Low-mass star formation in R Coronae Australis: Observations of organic molecules with the APEX telescope
F. L. Sch\"oier\inst{1}
J.~K. J{\o}rgensen\inst{2}
K.~M. Pontoppidan\inst{3}
A.~A. Lundgren\inst{4}
This paper presents new APEX submillimetre molecular line observations of three low-mass protostars, IRS7A, IRS7B, and IRAS32, in the R~Coronae~Australis molecular cloud complex. The molecular excitation analysis is performed using a statistical equilibrium radiative transfer code. The derived beam averaged fractional abundances vary by less than a factor of two among the three sources, except those of H$_2$CO and CH$_3$OH, which show differences of about an order of magnitude. The molecular abundances are similar to those typically found in other star-forming regions in the Galaxy, such as the $\rho$~Oph and Perseus molecular clouds. There is a marked difference in the kinetic temperatures derived for the protobinary source IRS7 from H$_2$CO ($\approx$\,40--60\,K) and CH$_3$OH ($\approx$\,20\,K), possibly indicating a difference in origin of the emission from these two molecules. %The derived gas-phase CO abundances for IRS7A and IRS7B are consistent with observed CO ice column densities if CO is completely frozen out at temperatures below 20\,K in the envelopes. % The similarity of RCrA--IRS7 to the well studied protobinary IRAS\,16293--2422 makes it an interesting candidate to harbour a 'hot core`. \keywords{Stars: formation -- ISM:molecules -- ISM:abundances -- astrochemistry}
schoeier.tex
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\begin{document}
%
   \title{Low-mass star formation in R Coronae Australis: Observations of organic molecules with the APEX telescope}
 
   \titlerunning{Low-mass starformation in Coronae Australis}

   \author{F. L. Sch\"oier\inst{1} \and
             J.~K. J{\o}rgensen\inst{2} \and
	   K.~M. Pontoppidan\inst{3} \and
	   A.~A. Lundgren\inst{4}}

   \offprints{F. L. Sch\"oier \\ \email{fredrik@astro.su.se}}

   \institute{Stockholm Observatory, AlbaNova University Center, SE-106 91 Stockholm, Sweden
     \and
	     Harvard-Smithsonian Center for Astrophysics, 60 Garden Street MS42, Cambridge, MA 02138, USA
             \and
	     Division of Geological and Planetary Sciences, California Institute of Technology, MS
                     150-21, Pasadena, CA 91125, USA
                     \and
                      European Southern Observatory, Casilla 19001, Santiago 19, Chile}    
             
   \date{Received; accepted}

   \abstract{This paper presents new APEX submillimetre molecular line observations of three low-mass protostars, IRS7A, IRS7B, and IRAS32, in the R~Coronae~Australis molecular cloud complex. The molecular excitation analysis is performed using a statistical equilibrium radiative transfer code. The derived beam averaged fractional abundances vary by less than a factor of two among the three sources, except those of H$_2$CO and CH$_3$OH, which show differences of about an order of magnitude.  The molecular abundances  are similar to those typically found in other star-forming regions in the Galaxy, such as the $\rho$~Oph and Perseus molecular clouds. There is a marked difference in the kinetic temperatures derived for the protobinary source IRS7 from H$_2$CO ($\approx$\,40--60\,K) and CH$_3$OH ($\approx$\,20\,K), possibly indicating a difference in origin of the emission from these two molecules.
   %The derived gas-phase CO abundances for IRS7A and IRS7B are consistent with observed CO ice column densities if CO is completely frozen out at temperatures below 20\,K in the envelopes. 
%   The similarity of RCrA--IRS7 to the well studied protobinary IRAS\,16293--2422 makes it an interesting candidate to harbour a 'hot core`.
   
