A&A 401, 613-624 (2003)
R. D. Alexander1, - M. M. Casali1 - P. André2 - P. Persi3 - C. Eiroa4
1 - UK Astronomy Technology Centre, Royal Observatory, Blackford Hill, Edinburgh, EH9 3HJ, UK
2 - Service d'Astrophysique, CEA/DSM/DAPNIA, C.E. Saclay, 91191 Gif-sur-Yvette Cedex, France
3 - IASF, CNR Roma, Via Fosso del Cavaliere, 00133 Roma, Italy
4 - Dpto. Física Teórica, C-XI, Facultad de Ciencias, UAM, Cantoblanco, 28049 Madrid, Spain
Received 28 August 2002 / Accepted 4 February 2003
We present the results of a mid-infrared (5-16.5 m) imaging spectroscopy survey of Young Stellar Objects (YSOs) and their surrounding environment in four low-mass star formation regions: RCrA, Ophiuchi, Serpens and Chamaeleon I. This survey was performed using ISOCAM and its Circular Variable Filters (CVF) and observed 42 YSO candidates: we were able to obtain complete 5-16.5 m spectra for 40 of these with a spectral resolving power of . A number of spectral features were measured, most notably the 9.7 m silicate feature, the bending modes of both water and CO2 ices at 6.0 and 15.2 m respectively and the well-known unidentified feature at 6.8 m. The strength of the unidentified feature was observed to correlate very well with that of the water ice bending mode and far less strongly with the CO2 ice bending mode. This suggests, in a manner consistent with previous observations, that the carrier of the unidentified feature is a strongly polar ice. Absorption profiles of the bending mode of CO2 ice are observed to show a significant long wavelength wing, which suggests that a significant fraction of the CO2 ice observed exists in a polar (H2O-rich) phase. The sources observed in RCrA, Oph and Serpens show similar spectral characteristics, whilst the sources observed in Cha I are somewhat anomalous, predominantly showing silicate emission and little or no absorption due to volatile ices. However this is consistent with previous studies of this region of the Cha I cloud, which contains an unusual cluster of YSOs. From comparisons of the strengths of the water ice and silicate bands we detect an apparent under-abundance of water ice towards the sources in Oph, relative to both RCrA and Serpens. This may be indicative of differences in chemical composition between the different clouds, or may be due to evaporation. Finally the CO2:H2O ice ratios observed towards the sources in Oph show significantly greater scatter than in the other regions, possibly due to varying local conditions around the YSOs in Oph.
Key words: ISM: dust, extinction - stars: circumstellar matter - stars: pre-main-sequence - ISM: molecules - ISM: lines and bands - infrared: ISM
Unfortunately, studies from the ground are constrained to work within atmospheric windows. In some cases, such as CO2, telluric absorption by the molecule at 15.2 m makes its study in astronomical objects impossible. Furthermore, the sensitivity of observations in the mid-IR is greatly reduced because telescope and atmospheric thermal emission greatly increase the photon shot noise. Consequently, ground-based studies longward of K band are generally restricted to observing relatively bright objects. The only solution to these problems is to move to an airborne or space platform, where even a modest sized telescope can outperform ground-based systems.
Mid-IR spectra of young stars at high resolution were successfully taken with the Infrared Space Observatory (ISO, see Kessler et al. 1996 for an overview of the ISO mission) by various investigators (e.g. W33A by Gibb et al. 2000; Elias 29 by Boogert et al. 2000). In an attempt to push mid-IR observations to even higher levels of sensitivity we made use of ISOCAM (Cesarsky et al. 1996) and the Circular Variable Filters (CVF) to image 42 YSO candidates in four well-known low-mass star formation regions: RCrA, Ophiuchi, Serpens and Chamaeleon I. The high sensitivity and good sample size was achieved by using the low CVF resolution and by observing dense clusters of YSOs, so that ISOCAM's imaging capability allowed simultaneous spectra to be obtained. It should be noted that the low spectral resolving power of the CVF (/ ) essentially prohibits the detection of gas-phase spectral features; this is a study of the molecular ices and dust that lie around these young stars. However it is possible to draw a number of conclusions about the physics and chemistry around these YSOs, with information about their composition, thermal behaviour and evolution readily obtainable from such spectroscopic observations.
5-16.5 m imaging spectroscopy was performed on six target fields in four low-mass star formation regions using the ISOCAM-CVF, on various dates between April 1996 and February 1997. This instrument made use of 2 variable width filters to take images at a number of different wavelengths, using a pixel gallium-doped silicon array. Exposure times of 2.1 s per CVF step were used. Each was exposed repeatedly and the images were sequenced (in time) in decreasing order of wavelength, resulting in a total exposure time of 2500 s per target region for each CVF set. This procedure produced 2 datacubes for each target region. The first covered the wavelength range 5.0-9.4 m (CVF1) and the other covered 9.3-16.3 m (CVF2), with a mean spectral resolving power of . A brief summary of the characteristics of the observed regions is given below.
