Limits on the Gas Disk Content of Two ``Evolved" T Tauri Stars
Alan G. Aversa
We derived upper limits of the circumstellar gas disk masses around the T~Tauri stars St~34 and RX~J0432.8+1735 in order to place constraints on theories of planet formation and to explore the evolution of the gas-to-dust ratio during the epoch of disk dissipation around young sun-like stars. Since sub-millimeter lines of $^{12}$CO trace of the cold, outer regions of circumstellar disks, we observed $^{12}$CO~$J=2-1$ emission with the 10~m Sub-Millimeter Telescope (SMT) for two carefully chosen targets. St~34 is a rare classical T Tauri star with an age of $8\pm3$~Myr, and RX~J0432.8+1735 is a rare weak-emission T Tauri star with far-infrared excess. Both exhibit radial space motion enabling us to distinguish disk emission from ambient cloud material. Assuming a $^{12}$CO excitation temperature of 20~K, a $^{12}$CO line-width of 5~km~s$^{-1}$, and optically-thin emission, we derive $3\sigma$ upper limits on the H$_2$ circumstellar disk mass for St~34 and RX~J0432.8+1735 to be $<4.20$~M$_\earth$ for both disks. Placing these results in the context of other studies, we discuss their implications on planet formation models.
gas_disks_12_feb.tex
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\title{Limits on the Gas Disk Content of Two ``Evolved" T Tauri Stars}
\author{Alan G. Aversa}
\affil{Steward Observatory, 933 N Cherry Ave., Tucson AZ 85721}

\begin{abstract}
We derived upper limits of the circumstellar gas disk masses around the T~Tauri stars St~34 and RX~J0432.8+1735 in order to place constraints on theories of planet formation and to explore the evolution of the gas-to-dust ratio during the epoch of disk dissipation around young sun-like stars. Since sub-millimeter lines of $^{12}$CO trace of the cold, outer regions of circumstellar disks, we observed $^{12}$CO~$J=2-1$ emission with the 10~m Sub-Millimeter Telescope (SMT) for two carefully chosen targets. St~34 is a rare classical T Tauri star with an age of $8\pm3$~Myr, and RX~J0432.8+1735 is a rare weak-emission T Tauri star with far-infrared excess. Both exhibit radial space motion enabling us to distinguish disk emission from ambient cloud material. Assuming a $^{12}$CO excitation temperature of 20~K, a $^{12}$CO line-width of 5~km~s$^{-1}$, and optically-thin emission, we derive $3\sigma$ upper limits on the H$_2$ circumstellar disk mass for St~34 and RX~J0432.8+1735 to be $<4.20$~M$_\earth$ for both disks. Placing these results in the context of other studies, we discuss their implications on planet formation models.
\end{abstract}

\begin{document}
\maketitle

\section{Introduction}
%T Tauri stars are newly formed stars which frequently have circumstellar disks of gas and/or dust. They fall into two categories: Weak-Line T Tauri Stars (WTTSs), characterized by low H$\alpha$ equivalent widths, and Classical T Tauri Stars (CTTSs), with higher H$\alpha$ equivalent widths indicative of ongoing gas accretion.

%Planet formation depends upon the gas content of the circumstellar disk from which giant planets form. Dust circumstellar disks disappear after about 3~Myr, with a 1-10~Myr dispersion \citep{meyer07}. To understand the timescales of planet formation, knowledge of how long gas circumstellar disks persist around their parent stars is necessary. Infrared campaigns with {\it Spitzer}, such as the Cores to Disks (c2d) program \citep{padgett06}, provide knowledge of the dust circumstellar disks based upon IR excesses at 70~ $\micron$, which traces cool dust at larger radii, and IR excesses at 24~ $\micron$, which traces dust at within a few AU of the parent star. IR observations do not, however, constrain the total gas content in circumstellar disks. This is what we investigate for our two T Tauri stars.

%If all circumstellar disks contain the materials necessary to form gas giant planets, such as the majority of extrasolar planets detected today, there must be sufficient amounts of cold gas at larger radii than detected by {\it Spitzer} infrared observations. A temperature-radius law such as $T\sim r^{-0.5\pm0.75}$,
%\begin{equation}
%\label{templaw}
%T(r)=10^3~\mathrm{K}\,\Big(\frac{r}{0.1~\mathrm{AU}}\Big)^{-0.5\pm0.75},
%\end{equation}
%shows that at a {\it Spitzer} wavelength of 70~ $\micron$, one predominately detects emission at radii of $\sim$3~AU. The rotational energy transitions of carbon monoxide (CO), a proxy for molecular hydrogen (H$_2$), enable one to trace cold gas disks out to radii many times larger.

%If gas content and dust content follow similar timescales \citep{damajanov07}

%par #1 - gas and dust rich disks around young sun-like
%stars are a common outcome of the star formation process
%and have a range of initial properties that are becoming
%understood. See intro of Kundurthy et al. (2006).

Circumstellar disks of gas and dust, a natural result of the conservation of angular momentum, are a common outcome of the star formation process. \citet{kenyon95} find that over half the low mass ($<$ 2 M$_{\odot}$) pre--main-sequence T Tauri stars in the Taurus-Auriga star formation region have more infrared emission than expected from a normal stellar photosphere, indicating the presence of a dusty circumstellar disk heated by the parent star as well as active accretion. T Tauri stars fall into two categories: Weak-Line T Tauri Stars (WTTSs), characterized by low H$\alpha$ equivalent widths, and Classical T Tauri Stars (CTTSs), with higher H$\alpha$ equivalent widths indicative of ongoing gas accretion. These circumstellar disks generally have the following properties \citep{dutrey07,andrews09}: mass surface densities $\Sigma(r)\propto r^{0\,\mathrm{to}\,-1.0}$, surface temperatures $T(r)\propto r^{0\,\mathrm{to}\,-0.6}$(depending on disk flaring), and Keplerian rotational velocities $V(r)\propto r^{0.5}$.

%par #2 - theories of planet formation attempt to explain
%the observational outcomes of this process and are also
%constrained by observations of disk evolution such as
%the end of gas accretion, evolution in dust disk properties,
%and molecular line observations of gas.

Gas giant planet formation depends upon the gas content of the circumstellar disk from which they form. Primordial inner disks traced by hot dust disappear after about 3~Myr with a range of $1-10$~Myr \citep{meyer07}. If gas content and dust content in disks dissipate through similar mechanisms \citep{damajanov07}, then we would expect gas to disappear on these same timescales. To understand the timescales of planet formation, we must understand how long a gas disk persists around its parent star. Accretion rates (higher in CTTSs than WTTSs) also trace the time evolution of gas content as gas must be present to accrete. \citet{durisen07} suggest that gravitational instabilities in gas disks could account for their rapid ($<3$~Myr) dissipation, locking up mass in planets. In contrast, simulations of the evolution of gas disk surface density for a 1~M$_\sun$ star due to photo-evaporation indicate that gas disks disappear in about 6~Myr \citep{alexander06,dullemond07}. Because there are various theoretical results for the gas dispersal mechanisms and associated timescales, we might expect to observe diverse properties for gas disks around T Tauri stars.

%par #3 - mm-wave observations of CO have some advantages
%and disadvantages in this regard, and provided that the
%right stellar samples can be identified, more work in
%this area is needed.

Ideally, sub-millimeter interferometric images of T~Tauri stars would yield the most information about circumstellar disks, including their orientation, geometry, and gas content. Yet observations of this sort have only been published for a few nearby stars \citep[e.g., Sub-Millimeter Array (SMA) observations of TW~Hya;][]{qi04}. Infrared photometric campaigns with {\it Spitzer}, such the Cores to Disks \citep{evans03} and the Formation and Evolution of Planetary Systems \citep[FEPS;][]{meyer06} {\it Spitzer} Legacy projects, provide knowledge of the dust circumstellar disks based upon IR excesses at many wavelengths including $24\micron$, which trace dust within a few AU of the parent star, and at $70\micron$, which trace cooler dust at larger radii. \citet{silverstone06} surveyed 74 stars with ages between 3-30 Myr finding no stars with mid--IR excess that were not gas rich accreting T Tauri stars. \citet{padgett06} however identified a handful of WTTS (lacking signatures of accretion) with evidence for mid-- and far--IR excess emission. Photometric IR observations do not, however, constrain the total gas content in circumstellar disks because of a potentially highly variable gas-to-dust ratio. Observing the rotational energy transitions of $^{12}$CO, a proxy for molecular hydrogen (H$_2$), and assuming the ISM abundance ratio $\left[\mathrm{H}_2/\mathrm{CO}\right]\approx10^4$ enables one to trace the majority of cold gas in disks out to radii many times larger. Is the timescale for gas dissipation similar to the timescale for dust dissipation around T Tauri stars? This is the central question we address with our new observations. 

%par #4 - what this paper is and how it is structured.