   \keywords{Stars: formation -- ISM:molecules -- ISM:abundances -- astrochemistry}
   }
   \maketitle
%
%________________________________________________________________

%
   \begin{figure*}
   \centering{   
   \includegraphics[width=11.7cm]{fig1.eps}
   \caption{New observations of molecular line emission  using the APEX telescope. The velocity resolution is 0.2\,km\,s$^{-1}$ in all the spectra.}
   \label{obs_apex}}
   \end{figure*}
%

\section{Introduction}
%
In the earliest, deeply-embedded stage, low-mass protostars are
 surrounded by a collapsing envelope and a circumstellar disk.
%through which material is accreted onto the growing star. 
In this phase, 
%the final
%mass of the star, the final mass of the disk, and 
the initial chemical composition of the disk is determined.  The chemistry of molecules in the circumstellar envelope 
is likely to be directly reflected in the molecular composition in the disk, which forms the building blocks from which planets may form.
An interesting aspect is the search for  low-mass
'hot cores`, i.e., regions where grain ice mantles are liberated in
the envelope close to the central protostar because of its heating
\citep[e.g.,][]{Ceccarelli00b, Schoeier02a, Cazaux03}. The search for hot cores around low-mass stars is not yet
conclusive, mainly due to limitations in the available observational
data. Also, outflows are likely to play an important role by affecting many of the same species thought to be sign posts of hot cores
\citep[e.g.,][]{Joergensen05b}. Moreover, the chemistry of traditional hot-core tracers such as
H$_2$CO may be affected by the freeze-out of CO \citep[e.g.,][]{Schoeier04a}, which is why combined studies of chemically related molecular species are important.

The R Coronae Australis dark cloud is a nearby \citep[170\,pc,][]{Knude98} region of active star formation \citep[e.g.,][]{Brown87, Chini03, Nutter05}.
Molecular line observations of protostellar objects in R Coronae Australis  have attained little attention, mainly due to its location on the southern hemisphere. 
Previous studies using single-dish telescopes with limited spatial resolution, which are mainly sensitive
to the chemistry of the outer cold regions of the envelopes, have been biased towards lower excitation lines. A wealth of molecular rotational transitions lying in the 279\,$-$\,381\,GHz window are accessible with APEX\footnote{This publication is based on data acquired with the Atacama Pathfinder Experiment (APEX). APEX is a collaboration between the Max-Planck-Institut f\"ur Radioastronomie, the European Southern Observatory, and the Onsala Space Observatory}, and can be used to constrain the 
physics, and in particular, the chemistry in the dense
($>$\,10$^{6}$\,cm$^3$) and warm ($>$\,30\,K)
material in the envelopes.
Observations of protostellar objects in Corona Australis are also interesting for testing the importance of the
surroundings on the formation of protostars, and will complement
 previous extensive studies in Perseus and $\rho$ Oph \citep{Joergensen02, Joergensen04c, Joergensen05b}. It is not clear how the
properties of the parental cloud affect 
%{\bf (are reflected in)} 
the properties of
the emerging protostar \citep[see][]{Jayawardhana01}.
% for example, suggest
%that Class 0 objects simply are Class I objects formed in denser
%environments instead of, as has been traditionally assumed, being objects in different evolutionary stages \citep{Andre00}. 
Likewise, the properties of the parental cloud may
reflect the chemical evolution of the protostar, for example, in the
degree of freeze-out and the 'initial abundances`.

The first step in a project aimed at obtaining an extensive data base of molecular line emission towards a large sample of protostars in the southern hemisphere, with a focus on the Corona Australis and Chameleon molecular clouds, is presented in this {\em Letter}. 
It is crucial to utilize APEX to build comparable samples of well-studied low-mass protostars for future ALMA observations

\section{Observations}
\label{sect_obs}
%
Observations of sub-millimetre molecular line emission 
%from H$_2$CO and CH$_3$OH 
towards the R Coronae Australis molecular cloud complex were performed September 4--6, 2005, using the 12\,m APEX telescope equipped with an SIS receiver (APEX--2a). Typical system temperatures were in the range 140\,$-$\,380\,K (DSB). The 1\,GHz FFTS with 8192 spectral channels was used as a backend, providing a spectral resolution of 122\,kHz (0.1 km\,s$^{-1}$). The observations were carried out using a position-switching mode, with the reference position located at +2$\arcmin$ in azimuth. 
The data were reduced in a standard way by removing, in most cases, a low order ($\leq$\,3) polynomial baseline and the raw spectra, stored in the $T_{\mathrm A}^{\ast}$  scale, converted to main-beam brightness temperature using $T_{\mathrm{mb}}$\,=\,$T_{\mathrm
A}^{\ast}/\eta_{\mathrm{mb}}$. $T_{\mathrm A}^{\ast}$ is the
antenna temperature corrected for atmospheric attenuation using the
chopper-wheel method, and $\eta_{\mathrm{mb}}$ is the main-beam
efficiency. A value of $\eta_{\mathrm{mb}}$\,$=$\,0.7 was assumed for all spectral settings.
Regular pointing checks were made on strong CO sources and typically found to be consistent within $\approx$\,3$\arcsec$. The FWHM of the main beam is 18\arcsec at 345\,GHz.
The uncertainty (including errors in pointing, calibration, and $\eta_{\mathrm{mb}}$) in the absolute intensity scale (main-beam brightness) is estimated to be about $\pm 15$\%.