The RCrA cloud is the name given to a molecular cloud of mass 120 (Harju et al. 1993) around the variable star R Coronae Australis, at a distance of 130 pc (Marraco & Rydgren 1981). It is a region of low-mass star formation located approximately 18 below the galactic plane (Olofsson et al. 1999), where it suffers little foreground obscuration. The previous K-band studies of Taylor & Storey (1984) and Wilking et al. (1997) and the sub-millimetre survey of Harju et al. (1993) revealed a number of embedded objects and other objects classified as YSOs. One field of view was observed in the RCrA cloud, using the 6 pixel field of view (pfov). It was centred on 190145 1, -36 58 34 , immediately to the south-west of RCrA itself.
The Ophiuchi dark cloud complex is a complicated structure of several large molecular clouds near the bright star Ophiuchi. Its relatively high galactic latitude ( ) combined with its relative proximity (d=160 pc, Whittet 1974) result in comparatively little foreground obscuration, making it an ideal region in which to study star formation. Initially identified by Grasdalen et al. (1973), who identified 41 sources in a K-band map, it has since been studied extensively both in the infrared (Elias 1978; Greene & Young 1992; Strom et al. 1995; Abergel et al. 1996; Bontemps et al. 2001) and at millimetre wavelengths (Loren et al. 1990; André & Montmerle 1994). In particular, Loren et al. (1990) showed that the cloud consists of a number of smaller sub-structures. They identified 6 dense cores (labelled A-F) and observed significantly higher clustering of YSOs in three of these cores (A, B, and E/F). Two fields were observed in Oph, both with the 6 pfov: one centred on 162621 4, -24 23 59 (field Oph A) and one at 162723 7, -24 40 39 (field Oph E). The positions correspond (approximately) to cores A and E/F above.
The Serpens molecular cloud was initially identified by Strom et al. (1976). It has since been very well surveyed, both in the infrared (Churchwell & Koornneef 1986; Gomez de Castro et al. 1988; Eiroa & Casali 1992; Giovannetti et al. 1998; Kaas 1999) and the sub-millimetre (Casali et al. 1993), with as many as 163 sources identified in the K-band map of Eiroa & Casali (1992). The cloud itself extends over a region of in CO line emission (Loren et al. 1979) and is at a distance of 260 pc (Straizys et al. 1996). In addition to survey imaging, many individual objects in the cloud have been studied in detail: Eiroa & Casali (1989) studied the multiple outflow source SVS4; SVS2 has been shown to illuminate the Serpens Reflection Nebula (Gomez de Castro et al. 1988; Huard et al. 1997); SVS20 has been shown to be a close binary system (Huard et al. 1997). Two fields were observed in Serpens. Field Ser B, taken with the 6 pfov, was of the region around SVS2 and SVS20 and was centred on 182957 6, 1 12 42 . Field Ser A was a higher resolution field (3 pfov) of the multiple outflow source SVS4, centred on 182957 2, 1 14 25 .
The Chamaeleon dark cloud complex is a complicated structure consisting of 3 large molecular clouds (designated Cha I, II, III by Hoffmeister 1963) and a number of smaller clumps and globules. Again, the relative proximity (d=160pc, Whittet et al. 1997) and high galactic latitude (16 ) of the clouds result in little foreground obscuration and make them ideal for star formation study. The region has been well-surveyed at millimetre wavelengths (Mattila et al. 1989; Henning et al. 1993) and more recently in the near- (Kenyon & Gómez 2001) and mid-infrared (Persi et al. 2000). A single field was observed in the Cha I cloud, with the 6 pfov, centred on 110939 4, -76 35 2 . The corresponds (approximately) to a previously observed dense core lying to the north of HD97300 (Jones et al. 1985; Persi et al. 1999).
Other spectral features were observed in individual sources, such as the apparent 10.5 m emission feature in RCrA IRS5 (see Fig. 2) or the apparent 14 m absorption feature in HH100-IR (see Fig. 4). These rarer features have been noted in the results tables where they occur, but the goal of this study was to look for general trends rather than individual curiosities. With this in mind, and also considering the large number of spectra obtained, it was decided to fit only the strongest and most commonly occuring features highlighted above.
