10^{-5}$, a few molecules containing elements of lower abundance have also been detected. For example, chlorine, with a solar system abundance of only 2$\\times 10^{-7}$ relative to hydrogen (Anders \\& Grevesse 1989), was detected more than a decade ago in the form of hydrogen chloride (Blake, Keene \\& Phillips 1985). Recent estimates (Neufeld \\& Green 1994; Schilke, Phillips \\& Wang 1995; Zmuidzinas et al.\\ 1995a) of the HCl abundances implied by observations of the HCl $J=1-0$ line toward Sgr B2 (Zmuidzinas et al.\\ 1995a) and the Orion Molecular Cloud 1 (Blake et al.\\ 1985; Schilke et al.\\ 1995) have yielded results for $n({\\rm HCl})/n({\\rm H}_2)$ in the range 0.3 -- 2 $\\times 10^{-9}$. If the elemental abundance of chlorine in these sources is the same as that in the solar system, then HCl accounts for only $0.1 - 0.7\\%$ of the total number of chlorine nuclei. Theoretical models for the chemistry of chlorine-bearing molecules (Schilke et al.\\ 1995) predict that HCl will account for $\\sim 30\\%$ of gas-phase chlorine. Thus the observed abundance of HCl can only be understood if the chlorine depletion in the dense ISM is large. The required depletion factors of $\\simgt 10^2$ greatly exceed the values inferred for the {\\it diffuse} ISM from UV absorption line studies (e.g.\\ Harris, Gry, \\& Bromage 1984 ). Prior to the observations reported in this Letter, HCl was the only known {\\it interstellar} molecule containing a halogen element\\footnote {Note, however, that NaCl, KCl, AlCl, and AlF have been detected in the circumstellar envelope of IRC+10216 (Cernicharo \\& Gu\\'elin 1987; Ziurys, Apponi \\& Phillips 1994), and that HF lines have been widely observed in the spectra of cool stars (e.g. Jorissen, Smith \\& Lambert 1992).}. Motivated by the earlier observations of HCl, by the fact that the solar system abundance of fluorine lies only a factor of 6 below that of chlorine (Anders \\& Grevesse 1989), and by the large H--F bond strength which suggests that hydrogen fluoride is likely to be a major reservoir of gas-phase fluorine, we have undertaken a search for interstellar hydrogen fluoride toward the strong far-infrared continuum source Sgr B2. Such a search would provide a unique probe of the chemistry of interstellar fluorine and of the fluorine depletion in the dense interstellar medium. Because of its large rotational constant -- the largest of any diatomic molecule other than molecular hydrogen or HeH$^+$ -- HF possesses a rotational spectrum that lies entirely shortward of the atmospheric windows within which most molecules show a rotational spectrum. Observations of HF rotational transitions are therefore possible only from airplane or space-based observatories. We have made use of the Long Wavelength Spectrometer (LWS; Clegg et al.\\ 1996) on board the Infrared Space Observatory (ISO; Kessler et al.\\ 1996) to search for hydrogen fluoride. The $J=1-0$ transition of HF lies longward of the wavelength range to which the photoconductive detectors used in LWS are sensitive, so we have carried out observations of the $J=2-1$ line at 121.6973 $\\mu$m. These observations are described in \\S 2 below. Our results are presented in \\S 3, and discussed in \\S 4. The detection of hydrogen fluoride reported here marks the first discovery of an {\\it interstellar} molecule containing the element fluorine and the first time that a new astrophysical molecule has been identified by means of observations in the far-infrared (30 -- 300 $\\mu$m) spectral region. ", + "conclusions": "" + }, + "9708/nucl-th9708056_arXiv.txt": { + "abstract": "We study the pionic decay of a possible dibaryon $d^{\\prime}\\to N+N+\\pi$ in the microscopic quark shell model. The initial $d^{\\prime}$ dibaryon wave function (J$^P$=0$^-$, T=0) consists of one 1$\\hbar\\omega$ six-quark shell-model $s^5p[51]_X$ configuration. The most important final six-quark configurations $s^6[6]_X$, $s^4p^2[42]_X$ and $(s^4p^2-s^52s)[6]_X$ are properly projected onto the NN channel. The final state NN interaction is investigated by means of two phase-equivalent - but off-shell different - potential models. We demonstrate that the decay width $\\Gamma_{\\rm{d'}}$ depends strongly on the short-range behavior of the NN wave function. In addition, the width $\\Gamma_{\\rm{d'}}$ is very sensitive to the mass and size of the $d^{\\prime}$ dibaryon. For dibaryon masses slightly above the experimentally suggested value $M_{\\rm{d'}}$=2.065 GeV, we obtain a pionic decay width of $\\Gamma_{\\rm{d'}}\\approx$ 0.18--0.32 MeV close to the experimental value $\\Gamma_{\\rm{d'}}\\approx$ 0.5 MeV. ", + "introduction": "During the last decade much attention has been devoted to theoretical and experimental investigations of the pionic double charge exchange ($\\pi$DCX) process on nuclei. Because this reaction $\\pi^++(A,Z)\\to (A,Z+2)+\\pi^-$ involves (at least) two nucleons in the nucleus, the $\\pi$DCX cross section depends sensitively on short-range NN-correlations in nuclei. Therefore, it provides a good testing ground for the nucleon-nucleon interaction at short range. Experiments on different nuclear targets have unambiguously confirmed the existence of a narrow resonance-like structure in the $\\pi$DCX cross-section at small incident pion energies $T_{\\pi}\\approx50$ MeV \\cite{bil91}. The position of this peak turned out to be largely independent of the studied nucleus. The height and width of this peak could not be explained by standard calculations based on the two-step process \\cite{kam} ($n+n+\\pi^+\\to n+p+\\pi^0\\to p+p+\\pi^-$). So far, these data could only be explained with the assumption of a non-nucleonic reaction mechanism \\cite{bil91,mart91} proceeding via an intermediate dibaryon resonance, henceforth called $d^{\\prime}$. The quantum numbers of the $d^{\\prime}$ dibaryon candidate were determined as J$^P$=0$^-$, T=0, and its free mass and hadronic decay width were suggested to be $M_{\\rm{d'}}$=2.065 GeV and $\\Gamma_{\\rm{d'}}\\simeq$ 0.5 MeV\\footnote{This value is uncertain by a factor of two \\cite{Clem95}.}. More than a decade ago Mulders {\\sl et al.\\ } \\cite{muld} predicted a dibaryon resonance with quantum numbers $J^P$=$0^-$, T=0 and a mass $M\\approx 2100$ MeV within the MIT bag model. Recently, this dibaryon candidate has been investigated in a series of works \\cite{wag95,buch95,iton} within the T\\\"{u}bingen chiral constituent quark model. These works emphasize the crucial role of the confinement mechanism for the existence of the $d^{\\prime}$. The quantum numbers J$^P$=0$^-$, T=0 of the $d^{\\prime}$ resonance prevent the decay into two nucleons and the only allowed hadronic decay channel of the $d^{\\prime}$ is the three-body decay into a $\\pi NN$ system with S-waves in each particle-pair \\cite{bil91,mart91}. Because the $d^{\\prime}$ mass $M_{d^{\\prime}}$ is only $\\approx 50$ MeV above the $\\pi NN$ threshold, the $d^{\\prime}$ decay width $\\Gamma_{d^{\\prime}}$ should be anomalously small owing to a very small phase volume of three-particle final states. We recall that the currently available experimental evidence of dibaryon excitations in nuclei is very limited \\cite{seth}. This is due to very large N-N decay widths of most dibaryon resonances, which renders them undetectable on the background of other hadronic processes at intermediate energy. At present, the experimental evidence for narrow dibaryons is reduced to a single candidate, the $d^{\\prime}(2065)$. In contrast to the deuteron, which consists of two on the average widely separated nucleons, there are indications \\cite{wag95,buch95}, that the $d^{\\prime}$ is a rather pure, compound six-quark system. Therefore, the dynamics of its hadronic decay into the $\\pi$NN system should be sensitive to the overlap region of the two outgoing nucleons; a situation that is ideal for understanding the role of quark degrees of freedom in the short-range nucleon-nucleon interaction (see e.g.\\ Ref.\\ \\cite{gold95} and references therein). Starting from this point (for alternative approaches see in Refs.\\ \\cite{kam,gar,sim}) we consider the $d^{\\prime}$ decay as a (quark) shell-model transition from one six-quark configuration to another one by emitting a pion. The quark line diagram of the decay is sketched in Figure \\ref{figure:obufig1}. The calculation of the transition matrix elements $d^{\\prime}\\to N+N+\\pi$ is similar to the calculation of $\\Delta$-isobar-decay matrix elements $\\Delta\\to N+\\pi$ (spin and isospin flip of a quark). In the case of the $d^{\\prime}$ decay only the initial dibaryon state is a definite six-quark configuration (the lowest shell-model configuration with quantum numbers J$^P$=0$^-$, T=0), whereas the final state consists of a continuum of NN-states which have to be projected onto a basis of six-quark configurations with quantum numbers J$^P$=0$^+$, T=1 of the NN $^1S_0$ wave. The main difficulty in comparing the calculated width $\\Gamma_{d\\prime}$ with experimental data is its sharp dependence on the energy gap between $M_{d^{\\prime}}$ and the $\\pi$NN threshold. A reliable result on $\\Gamma_{d^{\\prime}}$ can be obtained only if the exact mass $M_{d^{\\prime}}$ in vacuum is known (e.g.