The observed spectra for the protostellar objects IRS7A ($\alpha_{2000}$\,=\,19:01:55.3, $\delta_{2000}$\,=\,--36:57:21), IRS7B ($\alpha_{2000}$\,=\,19:01:56.4, $\delta_{2000}$\,=\,--36:57:27), and IRAS32 ($\alpha_{2000}$\,=\,19:02:58.7, $\delta_{2000}$\,=\,--37:07:34) are presented in Figs.~\ref{obs_apex} and \ref{obs_ch3oh} (the positions are obtained from IRAC4 8\,$\mu$m pipeline data in the Spitzer Space Telescope archive: AOR ID 3650816). IRS7B is identified as a class 0 protostar, whereas  IRS7A and IRAS32 are of class I \citep[see][ and references therein]{Nutter05}.

Results of Gaussian fits to the spectra are reported in Table~\ref{intensities}. The systemic (LSR) velocities obtained from the Gaussian fits are 5.3, 5.6, and 5.9\,km\,s$^{-1}$ (LSR) for IRS7A, IRS7B, and IRAS32, respectively. The spectra are generally well represented by a single Gaussian, with the major exception being the observed HCN $J$\,=\,4\,$\rightarrow$\,3 line emission towards IRS7A and IRS7B. In these two cases, not only is the HCN line significantly broader, compared to other molecular lines, but also a central absorption dip is present, at least in the IRS7A spectrum. The HCN spectra presented here 
are very similar to those of HCO$^{+}$ $J$\,=\,4\,$\rightarrow$\,3 obtained by \citet{Anderson97b}. They interpret the emission as arising in a large rotating (circumbinary) disk, and the self-absorption feature is caused by absorption in a dense and quiescent foreground of the cloud. In addition, for the two IRS7 sources,  our observed C$^{17}$O and, in particular, C$^{18}$O,  $J$\,=\,3\,$\rightarrow$\,2 spectra have weak wings, possibly related to the known molecular outflow in this region, and are somewhat better reproduced by a sum of two Gaussians.

%
   \begin{figure}
   \centering{   
   \includegraphics[width=6.3cm]{fig2.eps}
   \caption{New observations of CH$_3$OH line emission using the APEX telescope. The velocity resolution is 0.5\,km\,s$^{-1}$ in both spectra. Note that the $7_0$\,$\rightarrow$\,$6_{0}$ E$^{+}$ line is blended with the CH$_3$OH $4_0$\,$\rightarrow$\,$3_{1}$ E$^{-}$ line from the upper sideband.}
   \label{obs_ch3oh}}
   \end{figure}
%

%
\begin{table}
\caption{Observational results$^{\mathrm{a}}$.}
\label{intensities}
$
\begin{array}{p{0.2\linewidth}ccccccccccc}
\hline
\noalign{\smallskip}
&
\multicolumn{3}{c}{\mathrm{IRS7A}} &&
\multicolumn{3}{c}{\mathrm{IRS7B}} &&
\multicolumn{3}{c}{\mathrm{IRAS\ 32}} 
\\