As many of these spectral features overlap, independent determinations of their strengths were not possible, so all of the features in the spectra were fitted simultaneously. This was achieved by fitting a spectrum of the form:
The fitting procedure was iterative, with repeated fitting converging to a final profile. The fitting algorithm took the form:
It should be noted that this allowed the continuum to drift away from the observed spectrum in some cases, due to absorption across the entire observed wavelength range. This is because the normalised silicate opacity over the range 5-7.5 m is approximately constant at 0.15, and towards the long wavelength end of the spectra the 18 m bending feature begins to "cut in''. (The normalised opacity at 16 m is approximately 0.3.) The effect of this can be seen in Fig. 2: although the continua are seen to deviate significantly from the observed spectra, appearing to diverge from the observed spectra in some cases, this is primarily due to silicate absorption across the entire wavelength range. With only the silicate features included, these "continuum plus silicate'' fits are seen to follow the observed spectra well. As a result, any possible systematic errors arising from the manner of the continuum fit will only affect the evaluation of the silicate optical depth; the measured strengths of the narrower features are not affected by the choice of continuum.
|Figure 2: Examples of successful fits to GY265 in Oph (upper) and RCrA IRS5 (lower): fitted spectra are shown as solid lines and continua as dashed lines. The dotted lines are the "continuum plus silicate'' fits given by .|
On inspection of the residuals from this fitting procedure a further feature was observed: a broad feature centred on 11.2 m, seen in both emission and absorption. This has a number of possible identifications, such as crystalline silicates (Bregman et al. 1987; Campins & Ryan 1989), the Unidentified (PAH) Band at 11.3 m (Boulanger et al. 1996), or a shoulder on the silicate profile. Given the variation in possible identifications, it was decided to fit this feature independently using an asymmetric Gaussian profile, once again taking the widths to be the mean of those observed ( nm). This feature was then included in the algorithm and the fitting procedure repeated: examples of successful fits are shown in Fig. 2. Another similar residual around 8.5 m was also observed, but the properties and strength of this were greatly affected by the choice of continuum point around 8 m. Whereas the 11 m feature was robust against small changes in the continuum this apparent 8.5 m feature was not, so it was not included in the fitting procedure.
The 15.2 m CO2 ice feature is an interesting one and has been used in the past as a diagnostic of the ice environment. Gerakines et al. (1999) found that two distinct phases of CO2 ice (polar and non-polar) show significantly different absorption profiles. The non-polar phase is characterised by a double-peaked profile, arising from the two-fold degeneracy of the bending mode, whereas the polar (H2O-rich) phase shows a significant long wavelength wing. The spectral resolution of the CVF is too low to resolve the double-peak structure, but the CO2 profiles were observed to vary noticeably from source to source. As a result it was decided to let the CO2 profile vary, so it was fitted using an asymmetric Gaussian profile with the widths as variable parameters.
There are two forms of error associated with this fitting procedure. Firstly there are the ordinary statistical errors associated with the fits, evaluated by combining the statistical errors on the data points used to fit the features and the intrinsic uncertainties in the individual feature fits. These are typically around 5-10% for the stronger features (silicate, H2O, CO2) and around 15-20% for the less well constrained features, although in the case of some poor S/N spectra they are much larger. In addition to this, there are systematic errors associated with poor fits. These are usually due to the presence of "extra'' features in the spectra, or because of poorly fitting feature profiles, and are less easy to quantify. Consequently, the errors reported in the results tables are the statistical errors only; where fits are flagged as poor additional systematic errors also apply.
In such cases, the narrower features at the short end of the spectrum (6.0, 6.8 and 7.6 m) were fitted in exactly the same way as before, only this time using a linear continuum through the points at 5.6 and 7.9 m. As absorption due to the silicates is approximately constant over this wavelength range this removed complications due to the quadratic continuum fit. When tested on spectra which were fitted well, such as those in Fig. 2, this produced results which were the same as before, within the fitting errors. Also, the manner in which the continuum was fitted often led to the CO2 feature fitting poorly, as seen in Fig. 3. As a result it was decided to measure the CO2 absorption in all the sources by fitting a linear continuum from 14.5-16.1 m, as the silicate absorption does not vary much over this range and the CO2 is not blended with any other spectral feature. This procedure was only applied to these narrower features, however, as the broad silicate, 11 and 13 m bands were deemed to be too heavily blended to deconvolve them meaning fully over a short section of the spectra. In addition, all features fitted in this manner are flagged as such in the results tables.
A total of 43 sources were detected: a complete list is presented in Table 1. The 6 pfov used here results in coordinates which are accurate only to - at best, so where sources have been identified with previous studies the coordinates presented are those from previous observations, where higher spatial resolutions or better S/N have resulted in more accurate positions. The coordinates evaluated here are only presented for sources which represent new or uncertain detections. In the 3 cases where sources have been identified as unresolved doubles the coordinates presented are those of the brighter source. Of the 43 sources detected 6 were not identified with detections from previous surveys (see Sect. 3.1.1) and 2 (Ser A 7 and Ser B 11) were determined to be different images of the same source, as the two Serpens fields overlap slightly. Consequently we were able to obtain complete 5.0-16.3 m spectra for a total of 40 sources, along with a further 2 partial spectra: a complete atlas of these 42 spectra is presented in Appendix A.