Expanding on previous works \citep[e.g.][]{pascucci06,vankempen07}, we search for $^{12}$CO~$J=2-1$ emission for two ``evolved" T~Tauri stars in Taurus to place constraints on their gas circumstellar disk masses. St~34 is evolved in the sense that it is $8\pm3$~Myr old, whereas RX~J0432.8+1735 is evolved in the sense that---although it is much younger ($1.0\pm0.5$~Myr)---it has already lost its inner accretion disk. In \S2, we describe our selection of sources and observations, then in \S3 we derive upper limits on our sources' gas disk masses. Lastly, in \S4 we describe the implications of our results, and in \S5 we list our conclusions.

\section{Observations}

\subsection{Selection of Sources}
%%%IS IT KENYON95 OR KENYON90 IN FIRST CITATION HERE?%%%
We observed the $^{12}$CO~$J=2-1$ emission of the Taurus region objects St~34 and RX~J0432.8+1735, T Tauri stars with known 24 or $70\micron$ excesses \citep[see Table~\ref{sourcetab}]{kenyon95,padgett06}. In order to detect gas in their circumstellar disks, we must be able to distinguish emission from the disk and surrounding molecular cloud. From available candidate ``evolved'' (in age or shape of spectral energy distribution) T Tauri stars, we selected those least likely to be contamined by the ambient CO emission from the parent molecular cloud based on the \citet{dame87} CO survey. The radial velocities of our sources differ from the systemic velocities of any CO emission in the vicinity of our sources by $\sim2-3$~km~s$^{-1}$.

\begin{deluxetable}{lccccccccccll}%INCLUDE A REFERENCE FOR SOLAR MOTION IN THIS TABLE
\tabletypesize{\scriptsize}
\tablewidth{0pt}
\rotate
\tablecaption{Candidate Source Summary\label{sourcetab}}
\tablehead{\colhead{Object} & \colhead{$\alpha$} & \colhead{$\delta$} & \colhead{Heliocentric RV} & \colhead{Source $V_\mathrm{LSR}$\tablenotemark{a}} & \colhead{Cloud $V_\mathrm{LSR}$\tablenotemark{b}} & \colhead{$T_\mathrm{eff}$} & \colhead{$\log{L^\star}$} & \colhead{Mass} & \colhead{Age} & \colhead{General} \\
\colhead{} & \colhead{(J2000)} & \colhead{(J2000)} & \colhead{(km~s$^{-1}$)} & \colhead{(km~s$^{-1}$)} & \colhead{(km~s$^{-1}$)} & \colhead{(K)} & \colhead{($\log{L_\sun}$)} & \colhead{($M_\sun$)} & \colhead{(Myr)} & \colhead{Reference}}
\startdata
RX~J0432.8+1735 & 04 32 53.23 & 17 35 33.68 & $18.6  $ & 7.3 & 10.2 & 3499\tablenotemark{c} & $\sim -6.1\times10^{-2}$\tablenotemark{g} & \nodata & $1\pm 0.5$\tablenotemark{d} & a\\
St~34  & 04 54 23.70 & 17 09 54.00 & $17.9 \pm 0.6 $ & 6.6 & 8.37 & 3415 & $1.03\pm0.06$ & $0.37\pm0.08$\tablenotemark{e} & $8\pm 3$\tablenotemark{e,f} & b
\enddata
\tablerefs{For RVs, see \citet{wichmann00,white05}. For general references, see (a) \citet{padgett06} and (b) \citet{white05}.}
\tablenotetext{a}{We corrected the heliocentric radial velocities (RVs) from the literature into local standard of rest ($V_\mathrm{LSR}$) radial velocities assuming the sun moves toward J2000 $(\alpha = 18\fh0, \delta = 30\fdg0)$ at 20~km~s$^{-1}$.}
\tablenotetext{b}{Determined by the $^{12}$CO~$J=1-0$ emission from the \citet{dame87} survey}
\tablenotetext{c}{\citet{wichmann00}}
\tablenotetext{d}{\citet{dantona97}}
\tablenotetext{e}{For both binary components}
\tablenotetext{f}{Isochronal age is given. The Li depletion age for both binary components is $>25$~Myr.}
\tablenotetext{g}{Estimated by integrating the spectral energy distribution (SED) of RX~J0432.8+1735 in \citet{padgett06}}
\end{deluxetable}

%\begin{figure}
%\label{dameplots}
%\plottwo{St34.png}{LkCa19.png}
%\plottwo{RXJ043281735.png}{blank.png}
%\caption{{\it Top}: St~34, LkCa~19, and {\it Bottom}: RX~J0432.8+1735, from \citet{dame87}. Greyscale shading with a logarithmic stretch corresponds to CO(1-0) column densities. A red arrow indicates the positions of our sources in the center of each frame, and magenta crosshairs indicate our ``off cloud" position, a position with minimal CO emission from the molecular cloud which is subtracted from the spectrum of our source to eliminate atmospheric contamination.}
%\end{figure}

\subsubsection{St~34}
St~34 (HBC~425) is a binary system of two CTTSs, separated by $\lesssim0.78$~AU \citep{white05} based on the orbital solution for the system \citep{downes88} in the Taurus-Auriga T association \citep{kenyon95}. Both components of the spectroscopic binary have roughly equal mass and spectral types of M3 \citep{white05}. \citet{white05} observed St~34 in the optical with the HIRES spectrograph at Keck and derived an isochronal age of $8\pm3$~Myr for both components of the binary. Since they did not detect any lithium ($^7$Li) in the spectrum, St~34 must have reached---assuming the stars are completely convective---an internal temperature $>2\times10^6$~K and since depleted all of its lithium. St~34 has a low accretion rate of $2.5\times10^{-10}$~M$_\sun$~yr$^{-1}$, and the maximum radial velocity difference between the two binary components of St~34 is 58.4~km~s$^{-1}$ \citep{white05}. St~34, being one of the oldest known pre--main-sequence (PMS) star still accreting from a proto-planetary disk, also has a low dust mass of $\sim 2\times10^{-10}$~M$_\sun$ for radii $\lesssim 0.7$~AU \citep{hartmann05}.

\subsubsection{RX~J0432.8+1735}
% Need dust/debris mass citation here (Meyer et al. (2008) didn't have this source)
RX~J0432.8+1735 is a WTTS of spectral type M2 \citep{martin99}. Based on the PMS tracks of \citet{dantona97} RX~J0432.8+1735 is estimated to be $1.0\pm0.5$~Myr old. \citet{padgett06} observed RX~J0432.8+1735 with {\it Spitzer} and noticed that its $24\micron$ flux is in excess of the expected photospheric value by a factor of 3. Its lack of IR excess $\le 12\micron$ suggests there may be a large inner hole in the disk. Based on {\it ROSAT} observations, \citet{carkner96} discovered that RX~J0432.8+1735 is also an X-ray source. As RX~J0432.8+1735 is classified as a WTTS star with no estimates of its accretion rate, we assume it is not accreting.

\subsection{Observing Procedure}
On 26-27 November 2007, we observed the $^{12}$CO~$J=2-1$ (230.53799~GHz) emission line of our two T~Tauri stars with the 10~m Heinrich Hertz Sub-Millimeter Telescope (HHT) on Mt.~Graham, Arizona. Observations were obtained with a 1~mm dual polarization (Vpol, Hpol), sideband-separating, ALMA prototype receiver. The upper sideband was tuned to $^{12}$CO~$J=2-1$ while the lower sideband was tuned to $^{13}$CO~$J=2-1$. We used the Forbes Filter Bank (FFB) backend in 4 IF mode, an upper and lower sideband each with 1~MHz and 250~kHz of spectral resolution, respectively. The channel width, $\Delta\nu_\mathrm{ch}$, of our spectrometer was $0.33$~km~s$^{-1}$. The 1~MHz resolution data were used to determine main beam efficiencies $\eta_\mathrm{mb}$, and the $250$~kHz resolution data were used to measure the $^{12}$CO line.