\cline{2-4}
\cline{6-8}
\cline{10-12}
\multicolumn{1}{c}{\mathrm{Transition}} &
\multicolumn{1}{c}{I} &
\multicolumn{1}{c}{T_{\mathrm{mb}}} &
\multicolumn{1}{c}{\Delta v} &&
\multicolumn{1}{c}{I} &
\multicolumn{1}{c}{T_{\mathrm{mb}}} &
\multicolumn{1}{c}{\Delta v} &&
\multicolumn{1}{c}{I} &
\multicolumn{1}{c}{T_{\mathrm{mb}}} &
\multicolumn{1}{c}{\Delta v}
\\
\noalign{\smallskip}
\hline
\noalign{\smallskip}
C$^{18}$O 3--2                            & 24.9 & 7.2$:$   & 3.1$:$  && 26.9 & 10.2$:$ & 2.3$:$ && 6.5 & 5.2 & 1.0 \\
C$^{17}$O 3--2                            &   7.5 &  2.5$:$   & 2.5$:$  &&   8.8 &    3.7$:$ & 1.9$:$ && 2.6 & 1.9 & 1.0 \\
C$^{34}$S 7--6                            &   1.4 &  0.47 & 2.5 &&   2.0 &    0.61 & 2.1 && {<}0.2 & \cdots & \cdots \\
CS 7--6                                          &  17.6 &  6.2 & 2.7 &&  17.6 &  8.1 & 2.0 && 4.6 & 2.8$:$ & 1.3$:$ \\
HCN 4--3                                       &  26.0 &  6.1$:$ & 3.9$:$ &&  30.4 &  9.2$:$ & 3.0$:$ && 2.7 & 1.2$:$ & 1.8$:$ \\
HNC 4--3                                       &  8.1 &  3.1 & 2.4 &&  11.3 &  5.9 & 1.8 && 1.8 & 1.1$:$ & 1.0$:$ \\
C$_2$H:                                        &\\
  \ 4$_5$--3$_4$                            &   7.7 &  2.6 & 2.8 &&   10.0 &  4.0 & 2.3 && 2.4 & 1.1$:$ & 2.1$:$ \\
  \ 4$_4$--3$_3$                           &   6.4 &  1.9 & 2.9 &&   8.4 &    2.2 & 2.4 && 2.1 & 0.7$:$ & 2.4$:$ \\
H$_2$CO:                                     &\\
  \ 5$_{23}$--4$_{22}$                  &   1.8 &  0.51 & 2.8 &&   1.7 &    0.81 & 2.0 && {<}0.2 & \cdots & \cdots \\
  \ 5$_{05}$--4$_{04}$                  &   6.2 &  2.4 & 2.5 &&   7.7 &    3.8 & 1.8 && 0.3$:$ & \cdots & \cdots \\
CH$_3$OH:                                 & \\
  \ 7$_0$--6$_0$ E$^+$$^{\mathrm b}$        &  2.3 & \cdots & \cdots &&  3.0 & \cdots & \cdots && {<}0.1 & \cdots & \cdots \\
  \ 7$_1$--6$_1$ E$^-$         & 1.5 & 0.50 & 2.6 &&  2.2 &  1.0  & 2.1 && {<}0.1 & \cdots & \cdots  \\
  \ 7$_1$--6$_1$ E$^+$        & 0.49 & \cdots & \cdots &&  0.37 & \cdots & \cdots &&  {<}0.1 & \cdots & \cdots \\
  \ 7$_2$--6$_2$ E$^{\pm}$ &  0.51 & 0.18$:$  & 3.0$:$ && 0.90 & 0.35$:$ & 2.6$:$  &&  {<}0.1 & \cdots & \cdots \\
 \ 7$_0$--6$_0$ A$^+$          & 1.8 & 0.64 & 2.4 && 3.3 & 1.2 & 2.3 &&  {<}0.1 & \cdots & \cdots \\
  \ 7$_3$--6$_3$ A$^{\pm}$  & 0.2$:$  & \cdots & \cdots &&  0.2$:$ & \cdots & \cdots && {<}0.1 & \cdots & \cdots  \\
\noalign{\smallskip}
\hline
\end{array}
$

$^{\mathrm a}$ $I$\,$=$\,$\int T_{\mathrm{mb}}\,dv$ is the integrated intensity in K\,km\,s$^{-1}$. $T_{\mathrm{mb}}$ is the peak intensity (in K) and $\Delta v$ the FWHM (in km\,s$^{-1}$) from Gaussian fits. A colon (:) marks an uncertain value due to low a S/N-ratio or a significant departure from a Gaussian line profile.