As noted by Blommaert et al. (2001) ghost images and stray-light are a significant problem when using the CVF, so care must be taken when identifying sources. Ghost images arise in two ways and are found either around the true image or directly opposite it (relative to the optical axis). Potential ghost images were identified by looking at the spectra of the objects, their flux densities and their positions relative to the instrument optics. In this manner 7 objects initially identified as sources were reclassified as ghost images (4 in Oph E, 2 in Cha I and 1 in RCrA). This leaves 6 "new'' detections in our survey, which we now address:
CK2 is the only one of these 3 sources which has been identified as a background field star (probably a background supergiant, Casali & Eiroa 1996). Unfortunately it fell directly on the bad column in the array here and so no useful spectroscopic data regarding CK2 was obtained. Both of the other two sources in Serpens show extremely deep absorption features due to ices and silicates and both have very red SEDs. Taking the K-band magnitudes of Giovannetti et al. (1998) and assuming the SEDs of these stars to be Rayleigh-Jeans spectra ( ) results in predicted mid-IR flux densities of 0.1 mJy for both sources, approximately 1000 times less than what is observed. Consequently it seems unlikely that these are background sources, so we interpret them to be deeply embedded objects: with this interpretation these 2 sources are 2 of the 3 most deeply embedded objects in this survey.
The spectra of the 43 sources identified in Table 1 were fitted using the method described above: the results of this procedure are presented in Table 2 and the fits and continua obtained are presented in Appendix A. The fitting procedure is both robust and unambiguous, converging to a good fit in most cases, but a few weaknesses exist. Firstly, as discussed in Sect. 2.3, the procedure only measures the 7 spectral features observed in almost all of the sources and consequently does not fit some rarer features, such as an apparent 14 m absorption band observed in source Oph A 2. Further, the profile fitted to each feature (except the CO2) is not allowed to vary from source to source. While this is broadly valid, a few exceptions led to some poor fits, such as the broadened 6.0 and 6.8 m bands in SVS2 (see Fig. 3), or a handful of sources which show broadened silicate features. However, given the low spectral resolution of the CVF
|Figure 4: The spectrum of HH100-IR, with the fitted spectrum (solid line) and continuum (dashed line). Note the poor fit to the CO2 feature at 15 m as a further example of why the CO2 feature was fitted independently.|
and the variation in the spectra over the large number of sources detected, this was found to be the most reliable and consistent method of spectral fitting.
In order to check the instrumental calibration and the validity of this fitting method, comparisons were made with previous observations of the bright, well-studied object HH100-IR (RCrA 5, see Fig. 4). Whittet et al. (1996) find a silicate optical depth of
which is somewhat less than the value of
measured here: this is almost certainly due to the fact that the iterative fitting procedure used here allows the continuum to drift away from the observed spectrum slightly. As discussed in Sect. 2.4, this is due to absorption across the entire wavelength range and is a benefit made possible by the broad wavelength coverage of the CVF. Whittet et al. (1996) also find a H2O ice column density of
cm-2, based on observations of the 3 m stretching mode. Adopting a band strength of
cm molecule-1 (Gerakines et al. 1995), we evaluate the column density N as:
The observed equivalent widths of the 6 m bending and 13 m libration modes of water ice showed no significant correlation at all. Further, the libration mode was not observed to correlate with any of the measured features. Whilst there is evidence for this in the literature (e.g. Bowey et al. 1998) the most likely explanation is that the libration mode is poorly fitted, due to the wide variation in possible profiles. The peak of this band has previously been found at 11 m in crystalline water, 12.5 m in amorphous water and at even longer wavelengths in mixtures with other molecules (Hagen et al. 1983; d'Hendecourt & Allamandola 1986; Cox 1989). Here a single profile was used (that of amorphous ice) and this probably resulted in poor fitting. Unfortunately this section of the spectrum is strongly blended with the silicate feature, so this problem cannot be remedied without making further ad hoc assumptions about the silicate profile(s).
|Figure 5: The strength of the unidentified feature plotted against the bending modes of both water (top) and CO2 (bottom) ices.|
The depth of the measured 11.2 m feature was found to correlate strongly with the depth of the silicate feature, with a direct (negative) proportionality providing a good fit to the data. Consequently the 11.2 m feature was interpreted to be an emissive shoulder on the silicate feature, rather than an independent feature due to another species: the presence of this feature narrows the silicate absorption profile slightly. The silicate profile has previously been found to vary depending on the composition and structure of the silicate grains (e.g. Demyk et al. 2000), so such a shoulder is not unexpected. However, 3 of the sources (RCrA 1, Ser A 6 and Cha I 3) show a significant 11.2 m emission feature without any significant silicate absorption feature: these features may be attributable to emission from crystalline silicates (Bregman et al. 1987; Campins & Ryan 1989).