\begin{deluxetable}{lccc}
\tabletypesize{\scriptsize}
\tablewidth{0pt}
\tablecaption{Observation of Main Beam Efficiencies $\eta$\label{etatab}}
\tablehead{\colhead{Planet} & \colhead{$\eta_\mathrm{Vpol}$} & \colhead{$\eta_\mathrm{Hpol}$} & \colhead{$\eta_\mathrm{Vpol}/\eta_\mathrm{Hpol}$}}
\startdata
%Uranus & $0.75\pm0.07$ & $0.73\pm0.09$ & $1.03\pm0.16$\\
%Uranus & $0.73\pm0.06$ & $0.64\pm0.08$ & $1.14\pm0.17$\\
Mars\tablenotemark{a} & $0.80\pm0.05$ & $0.72\pm0.04$ & $1.10\pm0.09$\\
Mars & $0.80\pm0.05$ & $0.71\pm0.04$ & $1.12\pm0.09$\\
Saturn & $0.93\pm0.04$ & $0.70\pm0.03$ & $1.33\pm0.08$\\
Saturn & $0.94\pm0.04$ & $0.69\pm0.03$ & $1.36\pm0.09$\\
Venus & $0.89\pm0.03$ & $0.66\pm0.03$ & $1.35\pm0.08$\\
Venus & $0.88\pm0.03$ & $0.66\pm0.03$ & $1.34\pm0.08$\\
Mars & $0.87\pm0.05$ & $0.66\pm0.04$ & $1.31\pm0.10$\\
Mars\tablenotemark{b} & $0.87\pm0.05$ & $0.66\pm0.04$ & $1.33\pm0.10$\\
Mars\tablenotemark{c} & $0.88\pm0.05$ & $0.68\pm0.04$ & $1.29\pm0.10$\\
Mars & $0.89\pm0.05$ & $0.68\pm0.04$ & $1.31\pm0.11$\\
Mars & $0.84\pm0.04$ & $0.70\pm0.04$ & $1.21\pm0.10$\\
Mars & $0.88\pm0.05$ & $0.68\pm0.04$ & $1.29\pm0.10$\\
Venus & $0.93\pm0.03$ & $0.68\pm0.03$ & $1.38\pm0.09$
\enddata
\tablenotetext{a}{All Mars brightness temperature errors assumed to be 5\%}
\tablenotetext{b}{End of first night}
\tablenotetext{c}{Beginning of second night}
\end{deluxetable}

Using CLASS in the GILDAS data reduction package, we estimated the main beam efficiencies by observing the planets shown in Table~\ref{etatab}. Typical sideband rejections, ignored in the calibration, were $>10$~dB. The main beam efficiency $\eta$ was computed following \citet{mangum93} and corrected for single-sideband observations:
\begin{equation}
\label{etaequ}
\eta=\frac{T_A^{\star}(\mathrm{planet})}{J(\nu_s,T_\mathrm{planet})-J(\nu_s,T_\mathrm{cmb})}\times \left[1-\exp{\left(-\ln{(2)}\frac{\theta_\mathrm{eq}\theta_\mathrm{pol}}{\theta_\mathrm{mb}^2}\right)}\right]^{-1},
\end{equation}
where
\begin{equation}
\label{Jequ}
J(\nu,T_b)=\frac{h\nu/k}{e^\frac{h\nu}{k T_\mathrm{mb}}-1}
\end{equation}
is the Planck function at brightness temperature $T_b$ and frequency $\nu$, $T_A^\star$ is the single-sideband antenna temperature of the planet, $T_\mathrm{planet}$ is the planet's observed brightness temperature, $T_\mathrm{cmb} = 2.73$~K, $\theta_\mathrm{eq}$ and $\theta_\mathrm{pol}$ are respectively the planet's equatorial and poloidal diameters in arcseconds, and $\theta_\mathrm{mb}=33\arcsec$ at $\nu=230$~GHz. We adopted an average Venus brightness temperature $T_b$ from \citet{kuznetsov82} of $287\pm20$~K. For all other planets' $T_b$, we used the JCMT online database\footnote{\url{http://www.jach.hawaii.edu/jac-bin/planetflux.pl}}. We derived a ratio $\eta_\mathrm{Vpol}/\eta_\mathrm{Hpol}$ of the two IF's mean main beam efficiencies for both nights of $1.24\pm0.04$. We used this ratio to scale the Hpol polarization's antenna temperature up to match the level of the Vpol polarization's antenna temperature. After fitting a baseline to each spectrum, we averaged the sum of the scaled Hpol brightness temperatures and the Vpol brightness temperatures: $\frac{1}{2}\langle T_A^\star(\mathrm{Hpol, scaled})+T_A^\star(\mathrm{Vpol})\rangle=T_A^\star(\mathrm{sum})$. Thus we computed the corrected main beam temperature as
\begin{equation}
\label{tmb_equ}
T_\mathrm{mb}=\frac{T_A^\star(\mathrm{sum})}{\eta_\mathrm{Vpol}}.
\end{equation}
The average beam efficiencies were $\langle\eta_\mathrm{Hpol}\rangle=0.68\pm0.01$ and $\langle\eta_\mathrm{Vpol}\rangle=0.88\pm0.01$ for both nights.

Since we had null detections for our two sources, we must assume a line-width to calculate upper limits on the integrated intensity. We assumed typcial a line-width of $\Delta\nu=10$~km~s$^{-1}$ ($=7.69$~MHz). If we assume the CO line is well described by a Gaussian line shape, then the uncertainty in the integrated intensity is given by
\begin{equation}
\label{sigmai}
\sigma_I = \sigma_{T_\mathrm{mb}}\sqrt{\frac{3\Delta\nu_\mathrm{ch}\Delta\nu}{\sqrt{\ln{2}}}},
\end{equation}
where $\sigma_I$ and $\Delta v$ are the CO line fluxes and the the full width at half maximum (FWHM) and $\Delta v_\mathrm{ch}$ is the channel spacing 0.33~km~s$^{-1}$; see Appendix I of \citet{schlingman}. The observations are summarized in Table~\ref{obssummary}.

\begin{deluxetable}{lcccccc}
\tabletypesize{\scriptsize}
\tablewidth{0pt}
\rotate
\tablecaption{Observational Summary\label{obssummary}}
\tablehead{\colhead{Source} & \colhead{Integration Time} & \colhead{$\sigma_{T_\mathrm{mb}}$} & \colhead{$\sigma_I$} & \colhead{$\log_{10}{(F_{\nu})}$} & \colhead{$\overline{N}(T\mathrm{ex}=10,20,100\mathrm{~K})$} & \colhead{$M_{\mathrm{H}_2}(T_\mathrm{ex}=10,20,100\mathrm{~K})$} \\
\colhead{} & \colhead{(sec)} & \colhead{(K)} & \colhead{(K~km~s$^{-1}$)} & \colhead{($\log_{10}{(\mathrm{W~cm}^{-2})}$)} & \colhead{(10$^{13}$~cm$^{-2}$)} & \colhead{(M$_\earth$)}}
\startdata
RX~J0432.8+1735 & 3600 & 0.019 & 0.046 & $< -23.02$ & $<12.7,7.47,6.53$ & $<3.59,1.10,1.84$\\
St~34  & 6120 & 0.019 & 0.046 & $< -23.02$ & $<12.7,7.47,6.53$ & $<3.59,1.10,1.84$
\enddata
\tablecomments{$T_\mathrm{mb}$ is the main beam corrected brightness temperature and $I$ is the corresponding intensity assuming a line width of 5~km~s$^{-1}$. $F$ is the $3\sigma$ line flux upper limit.}
\tablecomments{That $\sigma_{T_\mathrm{mb}}$ for both objects is the same is a fluke.}
\end{deluxetable}

\section{Results \& Analysis}
The main-beam corrected spectra of our observations are shown in Figure~\ref{spec}.

\begin{figure}
\begin{center}
\plottwo{st34.eps}{rxj0432.eps}
\end{center}
\caption{Spectra of St~34 ({\it left}) and RX~J0432.8+1735 ({\it right}). The vertical line labeled ``Cloud" represents the estimated background cloud velocity determined by the \citet{dame87} CO~$J=1-0$ survey, and the vertical line labeled with the source's name indicates that source's $V_\mathrm{LSR}$. For St~34, we show with an arrow the difference between the radial velocities of its binary components with an arrow centered on the systemic velocity.}
\label{spec}
\end{figure}

While 9$\arcmin$ is a rough spatial scale for comparison to the 33$\arcsec$ beam of the SMT, and since we did not detect a $^{12}$CO line in any of our sources, the on-cloud results from SMT are consistent with \citet{dame87}. It is unlikely that high spatial frequency variations of 33$\arcsec$ scales over 9$\arcmin$ regions have systemic velocity shifts of 2-3~km$^{-1}$ %\citep{skrutskie93} <-- DOES SKRUTSKIE ET AL. (1993) SUPPORT THIS CLAIM?

Since we did not detect any $^{12}$CO line, we convert our $3\sigma$ noise into an upper limit on the flux. A knowledge of the flux will enable us to estimate upper limits on gas disk mass.