\\ from electroexcitation of the $d^{\\prime}$ on the deuteron at large momentum transfers \\cite{schep}). At present, we have only indirect data in the nuclear medium \\cite{bil91}. Due to the absence of vacuum data, we investigate the problem of the $d^{\\prime}$ decay width starting from theoretical quark-model results \\cite{wag95,buch95} for $M_{d^{\\prime}}$ and the hadronic $d^{\\prime}$ size parameter $b_6$. Our first calculation for $\\Gamma_{d^{\\prime}}$ was published in Ref.\\ \\cite{iton}. The aim of the present work is to improve mainly on three important effects which were neglected in Ref.\\ \\cite{iton}: a) antisymmetrization of the final NN-state on the quark level taking into account the effect of quark exchange between the two nucleons at short range, b) insertion of a complete basis of final six-quark states including besides the non-excited $s^6$ shell-model state all Pauli-allowed excited configurations $s^4p^2$ and $s^52s$, which have a non-vanishing overlap with the final NN-state and can be populated via the emission of the pion from the initial $d^{\\prime}$ dibaryon, and c) inclusion of the final state interaction (f.s.i.) for the two-nucleon system. ", + "conclusions": "\\noindent The total hadronic decay width of the possible $d^{\\prime}$ dibaryon $\\Gamma_{d^{\\prime}}$ contains three partial widths \\begin{equation} \\Gamma_{d^{\\prime}}= \\Gamma_{\\pi^-pp} + \\Gamma_{\\pi^0pn} + \\Gamma_{\\pi^+nn} = 3 \\Gamma_{\\pi^-pp} \\label{gtot} \\end{equation} which are equal to each other $\\Gamma_{\\pi^-pp}=\\Gamma_{\\pi^0pn}=\\Gamma_{\\pi^+nn}$, when we neglect isospin breaking effects. The partial $\\pi^-pp$ decay width $\\Gamma_{\\pi^-pp}$ is defined by the standard expression \\cite{iton} \\begin{eqnarray} \\Gamma_{\\pi^-pp} &=& 2\\pi \\int d^3q\\int d^3k \\, \\delta\\!\\left( M_{d^{\\prime}}-2M_N-\\frac{k^2}{4M_N}-\\frac{q^2}{M_N}- \\sqrt{m_{\\pi}^2+k^2} \\right) \\nonumber \\\\ &\\times & \\Bigl| \\langle\\Psi_{NN}({\\bf q}),\\pi^- \\, | \\, \\hat {\\cal{O}}_{\\pi q}({\\bf k}) \\, | \\, d^{\\prime} \\rangle \\Bigl|^2 \\; , \\label{gampi} \\end{eqnarray} where ${\\bf q}=\\frac{{\\bf q}_1-{\\bf q}_2}{2}$ is the relative momentum of the two final protons and ${\\bf k}$ is the momentum of emitted pion in the c.m.\\ of the $d^{\\prime}$ dibaryon. The $\\delta$-function conserves the energy in the decay, while the integration over the momentum conserving $\\delta^{(3)}({\\bf q}_1+{\\bf q}_2+{\\bf k})$ has already been exploited \\cite{iton} in Eq.\\ (\\ref{gampi}). The integration over three-particle phase space leads to the following result for the partial $d^{\\prime}\\rightarrow \\pi^-pp$ decay width \\begin{eqnarray} \\Gamma_{\\pi^-pp} &=& \\frac{2^5 10^2\\sqrt{6}}{3^8\\sqrt{\\pi}} \\frac{f_{\\pi q}^2}{4\\pi}\\frac{1}{m_{\\pi}^2} \\left( \\frac{2b_6/b_N}{1+b_6^2/ b_N^2} \\right)^{\\!12} \\nonumber \\\\ &\\times & \\int\\limits_0^{q_{max}} \\frac{2M_N(k_0b_6)^5}{2M_N+\\sqrt{m_{\\pi}^2+k_0^2}} \\, \\exp \\left ({-\\frac{5}{12}k_0^2b_6^2} \\right ) \\left[ I_{NN}^{(0)}(q)+\\sqrt{\\frac{2}{27}} \\left( 1-\\frac{k_0^2b_6^2}{24} \\right) I_{NN}^{(2)}(q) \\right]^2 q^2\\,dq \\; . \\label{result} \\end{eqnarray} Here, energy conservation relates the pion momentum $k_0$ to the NN relative momentum $q$ via $$ k_0(q) = \\left\\{ 4M_N \\left[ \\left( M_{d^{\\prime}} - \\frac{q^2}{M_N} \\right) - \\sqrt{ \\left( M_{d^{\\prime}} - \\frac{q^2}{M_N} \\right)^2 - \\left( M_{d^{\\prime}} -2M_N - \\frac{q^2}{M_N} \\right)^2+m_{\\pi}^2 } \\right] \\right\\}^{1/2} $$ and for $q_{max} = \\sqrt{M_N(M_{d^{\\prime}}-2M_N-m_{\\pi})}$, all available decay energy is converted to kinetic energy in the relative NN-system, and none to the pion $E_\\pi = m_\\pi, k_0=0$. The calculated decay widths are shown in Table \\ref{table2}, where we have introduced the abbreviations p.w., T and U. Here, p.w.\\ refers to a calculation employing a plane-wave final N-N state (\\ref{pw}), while T and U refer to calculations using the Tabakin \\cite{tab} (T) and Ueda et al.\\ \\cite{ued} (U) separable NN potentials for the final state interaction. In parentheses we give the results obtained in the approximation of using only one intermediate six-quark configuration $s^6$ (n=0). With the exception of the results for the Ueda NN-potential for sets 1, 2 and 4 (for which the $d^{\\prime}$ mass is 400 -- 650 MeV above the $\\pi NN$ threshold), the inclusion of all Pauli principle allowed intermediate 2$\\hbar\\omega$ shell-model configurations tends to increase the decay width by some 20 -- 30 $\\%$. The largest effect is obtained for $d^{\\prime}$ masses rather close to threshold, exemplarily shown for sets 3 and 5. It can be seen from Table \\ref{table2} and Fig.\\ \\ref{figure:obufig3}, that the pionic decay width of the $d^{\\prime}$ is very sensitive to the dibaryon mass $M_{d^{\\prime}}$, which determines the available phase space of the three-body $\\pi NN$ decay. The sensitivity grows dramatically near the $\\pi NN$ threshold (2016 MeV). If we extrapolate the results of Table \\ref{table2} to the experimental value of $M_{d^{\\prime}}$ = 2065 MeV, we obtain a very strong reduction of $\\Gamma_{d^{\\prime}}$ as compared with the quite realistic variants (sets 3 and 5) in Table \\ref{table2}: $$ \\begin{array}{cccl} \\Gamma_{d^{\\prime}}^{p.w.} = 0.032\\,{\\rm MeV}, & \\Gamma_{d^{\\prime}}^{T} = 0.046\\,{\\rm MeV}, & \\Gamma_{d^{\\prime}}^{U} = 0.083\\,{\\rm MeV}, & \\; \\mbox{if} \\; b_N \\mbox{=0.595 fm and}\\; b_6 \\mbox{=0.95 fm} \\\\ \\Gamma_{d^{\\prime}}^{p.w.} = 0.018\\,{\\rm MeV}, & \\Gamma_{d^{\\prime}}^{T} = 0.045\\,{\\rm MeV}, & \\Gamma_{d^{\\prime}}^{U} = 0.040\\,{\\rm MeV}, & \\; \\mbox{if} \\; b_N \\mbox{=0.6 fm and}\\; b_6 \\mbox{=1.24 fm}\\; . \\end{array} $$ This strong dependence of $\\Gamma_{d^{\\prime}}$ on the value of $M_{d^{\\prime}}$ is a consequence of the high power of $(k_0b_6)^5$ in the integrand of Eq.\\ (\\ref{result}). The origin of this $k_0^5$ behavior (compared with a $k_0^3$ behavior in case of the $\\Delta$-isobar decay) comes, as explained above, from the necessity to excite (or de-excite) a p-wave quark for the production of a pion. Note that for small $q_{max}$ (when $M_{d^{\\prime}}$ is close to the $\\pi$NN threshold) the function $k_0(q)$ is linear in the factor $\\sqrt{q^2_{max}-q^2}$ and can be written as $k_0(q)\\approx q_{max}\\sqrt{{4m_\\pi} (1-q^2/q_{max}^2)/M_{d^{\\prime}}}$. Therefore, for small $q_{max}$ the integral in Eq.\\ (\\ref{result}) behaves as $q_{max}^8$. The second high-power factor in Eq.\\ (\\ref{result}) is the scale factor $$ \\left( \\frac{2 b_6/b_N}{1+b_6^2/b_N^2} \\right)^{\\! 12}, $$ which depends sensitively on the ratio $b_6/b_N$. However, this sensitivity is considerably reduced by the factor $b_6^5$ in the integrand. The product $$b_6^5 \\left( \\frac{2 b_6/b_N}{1+b_6^2/b_N^2} \\right)^{\\! 12}$$ is a quite smooth function of $b_N/b_6$. For $b_N$=0.6 fm this product varies from 0.078 fm$^5$ to 0.158 fm$^5$, if $b_6$ varies from 0.6 fm to 1.24 fm. For small $q_{max}$, f.s.i.\\ make an important contribution to the $d^{\\prime}$ decay width because of the large scattering length in the $^1S_0$ wave $a_s=-23.7$ fm. The f.s.i. enhances the decay width for example by about $85\\%$ for set 5 in Table \\ref{table2}. At the experimental mass $M_{d^{\\prime}}$=2065 MeV, the hadronic decay width is more than doubled by the final state interaction. It is interesting that in the case of the Tabakin model with a nodal NN wave function at short range, the contribution from f.s.i.\\ is smaller than for the Ueda model and can even decrease the width compared to the plane wave result (cf.\\ set 4). This is a direct consequence of an approximate orthogonality of the nodal wave function of the Tabakin model to the projection of the intermediate $s^6$ configuration (i.e.\\ the h.o.\\ function $\\varphi_{00}$) of Eq.\\ (\\ref{appr}) onto the NN channel. This can easily be seen from Fig.\\ \\ref{figure:obufig2}, where both wave functions are shown. The approximate orthogonality of the functions $\\varphi_{00}$ and $\\Phi_{NN}^{Tabakin}$ in the integrand of Eq. (29) reduces considerably the overlap factor $I_{NN}^{0}(q)$, which gives the dominant contribution to the $d^{\\prime}$ decay width (see values in parenthesis in Table \\ref{table2}). As it can be seen in Fig.