As mentioned in Sect. 2.4, the CO2 ice profile is a very sensitive diagnostic of the ice environment. Whilst the spectral resolution of the CVF was not sufficient to study the profiles in great detail, it is worth noting that a significant long wavelength wing was present in almost all the observed absorption profiles. This would seem to indicate that a large fraction of the CO2 ice observed exists in a polar (H2O-rich) phase, which is characterised by this long wavelength wing (Gerakines et al. 1999). However quantifying the relative abundances of the different phases was not possible at this low spectral resolution.
|Figure 6: Examples of observed silicate profiles: emission profiles are generally broader than those seen in absorption and composite emission/absorption profiles are also seen. "Type'' refers to the classification scheme described in Sect. 3.2.5.|
|Figure 7: Composite silicate profiles fitted to the spectrum of SVS20. The solid line shows a composite fit with and , the dashed line a composite fit with and (and a different continuum). The single profile fit reported in the results table, with , is shown as the dotted line.|
It should be noted that while it is possible to fit such profiles to these sources, they are not well constrained. As can be seen by comparing the 2 composite profiles in Fig. 7 there is a degeneracy between the strength of this "flat-topped'' silicate profile and the continuum strength, and a great deal of information is assumed about the silicate profiles themselves. Consequently this investigation was not pursued for all the sources. Where the values in Table 2 represent the correct optical depth at 9.7 m, but are usually in error towards the wings of the silicate feature. Such values are useful in studies of general trends, but they do not accurately describe the entire silicate feature: a single number cannot fully describe the complicated nature of these features.
This also provides more direct evidence, similar to that of Whittet et al. (1988), that the extensive scatter in measurements of the ratio towards YSOs may be caused by complications due to silicate emission. Direct measurements of the silicate optical depth will significantly under-estimate the depth of foreground material in cases (such as SVS20) where optically thin silicate emission is also present. Rieke & Lebofsky (1985) derive , whilst noting the problem posed by "intrinsic'' silicate emission, and such emission does indeed lead to errors when evaluating AV from the silicate depth. As an example, the ISOCAM survey of Bontemps et al. (2001) derives values of for GY252 ( Oph E 1) and 24 for GY262 ( Oph E 3); the X-ray survey of Imanishi et al. (2001) derives slightly lower values of 24 and 18 respectively. However GY252, as seen in Fig. 6, clearly shows a "composite'' silicate profile whereas GY262 does not, and the respective silicate optical depths of 0.52 and 1.84 are clearly inconsistent with a single ratio. Unless the emitting and absorbing components of such silicate features can be separated unambiguously obtaining a single ratio will not be possible.
Comparing the index-based spectral classes to the spectral feature grouping above, we see that the majority of the class I sources belong to group a) and the majority of the class II sources to group b). However it is notable that a significant number (6) of the 23 group a) objects are either class II or transition objects. This in keeping with the convention that class I objects are the youngest (Lada 1987), but it may be significant that there are two distinct types of spectrum which both result in class II SEDs. Objects showing strong foreground absorption and objects showing little foreground absorption can both present class II continuum SEDs. This could be due to geometry of the YSOs relative to the absorbing clouds, or could be an effect intrinsic to the sources: spectroscopy along a pencil-beam cannot distinguish these. Whilst a larger sample size would constrain this problem further, it does seem that mid-IR observations alone do not provide strong constraints on the "true'' spectral class of YSOs.
In general, the most striking regional variation is between the sources in Cha I and those in the other three regions. All of the sources in Cha I, except for ISO-ChaI 192 (Cha I 2), show little or no absorption due to volatile ices and show silicates either in emission or an emission/absorption composite. As it lay outside the CVF1 images, due to rotation of the spacecraft between the two sets of observations, we only have a 9.3-16.3 m spectrum for ISO-ChaI 192. However it shows CO2 ice absorption and also appears to show deep silicate absorption ( -4), so it is probably an embedded object. Persi et al. (1999) found that ISO-ChaI 192 lies in a dense core and that it is probably the engine for the bipolar outflow detected by Mattila et al. (1989). Jones et al. (1985) found that several other YSOs have formed around the edge of this core, rather than at its centre: these were apparently formed during a burst of star formation triggered by a wind from the nearby star HD97300. Our results are consistent with this: ISO-ChaI 192 appears to be heavily embedded, and has a very red SED. The rest of the surrounding objects show similar spectral characteristics, and appear to be surrounded by far less circumstellar/foreground material than ISO-ChaI 192.