%\subsection{Converting $\Delta T$ to $F$}
%To convert main beam corrected, 3~$\sigma$ brightness temperature noises ($3\sigma_{T_\mathrm{mb}}$) into line flux upper limits, we first convert our brightness temperatures $T_b$ into specific intensities $I_\nu$ according to the definition of $T_b$ from the Rayleigh-Jeans law:
%\begin{equation}
%\label{T_b def}
%T_b\equiv \frac{c^2}{2 k\nu^2}I_\nu.
%\end{equation}
The observed flux is the double-integral of the observed intensity $I_\nu$ over frequency $\nu$ and solid angle $\Omega$:
\begin{equation}
\label{double integral}
F=\int\int I_\nu\,d\nu\,d\Omega=\int\int\frac{2 k \nu^3 T_b}{c^3}\,dv\,d\Omega.
\end{equation}
%Assuming the beam is filled according to a $1 \sigma$ Gaussian,
%\begin{equation}
%\label{Omega}
%\Omega=\frac{\pi\theta^2}{4\ln{2}},
%\end{equation}
%where $\ln{2}$ accounts for the difference between a Gaussian-filled beam and a uniformly-filled beam. Now for the $\nu$ integral, we change to velocity coordinates according to the Doppler relation
%\begin{equation}
%\label{Doppler}
%d\nu=\frac{\nu}{c}dv
%\end{equation}
%and integrate over our line, assumed to be Gaussian with FWHM of $\Delta v$. The difference between a Gaussian line of FWHM of $\Delta v$ and a top-hat function of width $\Delta v$ is $\sqrt{4\ln{2}/\pi}\approx0.939$, thus integrating over the Gaussian line yields
%\begin{equation}
%\label{v integral}
%\int I_\nu\,d\nu=\int\frac{2 k \nu^3 T_b}{c^3}\,dv=\frac{2 k \nu^3 T_b}{c^3}(0.939\times\Delta v).
%\end{equation}
Assuming the brightness temperature does not vary substantially over the telescope beam and that the line is a Gaussian with line-width $\Delta v$, then the upper limit on the $3\sigma$ line flux $F$ is
\begin{equation}
\label{upper lim}
F < \frac{2 k \nu^3 (3\sigma_{T_\mathrm{mb}})}{c^3}\frac{\pi\theta^2}{4\ln{2}}\sqrt{\frac{4\ln{2}}{\pi}}\times\Delta v=(5.05\times10^{-15})\sigma_{T_\mathrm{mb}}\mathrm{erg~s}^{-1}\mathrm{~cm}^{-2}.
\end{equation}
The $3\sigma$ upper limits on $F$ are listed in Table~\ref{obssummary}.

%\subsection{Estimating Disk Mass Upper Limits}

%\subsubsection{Column Depth}
%To estimate the column density $N(T_\mathrm{ex})$ and eventually the mass of the circumstellar disks, we begin with the radiation temperature $T_R$, where $T_R=f T_\mathrm{mb}$ and $f$ is the beam filling factor. The radiation temperature $T_R$ is given by \citet{ulich76} as
%\begin{equation}
%\label{T_R equ}
%T_R=\big\{J_\nu(T_\mathrm{ex})-J_\nu(T_\mathrm{cmb})\}\big(1-e^{-\tau}\big),
%\end{equation}
%where $J_\nu$ is the Planck function in temperature units, $T_\mathrm{ex}$ is our line's excitation temperature, $T_\mathrm{cmb}=2.7$~K, and $\tau$ is optical depth. In local thermodynamic equilibrium (LTE), Kirchoff's law relates the emission coefficient $j_\nu$ to the absorption coefficient $\alpha_\nu$ as
%\begin{equation}
%\label{Kirchoff}
%j_\nu=B_\nu(T_\mathrm{ex})\alpha_\nu,
%\end{equation}
%i.e., the source function $S_\nu\equiv j_\nu/\alpha_\nu$ becomes the Planck function $B_\nu$.
% The Boltzmann law then relates the number of CO molecules in a given state $u$ to the total number of CO molecules by
%\begin{equation}
%\label{T_R equ}
%N_u=N\frac{g_u}{Q(T_\mathrm{ex})}\exp{\left(-\frac{E_u}{k T_\mathrm{ex}}\right)},
%\end{equation}
%where $N_u$ is the number of CO molecules in state $u$, $g_u=5$ is the statistical weight of the state $J_u=2$, $Q(T_\mathrm{ex})$ is the partition function at temperature $T_\mathrm{ex}$, $E_u$ is the energy of state $u$, and $k$ is the Boltzmann constant. Using the previous three equations, we derive the column density $\overline{N}(T_\mathrm{ex})$ in the optically thin limit to be
Similarly, we can derive the column density in the optically thin limit to be
\begin{equation}
\label{col depth equ}
\overline{N}(T_\mathrm{ex})=\frac{8\pi k\nu^2}{h c^3 g_u A_{ul}}\mathcal{F}(T_\mathrm{ex}, E_u, \nu)\int T_\mathrm{mb}\,dv,
\end{equation}
where
\begin{equation}
\label{F def}
\mathcal{F}(T_\mathrm{ex}, E_u, \nu)\equiv\frac{J_\nu(T_\mathrm{ex})Q(T_\mathrm{ex})\exp{\left(\frac{E_u}{k T_\mathrm{ex}}\right)}}{J_\nu(T_\mathrm{ex})-J_\nu(T_\mathrm{cmb})},
\end{equation}
$A_{ul}$ is the Einstein A coefficient (spontaneous emission) and has units of s$^{-1}$.
%To find the column density in the optically thick limit, we use the result from \citet{goldsmith99} who proved that the column density in the optically thick limit is a factor of $C_\tau\equiv\tau_\mathrm{tot}/(1-e^{-\tau_\mathrm{tot}})$ larger than the column density in the optically thin limit.

Similar to the analysis of \citet{pascucci06}, we assume an excitation temperature $T_\mathrm{ex}\approx20$~K. Then in our case for $^{12}$CO~$J=2-1$, $\mathcal{F}(T_\mathrm{ex}, E_u, \nu)\approx28.00$. The Einstein $A_{ul}=6.91\times10^{-7}$~s$^{-1}$ and partition function $Q(20\mathrm{~K})=15.9$ (CDMS\footnote{\url{http://www.ph1.uni-koeln.de/vorhersagen/}}).
%The Einstein $A_{ul}$ for a linear dipole molecule in a $^1\Sigma$ electronic state is
%\begin{equation}
%\label{F def}
%A_{ul}=\frac{64\pi^4 \nu^3}{3 h c^3}\mu^2\frac{J}{2J+1}=6.91\times10^{-7}\mathrm{~s}^{-1},
%\end{equation}
%where $\mu=0.10908$~Debye is the electric dipole moment for $^{12}$CO.

%\subsubsection{Gas Masses in the Optically Thin and Thick Limits}
%Assuming $\int T_R\,dv=0.939 \sigma_{T_\mathrm{mb}}\times 10^5$~cm~s$^{-1}$, where 0.939 accounts for a Gaussian line profile and $10^5$~cm~s$^{-1}$ is the line width,
We compute and tabulate in Table~\ref{obssummary} the $^{12}$CO number densities and H$_2$ gas masses in the optically thin limit. Gas disk masses were derived from \citet{scoville86},
\begin{equation}
\label{scoville mass eqn}
M_{\mathrm{H}_2}<\overline{N}(T_\mathrm{ex})\times\Bigg\{\left[\frac{\mathrm{H}_2}{\mathrm{CO}}\right]\mu_G m_{\mathrm{H}_2} \frac{\pi\theta^2}{4} d^2\Bigg\}\mathrm{~g}=110\sigma_{T_\mathrm{mb}}\mathrm{~M}_\earth,
\end{equation}
where $\left[\mathrm{H}_2/\mathrm{CO}\right]\approx10^4$, $\mu_G=1.36$ is the mean molecular weight, $m_{\mathrm{H}_2}$ is the mass of an H$_2$ molecule, and $d\approx140$~pc is the distance to Taurus.

%Assuming the gas disk is optically think in a tiny, but beam diluted,

%The quantity $(2.90\times10^{-15})\sigma_{T_\mathrm{mb}}$ comes from Equation~\ref{upper lim}.

%- [ ] How does gas content evolve over time?
% - [ ] What's the oldest system known with gas?
% - [ ] How does this compare to our sources' ages?
%- [ ] How does the gas content relate to dust content?
%- [ ] How do our null detections fit in to this?

\section{Discussion}
%How does the gas disk mass evolve over time?

%Older
%\citet{briceno01} find that, for 168 main sequence stars in Orion with masses between 0.3 and 0.6~$M_\sun$, equivalent widths of H$\alpha$ emission indicate that accretion ceases after a few Myr. \citet{eisner03} corroborate this result with their observations of the $\sim0.3$~Myr cluster NGC~2024. Indeed \citet{durisen07} argues that gravitational instabilities in gas disks account for their rapid ($<3$~Myr) dissipation, whereas simulations of the evolution of gas disk surface density for a 1~M$_\sun$ star due to photo-evaporation indicate that gas disks disappear in about 6~Myr \citep{alexander06,dullemond07}.

%If gas content and dust content follow similar timescales \citep{damajanov07}, then it would be surprising for our sources to have a non-zero gas disk content if they do indeed have dust disks as identified by {\it Spitzer}.