\\ \\ref{figure:obufig3}, the disagreement between the Tabakin and Ueda models grows with increasing dibaryon mass $M_{d^{\\prime}}$ (the influence of the large scattering length $a_s$, which is common for both models, becomes negligible compared to the effect of the larger phase space). For sets 1, 2 and 4 in Table \\ref{table2}, the Tabakin model leads again to values of $\\Gamma_{d^{\\prime}}$, which are even smaller than $\\Gamma_{d^{\\prime}}$ in the plane wave approximation neglecting f.s.i." + }, + "9708/astro-ph9708048_arXiv.txt": { + "abstract": "We have measured fluxes or flux limits for 31 of the 79 cluster candidates in the Palomar Distant Cluster Survey (PDCS) using archival ROSAT/PSPC pointed observations. Our X-ray survey reaches a flux limit of $\\simeq 3 \\times 10^{-14}$ erg s$^{-1}$ cm$^{-2}$ (0.4 - 2.0 keV), which corresponds to luminosities of $L_x\\simeq 5 \\times 10^{43}$ erg s$^{-1}$ (${\\rm H_o}$ = 50 km s$^{-1}$ Mpc$^{-1}$, ${\\rm q_o}$ = $\\frac{1}{2}$), if we assume the PDCS estimated redshifts. Of the 31 cluster candidates, we detect six at a signal-to-noise greater than three. We estimate that $2.9^{+3.3}_{-1.4}$ (90\\% confidence limits) of these six detections are a result of X-ray emission from objects unrelated to the PDCS cluster candidates. The net surface density of X-ray emitting cluster candidates in our survey, $1.71^{+0.91}_{-2.19}$ clusters deg$^{-2}$, agrees with that of other, X-ray selected, surveys. It is possible, given the large error on our contamination rate, that we have not detected X-ray emission from any of our observed PDCS cluster candidates. We find no statistically significant difference between the X-ray luminosities of PDCS cluster candidates and those of Abell clusters of similar optical richness. This suggests that the PDCS contains objects at high redshift similar to the low redshift clusters in the Abell catalogs. We show that the PDCS cluster candidates are not bright X-ray sources, the average luminosity of the six detected candidates is only $\\bar{L_x}=0.9\\times10^{44}$ erg s$^{-1}$ (0.4-2.0 keV). This finding is in agreement with previous X-ray studies of high redshift, optically selected, rich clusters of galaxies. ", + "introduction": "One of the focal points of modern observational cosmology is the study of how structure evolves. Clusters of galaxies provide an excellent probe of such evolution because they represent the largest gravitationally bound objects and can be identified over a large range of redshifts. X-ray surveys for clusters of galaxies have produced catalogs of physically massive objects, but most high redshift X-ray surveys are limited to small ($\\sim 30$) samples (\\cite{henry92} 1992; \\cite{castander95} 1995; \\cite{collins97} 1997). By contrast, the availability of large format CCDs and sophisticated cluster finding algorithms means that it is now possible to produce optically selected catalogs, such as the Palomar Distant Cluster Survey (PDCS, \\cite{postman96} 1996), which contain $\\sim 100$ distant clusters with a large range of richnesses. In this paper we seek to compare the distant clusters found in automated optical surveys with those found with more traditional methods by studying the X-ray properties of the PDCS and by comparing the PDCS to lower redshift cluster samples. Previous studies (\\cite{henry82} 1982; \\cite{bower94} 1994; \\cite{castander94} 1994; \\cite{nichol94} 1994; \\cite{sok96} ) have demonstrated that optically rich clusters at high redshift are often X-ray faint. Many of the clusters examined have velocity dispersions consistent with massive, bound, systems (\\cite{bower97} 1997). The low luminosity of these systems has been cited as evidence for X-ray evolution in the cluster population (\\cite{bower94} 1994). However, this interpretation may be premature, since the high redshift optical samples studied to date are extremely small and heterogeneously selected. With our survey we are able to improve on previous X-ray studies of high redshift, optically selected, clusters. Our sample has the combined advantages of size (31 clusters in all), an objective cluster catalog (the PDCS) and random selection for the X-ray observations. We are also able to take advantage of two recent, low redshift ($z<0.2$) X-ray surveys (\\cite{bh93} 1993; \\cite{burg94} 1994) of Abell clusters (\\cite{abell58} 1958; \\cite{aco89} 1989) which are both statistically complete and contain a large ($\\sim 200$) number of objects. Our sample should provide a fair representation of the X-ray properties of the PDCS as a whole and allow us to make direct comparisons with the \\cite{bh93} (1993) and \\cite{burg94} (1994) surveys. Our survey will help us understand what sorts of galaxy overdensities were selected as clusters by the PDCS and enable us to test for evolution in the cluster population. In Sec. 2, we review the properties of the Palomar Distant Cluster Survey. Section 3 describes our analysis of the fourteen archival ROSAT/PSPC pointings that contain our X-ray data. We describe possible sources of contamination and compare with other published work in Sec. 4. We discuss our results in Sec. 5, and make conclusions in Sec. 6. Throughout this paper we assume ${\\rm H_o}$ = 50 km s$^{-1}$ Mpc$^{-1}$ and ${\\rm q_o}$ = $\\frac{1}{2}$. ", + "conclusions": "We have searched for X-ray emission around 31 distant, optically selected, rich clusters of galaxies taken from the PDCS sample of \\cite{postman96} (1996). We have found six possible coincidences between PDCS clusters and X-ray sources detected in ROSAT/PSPC pointings. Our statistical analysis, and explicit searches for alternative optical or radio counterparts to the X-ray sources, are consistent with 3 of these 6 being detections of contaminating, non-cluster, sources. Our survey demonstrates that PDCS cluster candidates are not strong ($L_x\\gtrsim 5 \\times 10^{43}$ erg s$^{-1}$) X-ray emitters. This observation is consistent with the results of previous studies of distant, optically selected, clusters (\\cite{bower94} 1994; \\cite{castander94} 1994; \\cite{nichol94} 1994; \\cite{sok96}) and highlights the possible existence of a population of optically rich clusters which are X-ray faint. The average luminosity of the six detections in our survey, $\\bar{L_x} = 0.9\\times 10^{44}$ ergs s$^{-1}$, and the estimated surface density of X-ray emitting clusters in the PDCS fields, 1.71 deg$^{-2}$, are consistent with the results of previous surveys for distant, X-ray selected, clusters (\\cite{castander95} 1995; \\cite{rosati95}; \\cite{scharf97}; \\cite{collins97} 1997). Taking both the upper limits and the detections, we find no statistically significant correlation between optical richness and X-ray luminosity in our sample, in contrast to the strong correlation ($\\geq 99.95\\%$ confidence level) measured previously for Abell clusters (\\cite{abram83}; \\cite{kowalski84} 1984; \\cite{bh93} 1993). This is not surprising given the large uncertainties in the redshifts of the PDCS. We find that our sample is in agreement with that of \\cite{bh93} (1993) when we compare optical to X-ray luminosity ratios, which is a redshift independent test. Our data are, therefore, consistent with the hypothesis that the PDCS and the Abell catalog sample the same section of cluster population, albeit at different redshifts. Furthermore, we find no evidence for evolution in the X-ray versus optical properties of rich clusters out to a redshift of $z=0.6$. The combination of the X-ray data presented here with the objective selection of the PDCS, makes this sample of 31 distant clusters ideal for further studies of cluster evolution. We are actively engaged in following up this sample with photometry and spectroscopy. These new data will allow us not only to determine which cluster candidates represent real physical systems, but also to estimate the volume density of X-ray emitting clusters and to probe the relationship between velocity dispersion and X-ray luminosity. These various studies will allow us to quantify the degree to which cluster properties are evolving with redshift." + }, + "9708/astro-ph9708081_arXiv.txt": { + "abstract": "The youngest stellar populations of a 'J'-shaped region inside the supergiant shell (SGS) LMC$\\,$4 have been analysed with CCD photometry in $B$, $V$ passbands. This region consists of 2 coherent strips, one from the east to west reaching about $400\\;$pc across the OB superassociation LH$\\,77$ and another extending about $850\\;$pc from south to north. The standard photometric methods yield 25 colour-magnitude diagrams (CMDs) which were used for age determination of the youngest star population by isochrone fitting. The resultant ages lie in the range from $9\\;$M\\Jahr\\ to $16\\;$M\\Jahr\\ without correlation with the distance to the LMC$\\,$4 centre. We therefore conclude that there must have been one triggering event for star formation inside this great LMC SGS with a diameter of $1.4\\;$kpc. We construct the luminosity function and the mass function of five regions consisting of 5 fields to ensure that projection effects don't mask the results. The slopes