$^{\mathrm b}$ Blended with the CH$_3$OH 4$_0$\,$\rightarrow$\,3$_1$ E$^-$ line from the upper sideband. The integrated intensity refers to the sum of both lines. 
\end{table}

\section{Excitation analysis}
\label{sect_model}
%

Presently, there is a lack of observational data required to carry out a detailed modelling of the dust emission,  mainly due to limited spatial resolution \citep[see][]{Nutter05}. This is usually the preferred method to determine the physical structure, such as density and temperature, of material present around protostars that is required for a detailed analysis of its molecular line emission  \citep[e.g.,][]{Joergensen02}.

A first estimate of the excitation conditions can be obtained assuming the lines to be emanating from the same homogeneous region. We have adopted RADEX, a statistical equilibrium radiative transfer code that uses an escape probability formalism to calculate the excitation and emission from the protostellar cores in our sample. RADEX is
comparable to the LVG method and provides a useful tool for rapidly
analysing a large set of observational data providing constraints on
physical conditions, such as density and kinetic temperature (Jansen et al.\ 1994\nocite{Jansen94}; van der Tak et al., in prep.).
An online version of RADEX, together with relevant molecular data used in the excitation analysis and summarized in \citet{Schoeier05a},  is made publicly available through the {\em Leiden Atomic and Molecular Database} (LAMDA){\footnote{\tt http://www.strw.leidenuniv.nl/$\sim$moldata}}.

%{\bf For H$_2$CO and CH$_3$OH, where more than one transition was observed, a first estimate of the excitation conditions can be obtained assuming the lines to be emanating from the same, homogeneous, region where the level populations are described by one single temperature. In the analysis no corrections for beam-filling effects are made and all derived parameters are averaged over the APEX beam.

%In the simplest approach, rotation-diagram method \citep[e.g.,][]{Blake87}, the emission is assumed to be optically thin and the excitation is described by the rotation temperature which may be quite different from the true kinetic temperature of the gas. The rotation-diagram method gives temperatures of 44\,K (35\,K) and 18\,K (15\,K) for H$_2$CO and CH$_3$OH, respectively, in the envelope around R CrA IRS7A (IRS7B). This could possibly indicate different origins for these traditional hot-core tracers. The derived column densities are  2\,$\times$\,10$^{13}$\,cm$^{-2}$ (3\,$\times$\,10$^{13}$\,cm$^{-2}$) and 2\,$\times$\,10$^{14}$\,cm$^{-2}$ (3\,$\times$\,10$^{14}$\,cm$^{-2}$) for H$_2$CO and CH$_3$OH, respectively.}

%In order to determine the molecular excitation in the circumstellar envelopes  {\bf in more detail} a detailed non-LTE radiative transfer code, based on the Monte Carlo method, was used.  The code is described in detail in \citet{Schoeier02a} and has been
%benchmarked, to high accuracy, against a wide variety of molecular-line radiative
%transfer codes in \citet{Zadelhoff02}.

In the excitation analysis, the number density of H$_2$ molecules was fixed to 3\,$\times$\,10$^6$\,cm$^{-3}$, typical of the warmer, more dense regions of the condensed cores. From the two lines of H$_2$CO, the kinetic temperatures were estimated to be $\approx$\,60\,K and $\approx$\,40\,K for IRS7A and IRS7B, respectively. The corresponding column densities derived were 2\,$\times$\,10$^{13}$\,cm$^{-2}$ (IRS7A) and  3.5\,$\times$\,10$^{13}$\,cm$^{-2}$ (IRS7B). In the case of CH$_3$OH, a full analysis using a $\chi^2$-statistic was used, and the kinetic temperatures derived were 23$\pm$3\,K and  22$\pm$3\,K for IRS7A and IRS7B, respectively. The column densities of CH$_3$OH derived were 1.4$\pm$0.4\,$\times$\,10$^{14}$\,cm$^{-2}$ (IRS7A) and 1.8$\pm$0.4\,$\times$\,10$^{14}$\,cm$^{-2}$ (IRS7B). The substantial difference in the derived kinetic temperatures from H$_2$CO and CH$_3$OH could possibly indicate different origins for these traditional hot-core tracers.