%St~34, being $8\pm3$~Myr old CTTS, indicates that it could have evaporated its disk or stored up its gas in planets much later than $\sim 3$~Myr. Being a CTTS, St~34 must still be accreting gas, but because of our null detection of its gas disk, most of its outer disk must have disappeared. Conversely, RX~J0432.8+1735, being a $1.0\pm0.5$~Myr old WTTS, must have either evaporated its disk or formed planets from its gas disk faster than normal. Its being a WTTS is consistent with our null detection of its gas content, although remnant amounts of gas less than a few $M_\earth$ could still exist in its outer disk.

%Newer
%Theories of gaseous planet formation require knowledge of 1) how much gas content there is in circumstellar disks and 2) how long this gas is available for planet formation. To answer the first question, upper limits on the amount of gas in circumstellar disks puts limits directly on the amount of mass available for gaseous planet formation. The second question is also necessary because if the timescale for planet formation is much longer than gas disk dispersion timescales, gaseous planets could not form in that particular system. Our results attempt to answer both questions; they put upper limits on the gas mass. Because we know the ages for most of our sources and can estimate $\langle\dot{M}\rangle$, the average accretion rate, we can determine how long and how much gas has been available for planet formation.

%According to \citet{meyer07}, inner circumstellar disk lifetimes are $\sim3$~Myr with a real $1-10$~Myr dispersion. This relatively large dispersion suggests that there are a range of initial conditions for the planet formation process. Thus we do not {\it a priori} expect that just because our systems have low gas disk content, they have not possibly already formed planets, large reservoirs of gas that does not readily contribute to the $^{12}$CO emission and thus possibly make our gas disk mass upper limits appear too small.

%Accretion rates, measured by H$\alpha$ equivalent widths, also help one to understand gas disk timescales. \citet{briceno01} find that, for 168 main sequence stars in Orion with masses between 0.3~$M_\sun$ and 0.6~$M_\sun$, equivalent widths of H$\alpha$ emission indicate that accretion ceases after a few Myr, and \citet{eisner03} corroborate this result with their observations of the $\sim0.3$~Myr cluster NGC~2024. It would seem unusual, then, for a CTTS, which has ongoing accretion, to be older than a few Myr, but perhaps some older CTTSs accrete the same amount of mass but at a lower accretion rate. A few Myr is a similar timescale to the $\sim3$~Myr for inner disk dispersion, yet there are older CTTSs that still have gas disks, like the CTTS St~34 in our sample. \citet{durisen07}, however, argue that gravitational instabilities in gas disks account for their rapid ($<3$~Myr) dissipation, whereas simulations of the evolution of gas disk surface density for a 1~M$_\sun$ star due to photo-evaporation indicate that gas disks disappear in about 6~Myr \citep{alexander06,dullemond07}. Thus again the large dispersion in gas disk dispersion timescales does not contradict our results.

%Mike's outline
%par #1 - there are a few ways in which gas mass observations
%constrain theories of planet formation and evolution, including
%amount of material available initially (w.r.t. classical core
%accretion as well as gravitational fragmentation), timescale
%for both these processes of forming gas giants, the formation
%of ice giants at larger radii, and even the dynamical evolution
%of forming terrestrial planets (see discussion in Pascucci et al.
%2006).

Theories of gaseous planet formation require knowledge of (1) how much gas there is in circumstellar disks initially and (2) the rate at which gas is depleted over time. To answer the first question, upper limits on the amount of gas in very young protostellar disk systems puts limits directly on the amount of mass available for gaseous planet formation in the systems observed. Answers to the second question, requiring large samples over a wide range of ages, are also necessary because of the competition between the timescale for planet formation and gas disk dispersion timescales \citep[for a recent review, see][]{meyer09}.

%par #2 - compare and contrast various ways of constraining
%gas masses (UV absorption lines, near-IR flurescent H2,
%mid-IR rovibrational lines, mm-wave lines of trace molecules
%with permanent electric dipole moments).

There are several ways of constraining gas disk masses, each with its own advantages and disadvantages. To understand gas disk timescales, one can also analyze H$\alpha$ emission line profiles and determine gas accretion rates
and thus constrain the gas mass surface density at the inner edge of the disk tracing perhaps global disk evolution \citep[e.g.][]{fedele09}. Using UV tracers of gas emission, \citet{ingleby09} describe {\it HST} observations searching 
for evidence of hot gas in emission finding no evidence for H$_2$ emission for WTTS in their sample. Ultraviolet absorption line from a continuum source \citep[e.g.][]{roberge05} can help constrain cold mass in disks, but it requires a continuum source that is bright in the far UV and an edge-on geometry; therefore, it is observationally feasible only in special circumstances. 
Near IR fluorescent H$_2$ traces gas with excitation temperatures $T_\mathrm{ex}>2000$~K \citep{bary03} and mid-IR ro-vibrational lines \citep[e.g.][]{najita07,thi01,pascucci06}, especially the $28.2~\mu$m and $17\mu$m {\it Spitzer} bands, trace gas up to 50-200~K. Our observations of rotational lines of CO trace cooler gas at larger orbital radii. The disadvantage of such measurements is that CO can freeze out at the coldest temperatures corresponding to the outer limits of the disk $\gtrsim 30$-100~AU. \citet{hughes10} present evidence for evolving gas to dust ratios in transitional disks which exhibit evidence for optically-thick outer disks but possess inner holes and gaps.
 

%\citet{briceno01} find that, for 168 main sequence stars in Orion with masses between 0.3~$M_\sun$ and 0.6~$M_\sun$, analysis of H$\alpha$ emission indicates that accretion ceases after a few Myr. A few Myr is a similar timescale to the $\sim3$~Myr for inner disk dissipation, yet there are older CTTSs that still have gas disks, like the CTTS St~34 in our sample. \citet{durisen07}, however, argue that gravitational instabilities in gas disks account for their rapid ($<3$~Myr) dissipation, whereas simulations of the evolution of gas disk surface density for a 1~M$_\sun$ star due to photo-evaporation indicate that gas disks disappear in about 6~Myr \citep{alexander06,dullemond07}. Core accretion models assert that the planet formation timescales $\tau_\mathrm{form}$ are of the same order as the disk depletion timescales $\tau_\mathrm{dep}$ \citep{ida08}. According to \citet{meyer07}, observed inner circumstellar disk lifetimes are $\sim3$~Myr with a real $1-10$~Myr dispersion. This relatively large dispersion suggests that there are a range of initial conditions for the planet formation process.

%par #3 - compare these results in the context of other
%observations of disk evolution

St~34, being $8\pm3$~Myr old, might have a lower gas surface density than typical CTTSs. It is still accreting gas \citep{white05}, albeit at a low rate, but retains at least a low density inner disk. Our non--detection in a search for cold gas implies that most of its outer disk must have disappeared or frozen out onto grains. \citet{white05} show that St~34 has an accretion rate $\langle\dot{M}\rangle=8.3\times10^{-10}$~M$_\sun$~yr$^{-1}$, so after 1~Myr much more gas than our upper limit of $4.20$~M$_\earth$ would accrete ($\sim276$~M$_\earth$). Perhaps St~34 recently lost its outer disk through photoevaporation \citep[e.g.][]{gorti09} and we are witnessing the ``last gasp'' of accretion onto the star. This is possible but not likely, as it requires current observations to be taking place at a very special time.

Conversely, RX~J0432.8+1735---being a much younger, $1.0\pm0.5$~Myr WTTS---must have either evaporated its disk or formed planets from its gas disk faster than normal. Its being a WTTS is consistent with our null detection of its gas content, although remnant amounts of gas less than a few $M_\earth$ could still exist in its outer disk. %???"Our null results and the $70\micron$ detection of \citet{padgett06} implies a gas-to-dust ratio of $<XXX$."??? Padgett et al. (2006) had no detection of RXJ... @ 70�m.
That RX~J0432.8+1735 is relatively young and does not have detectable gas content could be significant considering that there are older systems, such as Hen 3-600, a binary system at between $1-10$~Myr of age with apparent WTTS and CTTS components \citep{jayawardhana99}; TW~Hya, a CTTS at 8-10~Myr \citep{webb99}; and DM~Tau at $\sim 8$~Myr \citep{guilloteau94}. Thus something about the disk evaporation physics is dramatically different in RX~J0432.8+1735 than these oldest systems.