lie in the expected range ($\\gamma \\in [0.22;0.41]$ and $\\Gamma \\in [-1.3;-2.4]$ respectively). The greatest values of the slope occur in the north, which is caused by the absence of a young, number-dominating star population. We have calculated the rate with which supernovae (SNe) have exploded in LMC$\\,$4, based on the finding that all stars are essentially coeval. A total of 5--7$\\cdot10^3$ supernovae has dumped the energy of $10^{54.5}\\;$erg over the past 10$\\;$Myr into LMC$\\,$4, in fact enough to tear the original star-forming cloud apart in the time span between 5 and 8$\\;$Myr after the starformation burst. We conclude that LMC$\\,$4 can have been formed without a contribution from stochastic self-propagating star formation (SSPSF), although the ring of young associations and \\HII\\ regions around the edge have been triggered by the events inside LMC$\\,$4. ", + "introduction": "In the Magellanic Clouds (MCs) giant loops of \\HII\\ regions can be recognized in deep pictures taken in H$\\alpha$ light. These \\HII\\ structures were divided by Goudis \\& Meaburn (1978) into two distinct groups: shell structures called giant shells (GSs) with diameters of 20--$260\\;$pc and the huge supergiant shells (SGSs) with diameters of 600--$1\\,400\\;$pc. Employing unsharp masking techniques and high contrast copying of the long exposed H$\\alpha$ images of Davies et al. (1976, hereafter DEM), Meaburn and collaborators (see Meaburn 1980) subsequently identified 85 GSs and 9 SGSs in the Large Magellanic Cloud (LMC) and 1 SGS in the Small Magellanic Cloud (SMC). Unlike GSs whose stucture can be explained by the combined activity of supernova explosions, stellar winds and radiation pressure of the central star association(s), SGSs (i.e. structures about $1\\;$kpc in diameter) need very effective, gigantic mechanisms for their creation. Such large scale features, which rival in size only with spiral structures, might be created, according to the appraisal of mechanisms by Tenorio-Tagle \\& Bodenheimer (1988), by the collision of high velocity clouds (HVCs) with the disk of the galaxy or by stochastic self-propagating star formation (SSPSF). The infall of an HVC would cause a well defined velocity deviating from the velocity of the original disk gas. Also, because of the short time scale of such a collision, an almost identical age is expected for all stars whose formation was triggered by that event. In the case of SSPSF (see Feitzinger et al. 1981) star formation propagates from one point (e.g. the centre of a SGS) in all directions (e.g. towards the rim of a SGS). Thus one should see in projection two velocity components in the \\HI\\ layer (a receeding and an approaching component) as well as a clear age gradient of the star populations inside a SGS. However, the observation of a central area of a SGS with little gas (i.e. a 'hole' in the \\HI\\ layer) and of an approximately $200\\;$pc thick shell of neutral material, which is ionized at the inner edge (visible as H$\\alpha$ filaments) by the radiation from the associations containing numerous early type stars (see Lucke \\& Hodge 1970; Lucke 1974) in the centre of the SGS, would be consistent with both scenarios. The importance of understanding the formation and the structure of SGSs is evident if one realizes how big they are. They may dominate large portions of a galaxy, in particular of the smaller irregular galaxies, such as the Magellanic Clouds are. Knowledge about the creation of SGSs leads to a better knowledge of the more recent development of the MCs and of their youngest star populations. The shape and size of LMC$\\,$4 is such, that SSPSF has been considered as the most likely explanation for the formation of this SGS (see Dopita et al. 1985). However, observations of star groups along the edge of LMC$\\,$4 showed that the ages of these groups were of the same order as the time needed to create the SGS in the case of SSPSF (see Vallenari et al. 1993; Petr et al. 1994). Reid et al. (1987) found from $V$, $I$ photometry of Shapley Constellation III (the southern half of the SGS LMC$\\,$4, McKibben Nail \\& Shapley 1953) no clear age gradient, also inconsistent with a global SSPSF model. Furthermore, Domg{\\\"o}rgen et al. (1995) presented an investigation of LMC$\\,$4 based on \\HI\\ data and IUE spectra. One result is that there is only one distinct velocity component towards us with $10\\;\\kms$ and just a diffuse rear component. This is not consistent with an undisturbed expanding shell, but it indicates a break-out of the SGS at the back side of the LMC. So LMC$\\,$4 resembles a cylinder rather than a sphere which should be expected in a galaxy with an \\HI\\ scale height well below $500\\;$pc. We note, however, that it is notoriously difficult to derive depth structure in a reliable way. \\begin{figure} \\epsfxsize=8.95cm \\centerline{\\epsffile{fig1.ps}} \\caption[]{H$\\alpha$ image of supergiant shell LMC$\\,$4 made from a scan of a photographic plate taken with the Curtis Schmidt telescope at Cerro Tololo (Kennicutt \\& Hodge 1986). The locations of the important objects indicated in the text and of the 5 regions of the 1st dataset (named a--e, see Sect. 4) are marked} \\end{figure} All these points show that we still do not well understand the history of LMC$\\,$4. To improve this situation, photometry of star groups {\\it inside} LMC$\\,$4 was taken in 1993 with the goal to derive ages. This paper presents that data and the results. Additionally the star formation history is further investigated in Sect. 4 by looking for overlapping age groups in the derived mass functions. ", + "conclusions": "The ages found for the stars inside LMC$\\,$4 lie in the range of $9\\;$Myr to $16\\;$Myr without a visible correlation with the distance to the LMC$\\,$4 centre (see Table 4 and Fig. 3). This means that there must have been one triggering event for star formation inside the greatest SGS of the LMC with a diameter of $1.4\\;$kpc, whereas at the rim of LMC$\\,$4 there is evidence for SSPSF (Petr et al. 1994; Petr 1994). Thus SSPSF can't be the creation mechanism of LMC$\\,$4 as claimed by the first models (see e.g. Dopita et al. 1985), even if one is able to find examples of SSPSF in this region on scales below $150\\;$pc." + }, + "9708/astro-ph9708032_arXiv.txt": { + "abstract": "We examine the distribution of masses of black holes in transient low mass X-ray binary systems. A Bayesian analysis suggests that it is probable that six of the seven systems with measured mass functions have black hole masses clustered near seven solar masses. There appears to be a significant gap between the masses of these systems and those of the observed neutron stars. The remaining source, V404~Cyg, has a mass significantly larger than the others, and our analysis suggests that it is probably drawn from a different distribution. Selection effects do not appear to play a role in producing the observed mass distribution, which may be explained by currently unknown details of the supernova explosions and of binary evolution prior to the supernova. ", + "introduction": "The strongest case for the existence of black holes in nature is provided by a subclass of X-ray transients. The strong episodic X-ray emission from these binary stars demonstrates the existence of an accreting compact object. Radial velocity measurements of the companion star can be used to determine the mass function, which is a strict lower limit to the mass of the compact accretor. In a number of cases, this lower limit is above the upper limit of neutron star stability of $3M_{\\odot }$ (McClintock \\& Remillard 1986, Casares \\& Charles 1992, Remillard, McClintock \\& Bailyn 1992, Bailyn et al. 1995, Filippenko et al. 1995, Remillard et al. 1996). In these cases the compact object must be a black hole. Studies of the supernovae explosions which presumably give rise to these black holes have now progressed to the point where meaningful statements can begin to be made about the expected mass distribution of the black holes. Studies of galactic chemical evolution strongly imply that stars with initial masses above $\\approx 30M_{\\odot }$ must ``swallow'' many of their heavy elements during (or shortly after) the supernova event, and therefore should form relatively massive black holes (Maeder 1992, Timmes, Woosley \\& Weaver 1995). Detailed models of the supernovae explosions themselves (e.g. Timmes, Woosley \\& Weaver 1996) can reproduce this result by applying a fixed amount of kinetic energy to the outer layers of the star, and determining how much material recollapses onto the central object. However, there remain many undertainties in the relation between the initial mass of the star and the mass of the compact remnant left behind by the supernova. The amount of kinetic energy available, how it is transferred to the ejected material, details of the