12 of the 20 objects in Serpens show deep absorption features, 7 show composite silicate profiles, and one (SVS2) shows silicate emission. It is clear that Serpens contains a wide range of YSOs, ranging from the most heavily embedded object observed (Ser B 11) to a source showing silicate emission. By contrast almost all of the sources in both RCrA and Oph show deep absorption features, with only 1 of the 7 sources in RCrA and only 2 of the 10 sources in Oph showing composite silicate profiles. Given that YSOs are generally expected to "sweep out'' their circumstellar material as they evolve (Shu et al. 1987), this seems to imply that the observed star formation is at a similar evolutionary stage in both RCrA and Oph, probably an earlier stage than that in Serpens. However, given the broad variation in the sources observed in Serpens, this conclusion remains somewhat tentative. It should also be noted that the observed fields (of ) are much smaller than the star-forming clouds, which typically extend over several square degrees of the sky, so the populations observed here may not be representative of the entire clouds.
|Figure 8: The strength of the bending mode of water ice plotted against the silicate optical depth. Different regions are distinguished by different symbols, as shown in the legend.|
|Figure 9: The ratio of the CO2:H2O ice equivalent widths plotted against silicate optical depth. This ratio is directly proportional to the ratio of ice column densities: adopting the band strengths from Gerakines et al. (1995) results in an equivalent width ratio of 1 corresponding to a column density ratio of 0.17.|
Figure 9 shows the ratio of the CO2:H2O ice equivalent widths (which is proportional to the corresponding ratio of column densities) plotted against the silicate optical depth. The results from Serpens and RCrA are fairly well correlated, with CO2:H2O ice column density ratios ranging from 0 to 0.16 (adopting the band strengths from Gerakines et al. 1995). There is a possible trend showing the CO2:H2O increasing with decreasing silicate optical depth, but this is somewhat unclear in these data.
The data from Oph, however, show a far greater scatter with CO2:H2O ice column density ratios ranging from 0 to 0.4. This is consistent with previous observations: for example Chiar et al. (1994, 1995) found a greater scatter in column densities of CO and H2O ices against AV in Oph than in Serpens, RCrA or Taurus. The reasons for this are unclear but it seems likely that, as suggested by Chiar et al. (1995), local conditions around the young stars in Oph play a role. Detailed study of the gas-phase abundances of CO2 and H2O towards these objects will further constrain this problem, but without such observations or further knowledge of the envelope structures we merely note the presence of this discrepancy here.
We thank Tom Greene for his work in planning the observations and for helpful comments. We thank Andy Longmore for useful comments about the manuscript and also thank Sylvain Bontemps for useful discussions. Part of this work formed part of the MPhys dissertation of RDA at the University of Edinburgh, and RDA also thanks the UKATC for funding a vacation studentship. Finally, we thank an anonymous referee for helpful advice which greatly improved the clarity of the paper.
|RCrA 1||19 01 41.5||-36 58 29||IRS2, TS13.1||2|
|RCrA 2||19 01 41.5||-36 58 59||17|
|RCrA 3||19 01 41.7||-36 59 54||H2||1|