To determine the column densities for the remaining molecular species, we adopted a kinetic temperature of 40\,K. The column densities obtained for C$^{34}$S in the case of IRS7A (1\,$\times$\,10$^{13}$\,cm$^{-2}$) and IRS7B (1.5\,$\times$\,10$^{13}$\,cm$^{-2}$) were used together with a 'standard` fractional abundance (5\,$\times$\,10$^{-11}$) of C$^{34}$S \citep[][ see Table~\ref{abundances}]{Joergensen04c} to estimate the total H$_2$ column densities. In the case of IRS7A, we obtain a H$_2$ column density of 2\,$\times$\,10$^{23}$\,cm$^{-2}$, and for IRS7B, a slightly higher value of 3\,$\times$\,10$^{23}$\,cm$^{-2}$.
 For IRAS32, no C$^{34}$S $J$\,=\,7\,$\rightarrow$\,6 line emission was detected, and instead the main isotope (with an adopted fractional abundance of 1\,$\times$\,10$^{-9}$)  was used to estimate the H$_2$ column density to 3\,$\times$\,10$^{22}$\,cm$^{-2}$. The derived fractional abundances for all molecular species observed are reported in Table~\ref{abundances}.

%The circumstellar envelopes  are assumed to be spherically symmetric and static with no global velocity field. The spectral lines are broadened only through the local line width as determined by the gas kinetic temperature and micro-turbulence. The micro-turbulence is assumed to have a width of  2.0, 1.7 and 0.9\,km\,s$^{-1}$ (FWHM) for  IRS7A, IRS7B, and IRAS32, respectively, based on model fits to observed lines. 
%As long as the lines are optically thin, which is the case for the majority of lines in this study, the adopted value of the micro-turbulence will not affect the derived abundances \citep{Joergensen02}.

%The density structure is assumed to follow a single power-law distribution, $n_{\mathrm{H_2}}$\,=\,$n_0(r/r_0)^{-1.7}$, where $n_0$ is the H$_2$ number density at $r_0$\,=\,1000\,AU. Such a density profile is found to reproduce observations well \citep{Joergensen02}. The adopted gas kinetic temperature structure is $T(r)$\,$=$\,$T_0(r/r_0)^{-0.4}$, which is relevant for the scales probed by the current observations. In all cases $T_0$\,=\,25\,K (at $r_0$\,=\,1000\,AU) was assumed. In the optically thin limit the adopted $T_0$ corresponds to a central luminosity of 7\,L$_\odot$.  The inner and outer radii were fixed to 2.0\,$\times$\,10$^{14}$\,cm (13\,AU) and 5.0\,$\times$\,10$^{16}$\,cm (3300\,AU), respectively. 

%To obtain the absolute density scale, $n_0$, average CS abundances from the analysis of  \citet{Joergensen04c} were used. For IRS7A and IRS7B a C$^{34}$S fractional abundance relative to H$_2$ of 2.0$\times$\,10$^{-10}$ and 6.8\,$\times$\,10$^{-11}$ were adopted, characteristic of a Class I and Class 0 protostar (based on the classification in Nutter et al.\ 2005\nocite{Nutter05}), respectively. Then $n_0$ was modified until a good fit to the observed  C$^{34}$S $J$\,=\,7\,$\rightarrow$\,6 line intensity was found.  Values for $n_0$ of 3\,$\times$\,10$^{6}$\,cm$^{-3}$ and 6\,$\times$\,10$^{6}$\,cm$^{-3}$ were derived for IRS7A and IRS7B, respectively. In the case of the Class I source \citep{Nutter05} IRAS32 no C$^{34}$S $J$\,=\,7\,$\rightarrow$\,6 line emission was detected and instead the main isotope (with an adopted fractional abundance of 4\,$\times$\,10$^{-9}$)  was used to fix $n_0$ at 1\,$\times$\,10$^{6}$\,cm$^{-3}$. With the density scale determined fractional abundances of the remaining molecules could be obtained and are reported in Table~\ref{abundances}. In all cases the fractional abundance was assumed to stay constant throughout the circumstellar envelope. No modelling was performed for the highly optically thick line of HCN.
 