pre-supernova evolution of massive stars (especially relating to convection and mass loss) and the possible influence of a binary companion are all poorly understood. In this paper, therefore, we take an empirical approach --- we attempt to use the available observational evidence for stellar black holes with low mass companions (high mass systems such as Cygnus~X-1 presumably having followed a different evolutionary path), and see what constraints can be made upon the underlying distribution of black hole masses. The mass of the black hole primary $m_1$ is related to the observed mass function $f(m)$ in the following way: \\begin{equation} m_1={{f(m)(1+q)^2}\\over {\\sin ^3 i}} \\end{equation} In order to determine the mass of the black hole, one needs to know the mass ratio $q$ and the inclination $i$ of the binary system. These parameters can be measured, although somewhat more indirectly than the mass function, as discussed in section~2 below. Section 2 also presents a compilation of the data obtained on those low mass X-ray transients with measured mass functions which do not display X-ray bursts (generally considered to be a neutron star signature). In section 3, we use these data to examine the distibution of black hole masses following the Bayesian analysis employed by Finn (1994) in his study of neutron star masses. Our results suggest that V404~Cyg is not drawn from the same population as the other six sources, and that the group of six have black hole masses which cluster around $7M_{\\odot }$ with a significant gap between their distribution and that of the neutron stars. In section 4 we speculate about the implications of our results. ", + "conclusions": "The mass distribution of stellar mass black holes can in principle be predicted by models of the evolution and supernova explosions of massive stars. Timmes et al. (1996 --- hereafter TWW) for example find that compact remnants of supernovae become significantly more massive than the pre-supernova iron core when the progenitor's initial mass exceeds $30M_{\\odot }$. Above this mass, the kinetic energy of the explosion is insufficient to unbind the entire mantle and envelope of the star, so some of the outer regions fall back hole the core, increasing the mass of the remnant. Studies of galactic chemistry confirm that stars above $30M_{\\odot }$ cannot return their entire outer regions to the interstellar medium (Maeder 1992, Timmes et al. 1995). Taking the reasonable assumption that the applied kinetic energy is only weakly dependent on the progenitor's initial mass, TWW find a smooth monotonic relation between remnant and progenitor mass (see their Figure 1). If the initial mass function is weighted toward lower mass stars, as is almost certainly the case, the resulting mass distribution of black holes should be smooth, and strongly biased toward masses near those of the neutron stars. While a sample including only seven examples cannot provide unambiguous statistical results, our work strongly suggests that the expected distribution is not observed. As discussed above, it is likely that one of the seven sources is not drawn from the same distribution as the other six. These six black holes may well be tightly clustered in mass near $7M_{\\odot }$, althought broader distributions are also possible. Finally, there is a significant gap between the lower mass limit of this distribution and the upper limit of the observed neutron star masses given by Finn (1994). The expected pile-up of sources at masses near to that of the neutron stars appears to be ruled out. Thus, although the small number of sources prevents us from drawing completely compelling conclusions, the evidence against the kinds of distributions implied by the work of TWW is strong enough to warrant serious consideration of the ways in which those distributions could be significantly modified. A possible caveat is that this last conclusion might arise from observational selection effects. At present, however, we cannot think of any that would prevent systems with black hole masses near the neutron star upper limit from being detected. These sources are easy to identify --- in outburst they are among the brightest objects in the X-ray sky. Also, the clear observation of similar systems containing accreting neutron stars strongly suggests that the mass of the compact object is not critical in the identification of these sources. Note that we have included \\it all \\rm transient systems with measured mass functions in our sample, except those like Cen~X-4 (McClintock \\& Remillard 1990) which display type I X-ray bursts, which are convincing signatures of the presence of a neutron star. In particular, we included GRO~J0422+32, which is often left off lists of confirmed black holes because its mass function is well below the limit of $3M_{\\odot }$ generally taken to be the firm upper limit for neutron star stability. It is conceivable that some effect suppresses the disk instability cycle over some mass range of the primary star, thus preventing the transient behavior needed to both identify the source in the first place, and then subsequently measure the mass function. Once again, however, the existence of transient sources with neutron star primaries and recurrence timescales similar to or smaller than those of the black hole systems makes this solution implausible. The physical effects which might influence the black hole distribution can be divided into two classes: those which involve the supernova explosion itself, and those related to the binary nature of the observed systems. Regarding the supernova explosions, if the relation between the amount of fall-back material to the mass of the precursor were close to a step function, then one might imagine that if \\it any \\rm significant amount of material were to fall back, there would be enough to bring the total remnant mass up to $\\approx 7M_{\\odot }$. In this case, even if the amount of fallback material increased with increasing precursor mass after the initial step, one would still expect a strongly peaked distribution due to the steep expected mass function of massive stars. One conceivable scenario which might lead to such a step function is if some chemical transition within the precursor immediately prior to the supernova occurred just exterior to $\\approx 7M_{\\odot }$. Then the higher density and atomic number interior to the transition might result in precisely that portion of the star recollapsing, while the outer regions are expelled. Unfortunately, evolutionary calculations of very massive stars are subject to a number of important uncertainties (Woosley \\& Weaver 1995) particularly relating to mass loss rates, so this suggestion is hard to evaluate. It may be that the observed distribution of black hole masses is not a result of type II supernova explosions in general, but rather a consequence of the binary nature of the observed systems. All of the orbital periods are sufficiently small that considerable mass transfer must have occurred prior to the supernova explosion. Since the precursor of the supernova was almost certainly more massive than its companion, it would have filled its Roche lobe first, resulting in dynamically unstable mass transfer. This in turn would result in a common envelope configuration, leading to the expulsion of much of the outer envelope of the more massive star and a dramatic decrease in the orbital separation on a very short timescale. Such abrupt mass loss has been shown to dramatically change the nature of the subsequent supernova explosion (Brown, Weingartner \\& Wijers 1996). However it is not clear how this set of circumstances could lead to a narrow range of remnant masses, or to a gap, since the effects of common envelope evolution are strongly dependent on the initial binary separation, which presumably is broadly distributed. The mass of the black hole may also be influenced by mass accretion subsequent to the supernova event. If this effect changed the mass of the primary significantly, one would expect to see broadly distributed masses for the black hole, since the system should have undergone varying amounts of accretion after the formation of the compact object. Such a broadened distribution cannot currently be ruled out, but confirmation of the suggested sharp mass distribution would argue strongly against significant post-SN mass enhancement. In most cases there are also evolutionary arguments against post-SN mass enhancement. Five of the seven systems in our sample contain late-type main sequence secondaries. These stars have masses $<1M_{\\odot }$, and it would require considerable fine tuning if they had all begun with much more massive secondaries. GRO~J1655-40 has a secondary in the Hertzsprung gap --- considerable ingenuity is required to reconcile this with the required low mass transfer rate for transient behavior (Kolb et