|RCrA 4||19 01 48.0||-36 57 19||IRS5, TS2.4||2|
|RCrA 5||19 01 50.7||-36 58 07||HH 100-IR, IRS1, TS2.6||2|
|RCrA 6||19 01 50.9||-36 57 37||17||Very close to RCrA|
|RCrA 7||19 01 52.4||-36 57 37||17||Very close to RCrA|
|Oph A 1||16 26 17.3||-24 23 49||SKS9, ISO-Oph 21||6|
|Oph A 2||16 26 18.8||-24 24 21||SKS11, ISO-Oph 26 ???||17||Uncertain ID|
|Oph A 3||16 26 21.5||-24 23 07||GSS30, GY6, SKS12, ISO-Oph 29||6|
|Oph A 4||16 26 23.7||-24 24 40||GY21, SKS16, ISO-Oph 37||6||Falls on bad column|
|Oph E 1||16 27 21.6||-24 41 43||GY252, SKS36, IKT43, ISO-Oph 132||6|
|Oph E 2||16 27 24.8||-24 41 02||ISO-Oph 137||6|
|Oph E 3||16 27 26.9||-24 39 23||GY262, IKT53, ISO-Oph 140||6|
|Oph E 4||16 27 27.1||-24 41 31||17||Very faint|
|Oph E 5||16 27 27.2||-24 40 49||GY265, IKT54, ISO-Oph 141||6|
|Oph E 6||16 27 28.2||-24 39 32||GY269, IKT57, ISO-Oph 143||6|
|Ser A 1||18 29 56.6||+1 12 56||SVS4/2, GEL5||11||Very faint|
|Ser A 2||18 29 56.7||+1 12 35||SVS4/3||11||Very faint|
|Ser A 3||18 29 57.5||+1 12 57||SVS4/5 and||11||Unresolved double|
|Ser A 4||18 29 57.9||+1 12 25||SVS4/7||11||Very faint|
|Ser A 5||18 29 57.9||+1 12 34||SVS4/8||11|
|Ser A 6||18 29 57.9||+1 12 43||SVS4/9 and||11||Unresolved double|
|Ser A 7||18 30 00.1||+1 13 03||GCNM130||13||Off frame 7.0-9.3 m|
|Ser B 1||18 29 52.9||+1 14 55||GCNM53||13|
|Ser B 2||18 29 55.7||+1 14 31||CK9, GEL4, EC74, GCNM76||13|
|Ser B 3||18 29 56.8||+1 14 46||SVS2, CK3, GEL6, EC82, GCNM87||13||Unresolved double|
|and GEL8, EC86, GCNM93|
|Ser B 4||18 29 57.2||+1 13 28||17||Peak in nebulosity?|
|Ser B 5||18 29 57.6||+1 15 31||GCNM100||13|
|Ser B 6||18 29 57.7||+1 14 05||SVS20, CK1, GEL10, EC90, GCNM98||13||Double Object|
|Ser B 7||18 29 58.1||+1 13 24||17||Peak in nebulosity?|
|Ser B 8||18 29 58.2||+1 15 21||CK4, GEL12, EC97, GCNM106||13||Falls on bad column|
|Ser B 9||18 29 58.7||+1 14 26||EC103, GCNM112||13|
|Ser B 10||18 29 59.2||+1 14 08||CK8, GEL13, EC105, GCNM119||13|
|Ser B 11||18 30 00.1||+1 13 03||GCNM130||13|
|Ser B 12||18 30 00.5||+1 15 20||CK2, EC118, GCNM136||13||Peaks on bad column|
|Ser B 13||18 30 02.1||+1 13 59||CK7, EC125, GCNM154||13||Very faint|
|Cha I 1||11 09 23.3||-76 34 35||C1-6, ISO-ChaI 189, KG82||15||Bad column in CVF1|
|Cha I 2||11 09 29.2||-76 33 30||ISO-ChaI 192, KG87||15||Out of frame in CVF1|
|Cha I 3||11 09 42.6||-76 35 01||C1-25, ISO-ChaI 199, KG93||15|
|Cha I 4||11 09 46.9||-76 34 49||C1-24, ISO-ChaI 204, KG97||15|
|Cha I 5||11 09 54.1||-76 34 26||C1-5, ISO-ChaI 223, KG109||15|
|Cha I 6||11 10 00.8||-76 34 59||WW Cha, C1-7, ISO-ChaI 231, KG116||15|
|8 m||8 m||Equivalent Line Widths (nm)||Source||Cont.|
|Source||flux||continuum||6 m||6.8 m||7.6 m||13 m||15.2 m||Group||Notes|
|RCrA 1||7.84(16)||7.79(16)||-0.02(1)||-0.06(1)||178(8)||17(1)||17(3)||247(25)||168(7)||0.56||0.58||b||Strong 11 m emission|
|2||0.95(2)||1.24(3)||0.72(3)||-0.27(4)||203(8)||57(4)||17(3)||0(0)||83(4)||-0.96||-1.35||a||Silicate too wide at long|
|, probable 13 m error|
|3||0.27(2)||0.28(2)||0.22(1)||-0.16(2)||147(15)||14(3)||0(0)||718(76)||0(0)||-1.11||-0.96||b||probably too small,|
|possible absorption bands|
|at 13 m, 14.25 m|
|4||1.74(110)||2.28(144)||0.71(3)||-0.07(1)||271(10)||81(6)||6(1)||586(59)||253(11)||1.23||1.43||a||Possible 10.5 m emission|
|6||0.94(3)||1.36(4)||0.95(4)||-0.15(2)||204(7)||61(4)||11(2)||73(7)||97(4)||-0.40||-0.42||a||Under-estimates 13 m|
|7||1.20(27)||1.62(37)||0.74(3)||-0.08(1)||191(9)||86(6)||18(3)||0(0)||95(4)||-0.18||-0.37||a||Possible 10.5 m emission,|