%
\begin{table}
\caption{Model results.}
\label{abundances}
\resizebox{\hsize}{!}{
$ 
\begin{array}{p{0.16\linewidth}ccccccc}
\hline
\noalign{\smallskip}
&
\multicolumn{1}{c}{\mathrm{IRS7A}} &&
\multicolumn{1}{c}{\mathrm{IRS7B}} &&
\multicolumn{1}{c}{\mathrm{IRAS32}} &&
\multicolumn{1}{c}{\mathrm{standard}^{a}} 
\\
\cline{2-2}
\cline{4-4}
\cline{6-6}
\cline{8-8}
\multicolumn{1}{c}{\mathrm{Molecule}} &
\multicolumn{1}{c}{f(\mathrm{X})/f(\mathrm{H}_2)} &&
\multicolumn{1}{c}{f(\mathrm{X})/f(\mathrm{H}_2)} &&
\multicolumn{1}{c}{f(\mathrm{X})/f(\mathrm{H}_2)} &&
\multicolumn{1}{c}{f(\mathrm{X})/f(\mathrm{H}_2)}
\\
\noalign{\smallskip}
\hline
\noalign{\smallskip}
C$^{18}$O                             & 6.0\times10^{-8}\phantom{:^{1}}          && 5.0\times10^{-8}\phantom{:^{1}}           && \phantom{<}1.0\times10^{-7}\phantom{:^{1}}          && 4.1\times10^{-8}\phantom{^{1}}  \\
C$^{17}$O                             & 2.0\times10^{-8} \phantom{:^{1}}         &&  1.5\times10^{-8}\phantom{:^{1}}          && \phantom{<}2.5\times10^{-8}\phantom{:^{1}}          && 1.4\times10^{-8}\phantom{^{1}}  \\
C$^{34}$S                             &  5.0\times10^{-11}$:$                             && 5.0\times10^{-11}$:$                              && \phantom{<}\cdots\phantom{:^{1}}                            &&  5.1\times10^{-11} \\
CS                                           &  8.0\times10^{-10}\phantom{:}               && 6.0\times10^{-10}\phantom{:}           && \phantom{<}1.0\times10^{-9}$:$\phantom{^{1}}    && 2.5\times10^{-9} \phantom{^{1}} \\
C$_2$H                                  & 5.0\times10^{-10}\phantom{:}           &&  4.0\times10^{-10}\phantom{:}           &&  \phantom{<}9.0\times10^{-10}\phantom{:}        && \cdots \phantom{:^{1}} \\
HCN                                        &  \cdots\phantom{:^{1}}                             && \cdots\phantom{:^{1}}                               && \phantom{<}\cdots \phantom{:^{1}}                          && 2.2\times10^{-9} \phantom{^{1}}\\
HNC                                        &  1.0\times10^{-10}\phantom{:}                && 1.0\times10^{-10}\phantom{:}          && \phantom{<}1.2\times10^{-10}\phantom{:}       && 2.5\times10^{-10}\\
p-H$_2$CO                           &  2.0\times10^{-10}\phantom{:}                && 1.0\times10^{-10}\phantom{:}          && \phantom{<}4.0\times10^{-11}\phantom{:}       && 6.9\times10^{-10}\\
A-CH$_3$OH                        & 6.0\times10^{-10}\phantom{:}           && 7.0\times10^{-10}\phantom{:}            && {<}7.0\times10^{-11}\phantom{:}                           && 1.8\times10^{-9}\phantom{^{1}} \\
E-CH$_3$OH                        & 7.0\times10^{-10}\phantom{:}           && 6.0\times10^{-10}\phantom{:}           && {<}1.3\times10^{-10}\phantom{:}                            && 1.3\times10^{-9}\phantom{^{1}} \\
\noalign{\smallskip}
\hline
\end{array}
$}
$^a$ Median values from  \citet{Joergensen02}, \citet{Joergensen04c}, and \citet{Joergensen05b}. Note that CH$_3$OH was only detected in known outflow sources.\\
A colon (:) marks an adopted value used to set the density scale.
\end{table}