%and the associated Table~\ref{figuretable}
In Figure~\ref{mgdisk vs age} we compare our sources' gas disk masses and ages with the gas mass upper limits of those sources $\le 30$~Myr from \citet{pascucci06} and with the gas mass determinations ({\it solid circles}) of BP~Tau \citep[$^{13}$CO~$J=2-1$;][]{dutrey03}; DL~Tau, DO~Tau \citep[$^{12}$CO~$J=2-1$;][]{koerner95}; and DM~Tau, DR~Tau, GG~Tau~a, GM~Aur, GO~Tau, LkCa~15, RY~Tau \citep[$^{12}$CO~$J=3-2$ and $^{13}$CO~$J=3-2$;][]{thi01}. We note that this is not an exhaustive compilation from the literature, but representative of recent results. Assuming a 1:10 gas-to-dust ratio \citep{dalessio05}, we would expect RX~J0432.8+1735 and St~34 to have at most 0.420~M$_\earth$ and 0.420~M$_\earth$ of dust, respectively, for $T_\mathrm{ex}=20$~K. For St~34, \citet{hartmann05} estimates a disk mass of 665~M$_\earth$ located in a circumbinary disks between the ``wall" (the region defined to surround the two components of the St~34 binary; $\sim 0.7$~AU) and $7$~AU. Assuming this mass is representative of a total disk mass in St~34 out to $7$~AU, this would be consistent with a small ($<10$~AU), optically-thick CO disk with a gas-to-dust ratio of $\sim100$; we do not detect this due to beam dilation \citep[cf.][]{pascucci06}. %This is very high for a T~Tauri object \citep[cf. Figure~13 of][]{thi01} and implies significant evolution from the ISM value of 100. 

%\begin{deluxetable}{lccccc}
%\rotate
%\tabletypesize{\scriptsize}
%\tablewidth{0pt}
%\tablecaption{Data of Figure~\ref{mgdisk vs age}\label{figuretable}}
%\tablehead{\colhead{Object} & \colhead{Stellar Age} & \colhead{$M_\mathrm{gas}$ Upper Limit} & \colhead{Method} & \colhead{Stellar Mass} & \colhead{References} \\ \colhead{} & \colhead{(Myr)} & \colhead{} & \colhead{(M$_\sun$)} & \colhead{}}
%\startdata
%\enddata
%\tablecomments{{\bf I am currently working on formatting the data to insert here.}}
%\end{deluxetable}

% show how our gas disk upper limits compare to the upper limits of \citet{pascucci06} and how they compare to the few disks with well-determined gas mass. We see that our two sources appear to break the trend that the younger the system, the more massive its gas disk is. Thus the question remains: What physical parameters in our systems, especially in the young RX~J0432.8+1735, are different compared to the others that do not have such a rapid disk dissipation?

%par #4 - point out any implications of these observations
%(if there are any) based on the above discussion.

St~34 has infrared excess for wavelengths longer than 3.6~$\mu$m \citep{hartmann05} and RX~J0432.8+1735 has infrared excess for wavelengths longer than 24~$\mu$m \citep{padgett06}, yet St~34 is a binary CTTS with an accreting inner disk and RX~J0432.8+1735 is a WTTS. Binaries tend to disrupt inner gas disks \citep{jensen96} and may decrease disk lifetimes \citep{monin07}. However, \citet{armitage96} have argued that close binaries affect angular momentum exchange in the natural evolution of accretion disks resulting in longer lived outer disks. Indeed \citet{thebault04} find that planet formation around binaries might require a long-lived but massive disk. Since circumbinary disks allow for long gas disk lifetimes, St~34 might have had more time to form planets.

%Since all the points in Figure~\ref{mgdisk vs age} such that the source moves at different velocity than the background cloud, that it is bright in CO, and/or that it is close enough to be resolved spatially by an interferometer array; the apparent trend that M$_\mathrm{gas}$ decreases with age is physical. Thus there remain two questions: ``What physical parameters in our systems, especially in the young RX~J0432.8+1735, are different compared to the others that do not have such a rapid disk depletion?" or, ``What is different about St~34 that makes its gas disk depletion timescale so long?"

\begin{figure}
\begin{center}
\plotone{mgasdisk_vs_age.eps}
\end{center}
\caption{Gas circumstellar disk mass versus age of selected sources: our RX~J0432.8+1735 and St~34 upper limits ({\it labeled}); upper limits from the \citet{pascucci06} sample ({\it unlabeled upper limits}) with ages $\le 30$~Myr (the Kelvin-Hemholtz contraction timescale for a 1~M$_\sun$ star); and exact mass determinations ({\it solid circles}) of BP~Tau \citep[$^{13}$CO~$J=2-1$;][]{dutrey03}; DL~Tau, DO~Tau \citep[$^{12}$CO~$J=2-1$;][]{koerner95}; and DM~Tau, DR~Tau, GG~Tau~a, GM~Aur, GO~Tau, LkCa~15, RY~Tau \citep[$^{12}$CO~$J=3-2$ and $^{13}$CO~$J=3-2$;][]{thi01}. We assume errors in stellar ages to be 50\%.}
\label{mgdisk vs age}
\end{figure}

\section{Conclusions}
%Our main conclusions are:
%\begin{itemize}
% \item We do not detect significant amounts of gas around three T~Tauri stars.
% \item St~34, a CTTS, is still accreting gas although it is $8\pm 3$~Myr old.
% \item The gas disk of RX~J0432.8+1735, a WTTS of $1.0\pm0.5$~Myr, has disappeared quickly.
%\end{itemize}

Assuming optically thin disks ($\tau\ll1$), an excitation temperature $T_\mathrm{ex}=20$~K, and a line-width $\Delta v=10$~km~s$^{-1}$, we do not detect significant amounts of gas around three T~Tauri stars: $<4.20$~M$_\earth$ for the PMS binary St~34 and $<4.20$~M$_\earth$ for RX~J0432.8+1735. St~34, a CTTS, is still accreting gas although it is $8\pm 3$~Myr old, and the gas disk of RX~J0432.8+1735, a WTTS of $1.0\pm0.5$~Myr, has disappeared quickly. 
Future observations of larger samples will be required to understand the diversity of disk lifetimes as a function of stellar properties. 