al. 1997), and it is not clear how much mass transfer might have already taken place. The remaining system is V404~Cyg, whose secondary is at the base of the giant branch. In this case, considerable material may indeed have been transferred, which might conceivably account for the anomalously high mass of this particular black hole. Thus, while there are a number of poorly understood effects which might alter the distribution of post-supernova remnant masses, it is not immediately obvious how these effects could combine to produce the kind of distribution favored by our analysis. Obviously, more examples of black hole systems and better measurements of known systems will result in more precise observational constraints using techniques such as the ones outlined here. But already it appears likely that we will have to consider new underlying mechanisms for the origin of the Galactic black hole mass distribution." + }, + "9708/astro-ph9708204_arXiv.txt": { + "abstract": "We present continuum observations of a small sample of high-redshift, radio-loud AGN (radio galaxies and quasars) aimed at the detection of thermal emission from dust. Seven AGN were observed with IRAM and SEST at 1.25\\,mm; two of them, the radio galaxies 1243+036 ($z \\sim 3.6$) and MG\\,1019+0535 ($z \\sim 2.8$) were also observed at 0.8\\,mm with the JCMT submillimetre telescope. Additional VLA observations were obtained in order to derive the spectral shape of the synchrotron radiation of MG\\,1019+0535 at high radio frequencies. MG\\,1019+0535 and TX\\,0211$-$122 were expected to contain a large amount of dust based on their depleted Ly$\\alpha$ emission. The observations suggest a clear 1.25-mm flux density excess over the synchrotron radiation spectrum of MG\\,1019+0535, suggesting the presence of thermal emission from dust in this radio galaxy, whereas the observations of TX\\,0211$-$122 were not sensitive enough to meaningfully constrain its dust content. On the other hand, our observations of 1243+036 provide a stringent upper limit on the total dust mass of $<10^8$\\,M$_{\\odot}$. Finally, we find that the spectra of the radio-loud quasars in our sample ($2 < z < 4.5$) steepen between rest-frame radio and the far-infrared. We discuss the main implications of our results, concentrating on the dusty radio galaxy, MG\\,1019+0535. ", + "introduction": "In recent years several attempts have been made to study the interstellar medium in high-redshift galaxies. Promising results have been obtained by millimetre and submillimetre observations, which sample the rest-frame far-IR emission of galaxies at $z > 2$ (Andreani et al.\\ 1993). The study of dust in distant galaxies is very important because it provides information on the physical state of the ISM, and its relation to other properties of the galaxies, such as activity and evolution. For example, if it is correct to assume that the rest-frame far-IR luminosity $L_{\\rm FIR}$ is a measure of the star-formation rate, then this could provide a way to estimate the evolutionary state of high-$z$ galaxies. Establishing whether dusty galaxies are common at high redshifts is also important in order to evaluate the effects of dust obscuration on surveys of quasars and protogalaxies carried out in optical bands (Smail, Ivison \\& Blain 1997). Several active galaxies with $z>2$ have been detected in submillimetre/millimetre bands, suggesting the presence of large amounts of dust in their host galaxies (M$_{\\rm d} \\sim 10^{8-9}$ M$_{\\odot}$; see Hughes, Dunlop \\& Rawlings 1997 for a recent review; Omont et al.\\ 1996b; Ivison et al.\\ 1998). At high redshift, dust is also thought to be present in damped Ly$\\alpha$ absorption systems (Pettini et al.\\ 1994; Pei, Fall \\& Bechtold 1991), in very red galaxies (Hu \\& Ridgway 1994), and in high-$z$ radio galaxies (Cimatti 1996 and references therein). A substantial amount of dust is also expected in theoretical models of the evolution of galaxies at high-$z$ (Mazzei \\& De Zotti 1996 and references therein). Finally, it is important to recall that molecular gas has been observed in a few distant active galaxies, allowing a direct estimate of the dust-to-gas mass ratio at large cosmological distances for the first time (see Omont et al. 1996a; Ohta et al.\\ 1996). We recall that particular caution is needed in the interpretation of the submillimetre/millimetre data of radio-loud objects. The existence of a thermal dust emission excess should be confirmed by checking whether the submillimetre/millimetre flux densities represent a real excess over the synchrotron spectrum. The radio galaxy B2\\,0902+343 ($z\\sim3.4$) is an example of the problem, where the observed 1.3-mm flux density is not due to thermal emission from dust, but to the tail of the radio synchrotron spectrum (Downes et al.\\ 1996 and references therein). Therefore, high-frequency ($\\sim8-43$\\,GHz) radio observations are extremely important to derive the shape of the synchrotron spectrum and to estimate its contribution to the millimetre region. In the present paper we present the results of millimetric observations of a small sample of radio galaxies and radio-loud quasars with $z>2$, and additional high-frequency VLA observations of MG\\,1019+0535, a radio galaxy at $z\\sim2.8$. Throughout this paper we assume $H_0=50$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $q_0=0.5$ and define $h_{50}=H_0/50$. ", + "conclusions": "We have presented observations of the rest-frame far-infrared continuum of a small sample of radio-loud AGN (4 radio galaxies and 3 quasars) with $22$ can be rather inhomogeneous. Finally, the three flat-spectrum radio-loud quasars ($2$23. However, following Patterson \\&\\ Halpern's (1990) suggestion that it was an intermediate polar (IP -- a magnetic white dwarf accreting material from a close companion), Kamata, Tawara \\&\\ Koyama (1991) discovered deep X-ray eclipses recurring with a 6.06-h orbital period. XY~Ari is now the only known IP with a deep X-ray eclipse (EX~Hya shows a partial X-ray eclipse while DQ~Her shows a deep eclipse but negligible X-ray flux; see Patterson 1994 for a review of these stars). The identification of XY~Ari in the infra-red (Zuckerman \\etal\\ 1993) allows modelling of the ellipsoidal variations in the light curve (Allan, Hellier \\&\\ Ramseyer 1996), tying down the system parameters. Kamata \\&\\ Koyama (1993) made the first attempt to use the X-ray eclipse to constrain the accretion geometry, using 5 ingresses and egresses observed with {\\it Ginga.} They concluded that the time of ingress/egress was comparable to the expected size of the white dwarf, and showed the potential of observing more eclipses at a higher count rate. Accordingly, I proposed an observation of 20 eclipse egresses with the larger PCA on the {\\it RXTE\\/} satellite (Bradt, Rothschild \\&\\ Swank 1993). The aim was to record changes of the eclipse egress with the phase of the 206-s white dwarf spin cycle. I chose to observe only egresses so that they were all directly comparable: at ingress the secondary star limb is at a different angle on the white dwarf, potentially providing more information, but requiring more eclipses to look for patterns. The motivation for the observations is that after 15 years of studying IPs there is no agreement on the pattern of accretion flow onto the white dwarf, with estimates of $f$, the fraction of the white dwarf onto which accretion occurs, differing by orders of magnitude. King \\&\\ Shaviv (1984), Norton \\&\\ Watson (1989) and Patterson (1994) all argue for $f >$ 0.1--0.3; while Rosen, Mason \\&\\ Cordova (1988), Hellier, Cropper \\&\\ Mason (1991) and Hellier (1995) claim $f$ = 0.01--0.001. This paper presents direct geometric evidence addressing this issue. \\begin{figure*}\\vspace{8cm} % \\caption{One of the 20 {\\it RXTE\\/} lightcurves of XY~Ari, showing the star emerging from eclipse and pulsing at the 206-s spin period. Each bin lasts 10 s. The errors can be judged from the data in eclipse, since the high background level means that photon noise is nearly constant.