\section{Results and discussion}
%
The excitation analysis gives results (see Table~\ref{abundances}) for the abundances of common molecules such as CO, CS, and HNC, and for more complex species, such as H$_2$CO and CH$_3$OH, that are consistent with results from other star-forming regions, such as $\rho$~Oph and Perseus \citep{Joergensen02, Joergensen04c, Joergensen05b}. The derived fractional abundances vary by less than a factor of two among the three sources for the majority of molecular species. In the present analysis, it is not possible to correlate the abundances with the identified protostellar class.
It is interesting to note that the abundances of both H$_2$CO and CH$_3$OH vary by about an order of magnitude among the three sample sources, possibly indicating a difference in central heating source or impact of outflows. We find the abundances of these molecules to be particularly low in IRAS32.

IRS7A and IRS7B form a protobinary system with a separation of about 18$\arcsec$ (3000\,AU). As such,  RCrA--IRS7 is also a cousin of the well-studied IRAS 16293--2422 and NGC--IRAS4 binaries in terms of the total luminosity. The fact that \citet{Joergensen05b} only detected CH$_3$OH emission from known outflow sources in their large survey of protostellar cores, together with the low kinetic temperature ($\approx$\,20\,K) we derive from our observed CH$_3$OH emission in contrast to what we get for H$_2$CO ($\approx$\,40--60\,K), suggests the impact of outflows in regulating the chemistry around IRS7. These temperatures are similar to those obtained for NGC--IRAS4 \citep{Maret04,Maret05}, but lower than those for IRAS 16293--2422 \citep{Dishoeck95}. However, in contrast to the NGC--IRAS4A binary in particular, IRS7 does not show CH$_3$OH lines that are broader than other observed molecular lines. Nonetheless, a large scale outflow thought to be associated with IRS7 has been observed in several other molecular species \citep{Levreault88, Anderson97a}.

%This could possibly indicate that shocks are less important in regulating the chemistry in RCrA--IRS7.


%The abundances of these molecules are particularly low in IRAS32.
%The various lines of H$_2$CO and CH$_3$OH detected towards the two IRS7 sources are reasonably well modelled using a constant abundance. 
%However, there is an apparent difference in the observed line ratios of various CH$_3$OH lines between the two IRS7 sources.

%IRS7A and IRS7B form a protobinary system with a separation of about 18$\arcsec$ (3000\,AU). As such RCrA--IRS7 is a cousin of the well studied IRAS 16293--2422 and NGC--IRAS4A binaries also in terms of the total luminosity. However, in contrast to the NGC--IRAS4A binary in particular, RCrA--IRS7 does not show CH$_3$OH lines that are broader than other observed molecular lines. This could possibly indicate that shocks are less important in regulating the chemistry in RCrA--IRS7.

%The derived CO abundances for IRS7A and IRS7B are consistent with the CO ice column densities derived by \citet{Pontoppidan03} if CO is effectively frozen out onto dust grains at radii where $T$\,$<$\,20\,K. The CO ice abundance in such a scenario is about 8\,$\times$\,10$^{-5}$ and 6\,$\times$\,10$^{-5}$ for IRS7A and IRS7B, respectively.

To make further progress, in particular to solidify the claims regarding the chemistry in the direct neighbourhood of the observed protostars, more detailed knowledge of the underlying physical structure is required. Such work is in progress and involves mid-IR observations from large ground-based telescopes and the Spitzer space telescope, and interferometric mm/sub-mm observations from ATCA and the SMA. 

%These observations will provide a spatial resolution that can match the scales of the warm and dense regions probed by the molecular line emission presented here. Of interest is also to study the content and distribution of ices around these protostars.

\begin{acknowledgements}
%
We would like to thank the anonymous referee who helped us improve the presentation of this work.
The authors are grateful to the staff at the APEX telescope. 
FLS acknowledges financial support from the Swedish Research Council.
The research of JKJ was supported by the NASA Origins Grant NAG5--13050. 
%
\end{acknowledgements}

\bibliographystyle{aa}

%\bibliography{Defs,YSO,Radtransf,Moldata}

\begin{thebibliography}{20}
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