\begin{thebibliography}{}
\bibitem[Alexander et al.(2006)]{alexander06} Alexander, R.~D., Clarke, C.~J., \& Pringle, J.~E.\ 2006, \mnras, 369, 229 
\bibitem[Andrews et al.(2009)]{andrews09} Andrews, S.~M., Wilner, D.~J., Hughes, A.~M., Qi, C., \& Dullemond, C.~P.\ 2009, \apj, 700, 1502 
\bibitem[Armitage \& Clarke(1996)]{armitage96} Armitage, P.~J., \& Clarke, C.~J.\ 1996, \mnras, 280, 458 
\bibitem[Bary et al.(2003)]{bary03} Bary, J.~S., Weintraub, D.~A., \& Kastner, J.~H.\ 2003, \apj, 586, 1136
%\bibitem[Brice{\~n}o et al.(2001)]{briceno01} Brice{\~n}o, C., et al.\ 2001, Science, 291, 93 
%\bibitem[Calvet et al.(2004)]{calvet04} Calvet, N., Muzerolle, J., Brice{\~n}o, C., Hern{\'a}ndez, J., Hartmann, L., Saucedo, J.~L., \& Gordon, K.~D.\ 2004, \aj, 128, 1294 
\bibitem[Carkner et al.(1996)]{carkner96} Carkner, L., Feigelson, E.~D., Koyama, K., Montmerle, T., \& Reid, I.~N.\ 1996, \apj, 464, 286 
\bibitem[D'Alessio et al.(2005)]{dalessio05} D'Alessio, P., Mer{\'{\i}}n, B., Calvet, N., Hartmann, L., \& Montesinos, B.\ 2005, Revista Mexicana de Astronomia y Astrofisica, 41, 61 
\bibitem[Dame et al.(1987)]{dame87} Dame, T.~M., et al.\ 1987, \apj, 322, 706 
\bibitem[Damjanov et al.(2007)]{damajanov07} Damjanov, I., Jayawardhana, R., Scholz, A., Ahmic, M., Nguyen, D.~C., Brandeker, A., \& van Kerkwijk, M.~H.\ 2007, \apj, 670, 1337 
\bibitem[D'Antona \& Mazzitelli(1997)]{dantona97} D'Antona, F., \& Mazzitelli, I.\ 1997, Memorie della Societa Astronomica Italiana, 68, 807
\bibitem[Downes \& Keyes(1988)]{downes88} Downes, R.~A., \& Keyes, C.~D.\ 1988, \aj, 96, 777 
\bibitem[Dullemond et al.(2007)]{dullemond07} Dullemond, C.~P., Hollenbach, D., Kamp, I., \& D'Alessio, P.\ 2007, Protostars and Planets V, 555 
\bibitem[Durisen et al.(2007)]{durisen07} Durisen, R.~H., Boss, A.~P., Mayer, L., Nelson, A.~F., Quinn, T., \& Rice, W.~K.~M.\ 2007, Protostars and Planets V, 607 
\bibitem[Dutrey et al.(2003)]{dutrey03} Dutrey, A., Guilloteau, S., \& Simon, M.\ 2003, \aap, 402, 1003 
\bibitem[Dutrey et al.(2007)]{dutrey07} Dutrey, A., Guilloteau, S., \& Ho, P.\ 2007, Protostars and Planets V, 495
%\bibitem[Eisner \& Carpenter(2003)]{eisner03} Eisner, J.~A., \& Carpenter, J.~M.\ 2003, \apj, 598, 1341
\bibitem[Evans et al.(2003)]{evans03} Evans, N.~J., II, et al.\ 2003, \pasp, 115, 965
\bibitem[Fedele et al.(2009)]{fedele09} Fedele, D., van den Ancker, M.~E., Henning, T., Jayawardhana, R., \& Oliveira, J.~M.\ 2009, arXiv:0911.3320 
%\bibitem[Goldsmith \& Langer(1999)]{goldsmith99} Goldsmith, P.~F., \& Langer, W.~D.\ 1999, \apj, 517, 209 
\bibitem[Gorti \& Hollenbach(2009)]{gorti09} Gorti, U., \& Hollenbach, D.\ 2009, \apj, 690, 1539
\bibitem[Guilloteau \& Dutrey(1994)]{guilloteau94} Guilloteau, S., \& Dutrey, A.\ 1994, \aap, 291, L23
\bibitem[Hartmann et al.(2005)]{hartmann05} Hartmann, L., et al.\ 2005, \apjl, 628, L147 
%\bibitem[Herbig(1988)]{herbig88} Herbig, G.~H.\ 1988, \apj, 331, 999 
\bibitem[Hughes et al.(2010)]{hughes10} Hughes, A.~M., et al.\ 2010, arXiv:1007.3267 
%\bibitem[Holman \& Wiegert(1999)]{holman99} Holman, M.~J., \& Wiegert, P.~A.\ 1999, \aj, 117, 621
%\bibitem[Ida \& Lin(2008)]{ida08} Ida, S., \& Lin, D.~N.~C.\ 2008, \apj, 673, 487
\bibitem[Ingleby et al.(2009)]{ingleby09} Ingleby, L., et al.\ 2009, \apjl, 703, L137
\bibitem[Jayawardhana et al.(1999)]{jayawardhana99} Jayawardhana, R., Hartmann, L., Fazio, G., Fisher, R.~S., Telesco, C.~M., \& Pi{\~n}a, R.~K.\ 1999, \apjl, 520, L41 
\bibitem[Jensen(1996)]{jensen96} Jensen, E.~L.~N.\ 1996, Ph.D.~Thesis,
%\bibitem[Kastner et al.(1997)]{kastner97} Kastner, J.~H., Zuckerman, B., Weintraub, D.~A., \& Forveille, T.\ 1997, Science, 277, 67
\bibitem[Kenyon \& Hartmann(1995)]{kenyon95} Kenyon, S.~J., \& Hartmann, L.\ 1995, \apjs, 101, 117
\bibitem[Koerner \& Sargent(1995)]{koerner95} Koerner, D.~W., \& Sargent, A.~I.\ 1995, \aj, 109, 2138 
\bibitem[Kuznetsov et al.(1982)]{kuznetsov82} Kuznetsov, I.~V., Fedoseev, L.~I., \& Shvetsov, A.~A.\ 1982, Radiofizika, 25, 247
%\bibitem[Lin et al.(2006)]{lin06} Lin, S.-Y., Ohashi, N., Lim, J., Ho, P.~T.~P., Fukagawa, M., \& Tamura, M.\ 2006, \apj, 645, 1297
\bibitem[Mangum(1993)]{mangum93} Mangum, J.~G.\ 1993, \pasp, 105, 117
\bibitem[Mart{\'{\i}}n \& Magazz{\`u}(1999)]{martin99} Mart{\'{\i}}n, E.~L., \& Magazz{\`u}, A.\ 1999, \aap, 342, 173
\bibitem[Meyer et al.(2006)]{meyer06} Meyer, M.~R., et al.\ 2006, \pasp, 118, 1690 
\bibitem[Meyer et al.(2007)]{meyer07} Meyer, M.~R., Backman, D.~E., Weinberger, A.~J., \& Wyatt, M.~C.\ 2007, Protostars and Planets V, 573 
%\bibitem[Meyer et al.(2008)]{meyer08} Meyer, M.~R., et al.\ 2008, \apjl, 673, L181 
\bibitem[Meyer(2009)]{meyer09} Meyer, M.~R.\ 2009, IAU Symposium, 258, 111
\bibitem[Monin et al.(2007)]{monin07} Monin, J.-L., Clarke, C.~J., Prato, L., \& McCabe, C.\ 2007, Protostars and Planets V, 395
\bibitem[Najita et al.(2007)]{najita07} Najita, J.~R., Carr, J.~S., Glassgold, A.~E., \& Valenti, J.~A.\ 2007, Protostars and Planets V, 507 
\bibitem[Padgett et al.(2006)]{padgett06} Padgett, D.~L., et al.\ 2006, \apj, 645, 1283 
\bibitem[Pascucci et al.(2006)]{pascucci06} Pascucci, I., et al.\ 2006, \apj, 651, 1177 
\bibitem[Qi et al.(2004)]{qi04} Qi, C., et al.\ 2004, \apjl, 616, L11
\bibitem[Roberge et al.(2005)]{roberge05} Roberge, A., Weinberger, A.~J., Redfield, S., \& Feldman, P.~D.\ 2005, \apjl, 626, L105
\bibitem[Schlingman et al. (in prep.)]{schlingman} Schlingman et al., in prep.
\bibitem[Scoville et al.(1986)]{scoville86} Scoville, N.~Z., Sargent, A.~I., Sanders, D.~B., Claussen, M.~J., Masson, C.~R., Lo, K.~Y., \& Phillips, T.~G.\ 1986, \apj, 303, 416 
\bibitem[Silverstone et al.(2006)]{silverstone06} Silverstone, M.~D., et al.\ 2006, \apj, 639, 1138 
\bibitem[Skrutskie et al.(1993)]{strutskie93} Skrutskie, M.~F., et 
al.\ 1993, \apj, 409, 422 
\bibitem[Th{\'e}bault et al.(2004)]{thebault04} Th{\'e}bault, P., Marzari, F., Scholl, H., Turrini, D., \& Barbieri, M.\ 2004, \aap, 427, 1097
\bibitem[Thi et al.(2001)]{thi01} Thi, W.~F., et al.\ 2001, \apj, 561, 1074
%\bibitem[Ulich \& Haas(1976)]{ulich76} Ulich, B.~L., \& Haas, R.~W.\ 1976, \apjs, 30, 247 
\bibitem[van Kempen et al.(2007)]{vankempen07} van Kempen, T.~A., van Dishoeck, E.~F., Brinch, C., \& Hogerheijde, M.~R.\ 2007, \aap, 461, 983
\bibitem[Webb et al.(1999)]{webb99} Webb, R.~A., Zuckerman, B., Platais, I., Patience, J., White, R.~J., Schwartz, M.~J., \& McCarthy, C.\ 1999, \apjl, 512, L63 
\bibitem[White \& Hillenbrand(2005)]{white05} White, R.~J., \& Hillenbrand, L.~A.\ 2005, \apjl, 621, L65 
\bibitem[Wichmann et al.(2000)]{wichmann00} Wichmann, R., et al.\ 2000, \aap, 359, 181
\end{thebibliography}
\end{document}