} \\special{psfile=xy1fig1.ps hoffset=-50 voffset=-63 hscale=70 vscale=70 angle=0} \\end{figure*} ", + "conclusions": "(1) The majority of the X-ray flux emerges from eclipse in $< 2$ s, implying that the accretion polecaps have areas of $<0.002$ as a fraction of the white dwarf surface. \\noindent (2) The accretion footprints are not fixed in position on the white dwarf, but can wander over an area $<0.01$ of the white dwarf surface. \\noindent (3) There is evidence that a minority of the flux arises from a much larger area, or that we can always see part of the accretion regions at both poles. \\noindent (4) The accretion regions are at a latitude in the range 43\\deg--63\\deg, which suggests a magnetic dipole offset of 8\\deg--27\\deg\\ (constraints on other system parameters are listed in Table~1). \\noindent (5) We see roughly equal amounts of X-ray flux from both upper and lower poles. There are, though, slight asymmetries so that the disappearance of one pole over the white dwarf limb does not entirely compensate for the appearance of the other pole. The net effect is a low-amplitude, complex spin pulse." + }, + "9708/astro-ph9708183_arXiv.txt": { + "abstract": "A gamma-ray burst (GRB) releases an amount of energy similar to that of a supernova explosion, which combined with its rapid variability suggests an origin related to neutron stars or black holes. Since these compact stellar remnants form from the most massive stars not long after their birth, gamma-ray bursts should trace the star formation rate in the Universe; we show that the GRB flux distribution is consistent with this. Because of the strong evolution of the star formation rate with redshift, it follows that the dimmest known bursts have $z\\sim6$, much above the value usually quoted and beyond the most distant quasars. This explains the absence of bright galaxies in well-studied gamma-ray burst error boxes. The increased distances imply a peak luminosity of $8.3\\times10^{51}\\un{erg}{}\\un{s}{-1}$ and a rate density of 0.025 per million years per galaxy. These values are 20 times higher and 150 times lower, respectively, than are implied by fits with non-evolving GRB rates. This means that GRBs are either caused by a much rarer phenomenon than mergers of binary neutron stars, or their gamma-ray emission is often invisible to us due to beaming. Precise burst locations from optical transients will discriminate between the various models for GRBs from stellar deaths, because the distance between progenitor birth place and burst varies greatly among them. The dimmest GRBs are then the most distant known objects, and may probe the Universe at an age when the first stars were forming. ", + "introduction": "The ability of the Italian-Dutch BeppoSAX satellite \\cite{psb:95} to accurately locate gamma-ray bursts (GRBs) with its Wide Field Cameras \\cite{jhzb:95} has led to the discovery of two optical transients, associated with GRB\\,970228 \\cite{pggks:97} and GRB\\,970508 \\cite{bond:97}. The detection of redshifted absorption lines in the optical transient associated with GRB\\,970508 \\cite{mdska:97} has established that it lies at cosmological distance, and here we assume they all do. A typical dim burst ($1\\un{ph}{}\\un{cm}{-2}\\un{s}{-1}$ lasting 10\\,s) at $z=1$ releases $3.5\\times10^{49}\\un{erg}{}\\un{sr}{-1}$ in gamma rays, hence the engine behind it must provide about $4.4\\times10^{52}f\\sub{b}/\\epsilon_{-2}\\un{erg}{}$. (Where needed, we assume $H_0=70\\un{km}{}\\un{s}{-1}\\un{Mpc}{-1}$ and $\\Omega_0=1$.) Here $f\\sub{b}$ is the fraction of the sky illuminated by the gamma-ray emission, and the efficiency of conversion of the initially available energy into gamma rays is $\\epsilon_{-2}$ percent. This is the natural energy scale of supernova explosions, in which of order $10^{53}$\\,ergs is released suddenly from the rest mass energy of a solar mass of material. The bulk is carried away in neutrinos, and about 1\\% becomes kinetic energy of the ejecta. The proposed GRB models related to end stages of massive stars are (i) merger of two neutron stars \\cite{paczy:86,gdn:87,elps:89}; (ii) merger of a neutron star and a black hole \\cite{paczy:91,mhim:93}; (iii) `failed supernova': the collapse of a massive star to a black hole surrounded by a dense torus of material that might result in a relativistic jet \\cite{woosl2:93}; (iv) a `hypernova': the collapse of a rapidly rotating massive star in a binary \\cite{paczy:97}; (v) collapse of a Chandrasekhar-mass white dwarf \\cite{usov:92}. Whether these are efficient enough at converting a fraction of the available energy to kinetic energy and then eventually to gamma rays (see below) is an open question, and the major unsolved issue in this class of burst models. In this paper we assume that somehow a variety of such a model manages this. The important point is that they all arise from massive stars which evolve into remnants within about 100\\,Myr. The binary mergers then usually take place within about 100\\,Myr of remnant formation \\cite{ps:96}, as does the white-dwarf collapse, because the favoured route has a high mass transfer rate \\cite{hbnr:92}. Since the expansion age of the Universe is already 1\\,Gyr at $z=4.4$, it is safe to neglect the delay between (binary) stellar birth and the GRB it eventually yields in the present context. {\\it The gamma-ray burst rate therefore traces the massive star formation rate.\\/} The star formation rate as a function of redshift has recently been studied extensively, and is determined observationally with some confidence \\cite{llhc:96,mfdgs:96,madau:96}: the luminosity density in the rest frame $U$ and $B$ band is combined with an IMF to deduce the star formation rate. The assumption of an IMF introduces an uncertainty in the deduced total star formation rate, but the basic data ($U$ and $B$ light density) are dominated by massive stars. Since GRBs come from massive stars, they may trace the UV light density in the Universe better than the total star formation rate, and our results are therefore less sensitive to the assumed IMF. A further potential source of uncertainty is dust extinction, which would cause a relative underestimate of the high-redshift star formation rate. Recent interpretations of the afterglows \\cite{mr:97,wrm:97,waxma:97} support the notion that the energy release is initially in the form of an ultrarelativistic explosion or `fireball' \\cite{cr:78,goodm:86} whose energy is largely converted to a blast of gamma rays via hydrodynamic collisions within it \\cite{px:94,rm:94} or with the ambient medium \\cite{rm:92}. Since the kinetic energy comes from a fairly standardised event, it is likely that the gamma-ray luminosity distribution of bursts is not too wide, so we shall treat them as standard candles. With the fact that GRB trace star formation, this has the important testable consequence that the redshift dependence of the GRB rate has no free parameters. Only two normalisations need to be fitted, namely the local GRB rate density $\\rho_0$ and the standard-candle 30--2000\\,keV luminosity $L_0$. ", + "conclusions": "Our results depend only on two key features of the star formation rate: rapid evolution up to $z\\approx 1$ and a more gentle variation beyond that up to $z \\approx 2.5$. Although little is known about the star formation rate between these redshifts, many indirect arguments suggest that the star formation rate did not evolve strongly between $z=1$ and $z=2.5$. Uncertainties are introduced by our lack of knowledge of the correction due to dust extinction and to a lesser extent by having to assume an IMF. However this does not change the qualitative nature of the results presented here. Furthermore, the distance corrections due to the wide range in spectral shapes \\cite{fb:95} that are the result of assuming that GRB are standard candles in the 30--2000\\,keV (rest frame) range are very important, and neglect of the spectral shape variation can lead to considerably different results. In view of this, and of the fact that we have not explored a range of parameters for the geometry of the Universe, our results obviously have a somewhat exploratory character. Other observational tests, such as the determination of redshifts from afterglow spectra and constraints from the lensing rate of GRBs, will provide additional constraints on parameters. The logical consequence of assuming that GRB are related to remnants of the most massive stars in any of the ways hitherto proposed is that the GRB rate is proportional to the formation rate of massive stars in the Universe (with the possible exception of NS-NS mergers if those simulations which indicate significant fractions of late mergers are correct). We show that this assumption is consistent with the GRB flux distribution. Compared to previous, non-evolving, models of cosmological bursts we find a twenty-fold increase of the required GRB luminosity, which suggests that the gamma-ray emission is significantly beamed in order that the emitted energy can be supplied by merger/collapse models. The redshifts of GRBs are also greatly increased, and the very dimmest known ones are at $z\\gsim6$, beyond the farthest known quasars. This makes them the most distant known objects, and their optical counterparts very valuable probes of the early evolution of stars and interstellar gas." + }, + "9708/astro-ph9708197_arXiv.txt": { + "abstract": "We present the preliminary results of Monte Carlo simulations aimed to investigate the effects of realistic dust extinction (absorption + scattering) on the colours of high-$z$ galaxies. In this paper, we concentrate on the case of spheroidal galaxies, and we obtain attenuation curves in