1112.4581.extracted_bib
\begin{thebibliography}{}
\bibitem[Alexander et al.(2006)]{alexander06} Alexander, R.~D., Clarke, C.~J., \& Pringle, J.~E.\ 2006, \mnras, 369, 229 
\bibitem[Andrews et al.(2009)]{andrews09} Andrews, S.~M., Wilner, D.~J., Hughes, A.~M., Qi, C., \& Dullemond, C.~P.\ 2009, \apj, 700, 1502 
\bibitem[Armitage \& Clarke(1996)]{armitage96} Armitage, P.~J., \& Clarke, C.~J.\ 1996, \mnras, 280, 458 
\bibitem[Bary et al.(2003)]{bary03} Bary, J.~S., Weintraub, D.~A., \& Kastner, J.~H.\ 2003, \apj, 586, 1136
%\bibitem[Brice{\~n}o et al.(2001)]{briceno01} Brice{\~n}o, C., et al.\ 2001, Science, 291, 93 
%\bibitem[Calvet et al.(2004)]{calvet04} Calvet, N., Muzerolle, J., Brice{\~n}o, C., Hern{\'a}ndez, J., Hartmann, L., Saucedo, J.~L., \& Gordon, K.~D.\ 2004, \aj, 128, 1294 
\bibitem[Carkner et al.(1996)]{carkner96} Carkner, L., Feigelson, E.~D., Koyama, K., Montmerle, T., \& Reid, I.~N.\ 1996, \apj, 464, 286 
\bibitem[D'Alessio et al.(2005)]{dalessio05} D'Alessio, P., Mer{\'{\i}}n, B., Calvet, N., Hartmann, L., \& Montesinos, B.\ 2005, Revista Mexicana de Astronomia y Astrofisica, 41, 61 
\bibitem[Dame et al.(1987)]{dame87} Dame, T.~M., et al.\ 1987, \apj, 322, 706 
\bibitem[Damjanov et al.(2007)]{damajanov07} Damjanov, I., Jayawardhana, R., Scholz, A., Ahmic, M., Nguyen, D.~C., Brandeker, A., \& van Kerkwijk, M.~H.\ 2007, \apj, 670, 1337 
\bibitem[D'Antona \& Mazzitelli(1997)]{dantona97} D'Antona, F., \& Mazzitelli, I.\ 1997, Memorie della Societa Astronomica Italiana, 68, 807
\bibitem[Downes \& Keyes(1988)]{downes88} Downes, R.~A., \& Keyes, C.~D.\ 1988, \aj, 96, 777 
\bibitem[Dullemond et al.(2007)]{dullemond07} Dullemond, C.~P., Hollenbach, D., Kamp, I., \& D'Alessio, P.\ 2007, Protostars and Planets V, 555 
\bibitem[Durisen et al.(2007)]{durisen07} Durisen, R.~H., Boss, A.~P., Mayer, L., Nelson, A.~F., Quinn, T., \& Rice, W.~K.~M.\ 2007, Protostars and Planets V, 607 
\bibitem[Dutrey et al.(2003)]{dutrey03} Dutrey, A., Guilloteau, S., \& Simon, M.\ 2003, \aap, 402, 1003 
\bibitem[Dutrey et al.(2007)]{dutrey07} Dutrey, A., Guilloteau, S., \& Ho, P.\ 2007, Protostars and Planets V, 495
%\bibitem[Eisner \& Carpenter(2003)]{eisner03} Eisner, J.~A., \& Carpenter, J.~M.\ 2003, \apj, 598, 1341
\bibitem[Evans et al.(2003)]{evans03} Evans, N.~J., II, et al.\ 2003, \pasp, 115, 965
\bibitem[Fedele et al.(2009)]{fedele09} Fedele, D., van den Ancker, M.~E., Henning, T., Jayawardhana, R., \& Oliveira, J.~M.\ 2009, arXiv:0911.3320 
%\bibitem[Goldsmith \& Langer(1999)]{goldsmith99} Goldsmith, P.~F., \& Langer, W.~D.\ 1999, \apj, 517, 209 
\bibitem[Gorti \& Hollenbach(2009)]{gorti09} Gorti, U., \& Hollenbach, D.\ 2009, \apj, 690, 1539
\bibitem[Guilloteau \& Dutrey(1994)]{guilloteau94} Guilloteau, S., \& Dutrey, A.\ 1994, \aap, 291, L23
\bibitem[Hartmann et al.(2005)]{hartmann05} Hartmann, L., et al.\ 2005, \apjl, 628, L147 
%\bibitem[Herbig(1988)]{herbig88} Herbig, G.~H.\ 1988, \apj, 331, 999 
\bibitem[Hughes et al.(2010)]{hughes10} Hughes, A.~M., et al.\ 2010, arXiv:1007.3267 
%\bibitem[Holman \& Wiegert(1999)]{holman99} Holman, M.~J., \& Wiegert, P.~A.\ 1999, \aj, 117, 621
%\bibitem[Ida \& Lin(2008)]{ida08} Ida, S., \& Lin, D.~N.~C.\ 2008, \apj, 673, 487
\bibitem[Ingleby et al.(2009)]{ingleby09} Ingleby, L., et al.\ 2009, \apjl, 703, L137
\bibitem[Jayawardhana et al.(1999)]{jayawardhana99} Jayawardhana, R., Hartmann, L., Fazio, G., Fisher, R.~S., Telesco, C.~M., \& Pi{\~n}a, R.~K.\ 1999, \apjl, 520, L41 
\bibitem[Jensen(1996)]{jensen96} Jensen, E.~L.~N.\ 1996, Ph.D.~Thesis,
%\bibitem[Kastner et al.(1997)]{kastner97} Kastner, J.~H., Zuckerman, B., Weintraub, D.~A., \& Forveille, T.\ 1997, Science, 277, 67
\bibitem[Kenyon \& Hartmann(1995)]{kenyon95} Kenyon, S.~J., \& Hartmann, L.\ 1995, \apjs, 101, 117
\bibitem[Koerner \& Sargent(1995)]{koerner95} Koerner, D.~W., \& Sargent, A.~I.\ 1995, \aj, 109, 2138 
\bibitem[Kuznetsov et al.(1982)]{kuznetsov82} Kuznetsov, I.~V., Fedoseev, L.~I., \& Shvetsov, A.~A.\ 1982, Radiofizika, 25, 247
%\bibitem[Lin et al.(2006)]{lin06} Lin, S.-Y., Ohashi, N., Lim, J., Ho, P.~T.~P., Fukagawa, M., \& Tamura, M.\ 2006, \apj, 645, 1297
\bibitem[Mangum(1993)]{mangum93} Mangum, J.~G.\ 1993, \pasp, 105, 117
\bibitem[Mart{\'{\i}}n \& Magazz{\`u}(1999)]{martin99} Mart{\'{\i}}n, E.~L., \& Magazz{\`u}, A.\ 1999, \aap, 342, 173
\bibitem[Meyer et al.(2006)]{meyer06} Meyer, M.~R., et al.\ 2006, \pasp, 118, 1690 
\bibitem[Meyer et al.(2007)]{meyer07} Meyer, M.~R., Backman, D.~E., Weinberger, A.~J., \& Wyatt, M.~C.\ 2007, Protostars and Planets V, 573 
%\bibitem[Meyer et al.(2008)]{meyer08} Meyer, M.~R., et al.\ 2008, \apjl, 673, L181 
\bibitem[Meyer(2009)]{meyer09} Meyer, M.~R.\ 2009, IAU Symposium, 258, 111
\bibitem[Monin et al.(2007)]{monin07} Monin, J.-L., Clarke, C.~J., Prato, L., \& McCabe, C.\ 2007, Protostars and Planets V, 395
\bibitem[Najita et al.(2007)]{najita07} Najita, J.~R., Carr, J.~S., Glassgold, A.~E., \& Valenti, J.~A.\ 2007, Protostars and Planets V, 507 
\bibitem[Padgett et al.(2006)]{padgett06} Padgett, D.~L., et al.\ 2006, \apj, 645, 1283 
\bibitem[Pascucci et al.(2006)]{pascucci06} Pascucci, I., et al.\ 2006, \apj, 651, 1177 
\bibitem[Qi et al.(2004)]{qi04} Qi, C., et al.\ 2004, \apjl, 616, L11
\bibitem[Roberge et al.(2005)]{roberge05} Roberge, A., Weinberger, A.~J., Redfield, S., \& Feldman, P.~D.\ 2005, \apjl, 626, L105
\bibitem[Schlingman et al. (in prep.)]{schlingman} Schlingman et al., in prep.
\bibitem[Scoville et al.(1986)]{scoville86} Scoville, N.~Z., Sargent, A.~I., Sanders, D.~B., Claussen, M.~J., Masson, C.~R., Lo, K.~Y., \& Phillips, T.~G.\ 1986, \apj, 303, 416 
\bibitem[Silverstone et al.(2006)]{silverstone06} Silverstone, M.~D., et al.\ 2006, \apj, 639, 1138 
\bibitem[Skrutskie et al.(1993)]{strutskie93} Skrutskie, M.~F., et 
al.\ 1993, \apj, 409, 422 
\bibitem[Th{\'e}bault et al.(2004)]{thebault04} Th{\'e}bault, P., Marzari, F., Scholl, H., Turrini, D., \& Barbieri, M.\ 2004, \aap, 427, 1097
\bibitem[Thi et al.(2001)]{thi01} Thi, W.~F., et al.\ 2001, \apj, 561, 1074
%\bibitem[Ulich \& Haas(1976)]{ulich76} Ulich, B.~L., \& Haas, R.~W.\ 1976, \apjs, 30, 247 
\bibitem[van Kempen et al.(2007)]{vankempen07} van Kempen, T.~A., van Dishoeck, E.~F., Brinch, C., \& Hogerheijde, M.~R.\ 2007, \aap, 461, 983
\bibitem[Webb et al.(1999)]{webb99} Webb, R.~A., Zuckerman, B., Platais, I., Patience, J., White, R.~J., Schwartz, M.~J., \& McCarthy, C.\ 1999, \apjl, 512, L63 
\bibitem[White \& Hillenbrand(2005)]{white05} White, R.~J., \& Hillenbrand, L.~A.\ 2005, \apjl, 621, L65 
\bibitem[Wichmann et al.(2000)]{wichmann00} Wichmann, R., et al.\ 2000, \aap, 359, 181
\end{thebibliography}