the range 0.1-1.2$\\mu$m for different dust spatial distributions, and for a range of values of the dust optical depth, geometrical thickness and inclination of the dust disk. We find that the resultant curves are strongly dependent on the dust geometrical distribution and optical depth. A serendipitous finding is that the strength of the 2200~\\AA~ absorption feature depends not only on the optical depth, but also on the dust geometrical distribution. As a first application, we test our results on two high-$z$ galaxies with extremely red colours (HR10 and HR14) in order to infer clues on their ages, dust content and dust spatial distribution. We confirm that HR10 must be a very dusty galaxy, and we suggest that its stellar component should be strongly embedded in the dust in order to reproduce the observed extremely red colours. ", + "introduction": "Internal dust might affect seriously the observed properties of distant objects, and observations have indeed found evidence of dust in high-$z$ objects. Several active galaxies with $z>$2 have been detected at submillimeter wavelengths, suggesting the presence of large amounts of dust (M$_{dust} \\sim 10^{8-9}$ M$_{\\odot}$; see Hughes 1996 for a recent review). The reddening of the background quasars and their metal abundances suggest the presence of dust in damped Ly$\\alpha$ absorption systems (Pettini et al. 1994; Pei \\& Fall 1996). Also, the UV polarization properties of high-$z$ radio galaxies can be explained in term of dust scattering, suggesting a significant amount of dust in their ISM (Cimatti 1996 and references therein). A substantial amount of dust is also expected in evolutionary models of spheroidal galaxies at high-$z$ (Franceschini et al. 1994; Mazzei \\& De Zotti 1996 and references therein). Although dust is likely to be present in most high-$z$ systems, no information is available neither on its spatial distribution, nor about how {\\it realistic} dust extinction can affect our view of distant galaxies. In fact, extinction is usually treated in a simplistic way, neglecting the contribution of scattering, and assuming naive spatial distributions (e.g. uniform foreground screens or infinite slabs). Only recently, more realistic models have been developed (Kylafis \\& Bahcall 1987; Bruzual et al. 1988; Witt et al. 1992; Byun et al. 1994; Wise \\& Silva 1996; Bianchi, Ferrara \\& Giovanardi 1996, hereafter BFG). Although these models have shown that dust scattering plays an important role by reducing the effects of dust absorption, no systematic studies of the effects of the extinction on the observed colours of high-$z$ galaxies have been performed. Understanding these effects is crucial in cosmology and galaxy evolution studies. For instance, a relevant problem arises in the age estimates of high-$z$ galaxies. When deep spectroscopy is not available, no information on the stellar continuum and absorption features are obtainable, and the age estimates are based solely on the fitting of the broad-band photometric Spectral Energy Distributions (SEDs) with synthetic stellar population spectra. In a simplistic scenario where the extinction is entirely due to absorption (the screen model), it is well known that the colours of a reddened young galaxy can mimic those of an unreddened old galaxy, producing a degeneracy of the age and dust extinction. However, what happens to the integrated colours of a galaxy when dust extinction is treated more realistically is poorly known. Witt et al. (1992) suggested that the effects of realistic dust extinction on the SEDs of high-$z$ galaxies are not strong. However, their claim was based on a single object (the radio galaxy B2 0902+34 at $z\\sim3.4$), and generalized to the whole population of high-$z$ galaxies. In addition, the continuum SED of B2 0902+34 turned out to be completely different from that used by Witt et al. (1992) because of a strong contamination of the $K$-band flux by the redshifted [OIII]$\\lambda\\lambda$4959+5007 (Eisenhardt \\& Dickinson 1992). Instead, recent studies (Franceschini et al. 1994) have emphasized the role of dust in the early stages of evolution of galaxies. The present optical surveys for high-$z$ galaxies are designed to select objects with strong emission lines, such as Ly$\\alpha$, (see Pritchet 1994 for a review; Thompson et al. 1995), or galaxies with a flat continuum spectrum and with a sharp Lyman-break (Steidel et al. 1996). However, because of those selection criteria, these surveys would miss the putative population of high-$z$ galaxies obscured by dust extinction. It is notable that the strongest evolutionary effects due to dust extinction are not expected for the galactic disks, or in disk dominated systems (Mazzei et al. 1992), but rather in early-type galaxies which, under appropriate circumstances, might experience a prolonged opaque phase (Mazzei \\& De Zotti 1996). The whole issue of the effects of dust on the SEDs of high-$z$ galaxies still remains an important, open question, especially in view of the future ISO and sub-mm data. Extinction effects become particularly relevant in the case of high-$z$ galaxies with very red colours, also called extremely red objects (EROs) (McCarthy et al. 1992; Hu \\& Ridgway 1994). In fact, these galaxies may be distant and old ellipticals, and could provide crucial clues on the first epoch of galaxy formation, $H_0$ and $q_0$, once their ages are derived accurately. However, they could also be very dusty high-$z$ galaxies whose intrinsic colours are strongly reddened by dust extinction. The main problem is that the degeneracy of the age and the dust extinction becomes maximized in this class of objects because of the similarity between the colours of a genuinely old galaxy and those induced by foreground screen reddening. The main questions are then : how much {\\it realistic} dust extinction can mimic the colours of an old galaxy ? If these galaxies are dusty, is it possible to learn something about the spatial distribution of the dust by modeling their SEDs ? Motivated by the general lack of information on the effects of dust extinction on the SEDs of high-$z$ galaxies, we have started an extensive study aimed at investigating the relevance of these effects. In this paper, we focus on the case of spheroidal galaxies, present the first results, discuss an application to the case of the extremely red galaxies, and show a serendipitous result about the strength of the 2200~\\AA~ dust absorption feature in external galaxies. More general results and applications will be presented in a forthcoming paper. ", + "conclusions": "We have presented the preliminary results of a set of simulations aimed to study the effects of realistic dust extinction on the colours of high-$z$ galaxies. As a first attempt, we have considered three different spatial distributions of the dust in spheroidal galaxies, adopting Galactic dust grain properties. A serendipitous finding is that the 2200~\\AA~ feature is generally much weaker than in the foreground absorption screen model. The weakening of the feature is particularly evident in the case of a disk of dust in a spheroidal galaxy, where the feature can be almost absent. We find that the equivalent width of the 2200~\\AA~ feature depends on the dust distribution, geometry and optical depth. We suggest that this effect is mostly a geometrical one, and that it might be one explanation of the failure of the detection of the 2200~\\AA~ feature in external dusty galaxies. However, other causes of the non-detection of 2200~\\AA~ feature, such as different chemical composition of the dust grains, cannot be ruled out by the present results. As a first application of our extinction curves, we consider the case of two extremely red galaxies at high redshift, HR10 and HR14, in order to constrain the ages of their stellar populations, the amount of dust, and the dust geometrical distribution. Our results confirm that HR10 must be a very dusty galaxy, and allow us to constrain the spatial distribution of the dust in this object. However, we find that a strong degeneracy between the age of the stellar population and the dust extinction is present in both galaxies. In the case of HR10, we find that its colours can be reproduced by a 1 Gyr old stellar population and $\\tau_V\\sim$6 (model B). Its dereddened rest-frame K-band absolute magnitude, $M_{K_{rest}}$=-27.5, is just beyond the highest luminosities of the local field galaxies. For HR14, although the redshift is poorly constrained, we find that its SED can be reproduced without extinction and with an old stellar population (2.5-4.0 Gyr) over the wide range 1.8$$1.8. However, if dust extinction is present, the age can be lowered down to 0.8 Gyr for $\\tau_V \\sim$2-3 and for a broad range of dust spatial distributions. It is important to note that the main results on HR14 are not strongly dependent on $z$, if 1.8$