|under-estimates 13 m|
|Oph A 1||0.42(2)||0.85(4)||2.11(11)||-0.27(4)||27(3)||72(7)||19(3)||927(94)||57(3)||-0.70||-0.20||a||Poor 5-8 m fit|
|2||0.21(2)||0.54(6)||2.68(19)||-0.70(10)||207(14)||253(53)||79(9)||2611(319)||0(0)||0.34||1.36||a||Definite 14 m band|
|3||6.03(7)||8.99(10)||1.05(4)||-0.20(3)||107(6)||0(0) ??||0(0)||10(3)||87(3)||1.61||1.84||a||Possible 7.1 m band|
|4||2.33(3)||2.74(4)||0.49(2)||-0.16(2)||0(0)||25(12)||27(4)||687(69)||23(1)||0.95||1.15||b||Silicate fit too wide|
|Oph E 1||2.88(3)||3.36(4)||0.52(2)||-0.21(3)||140(5)||60(4)||27(4)||1206(121)||98(4)||0.15||0.43||b||Silicate too wide|
|2||0.23(2)||0.56(5)||2.55(12)||-0.36(5)||255(16)||93(7)||53(8)||1881(196)||324(20)||-0.26||0.51||a||Deep 7.6 m band|
|3||0.34(38)||0.65(72)||1.84(12)||-0.27(4)||131(10)||134(13)||15(2)||1536(160)||155(7)||-0.06||0.47||a||Very deep 6.8 m band|
|4||0.10(2)||0.17(4)||1.42(8)||-0.23(4)||291(41)||0(0)||29(5)||1561(279)||501(49)||-0.55||-0.01||a||Very noisy: probably only|
|class, CO2, good data|
|Ser A 1||0.06(1)||0.15(3)||2.14(31)||-0.60(9)||845(849)||356(161)||0(0)||2386(311)||298(33)?||-0.24||0.47||a||Noisy and faint, possible|
|12, 14 m bands|
|2||0.06(1)||0.17(3)||2.39(61)||-0.76(10)||965(878)||298(130)||0(0)||2679(437)||0(0)||0.43||1.35||a||Noisy, possible absorption|
|bands at 11-16 m|
|3||0.64(2)||1.24(4)||1.70(7)||-0.32(4)||523(20)||165(11)||19(3)||1633(163)||285(13)||1.35||1.91||a||Possible narrow absorption|
|band at 9.7 m - CH3OH?|
|4||0.07(1)||0.15(2)||2.14(34)||-0.46(7)||847(688)||271(69)||0(0)||2426(552)||269(51)?||0.42||1.17||a||Noisy, possible absorption|
|bands at 11-16 m|
|5||0.17(1)||0.30(2)||1.42(12)||-0.37(5)||435(26)||165(11)||8(1)||1538(177)||186(11)||0.25||0.71||a||Possible 12, 14 m bands|
|6||0.46(1)||0.45(1)||-0.18(1)||-0.22(3)||272(17)||97(6)||0(0)||803(80)||223(9)||0.44||0.50||b||Strong 11 m emission|
|Classification scheme proposed in Sect. 3.2.5.|
|Short wavelength features fitted using linear continuum, fitted from 5.6-7.9 m, in cases where the automatic fitting procedure produced errors (see Sect. 2.5).|
|Feature fitted manually as spectrum was incomplete.|
|Source fell on bad column on array: data from these sources is not considered to be especially accurate.|
|No K-band data: evaluated using m.|
|8 m||8 m||Equivalent Line Widths (nm)||Source||Cont.|
|Source||flux||continuum||6 m||6.8 m||7.6 m||13 m||15.2 m||Group||Notes|
|Ser B 1||0.19(1)||0.57(3)||2.76(16)||-0.39(5)||703(44)||215(31)||0(0)||962(99)||267(14)||2.02||2.63||a||Very deep absorption features|
|2||0.19(1)||0.16(1)||-0.41(2)||-0.23(3)||91(5)||39(3)||25(4)||0(0)||67(10)||0.14||0.04||b||Large discontinuity at join,|
|poor fit from 8-12 m|
|3||1.37(2)||0.77(2)||-1.51(5)||-0.10(1)||208(16)||33(2)||43(6)||1198(121)||202(8)||1.14||1.05||c||Strong silicate emission,|
|fit too narrow|
|4||0.33(1)||0.28(1)||-0.37(1)||-0.19(3)||154(9)||45(4)||32(5)||0(0)||102(5)||-1.03||-1.14||b||Poss. 13 m absorption band|
|5||0.10(1)||0.08(1)||-0.57(3)||-0.22(3)||208(8)||57(4)||29(4)||0(0)||59(4)||0.06||-0.09||b||Long fit probably wrong|
|6||6.16(14)||5.31(13)||-0.33(1)||-0.12(2)||137(5)||38(2)||24(4)||0(0)||103(4)||-0.11||-0.18||b||Poor 8-11 m fit|
|7||0.35(1)||0.32(1)||-0.14(1)||-0.15(2)||152(7)||50(4)||23(3)||0(0)||100(4)||-1.13||-1.10||b||Large discontinuity at join|
|9||0.35(1)||0.53(2)||1.02(4)||-0.20(3)||237(8)||38(4)||25(4)||1119(114)||144(6)||0.99||1.33||a||Dubious 6.8 m and 7.5 m fits|
|10||0.36(1)||0.33(1)||0.04(1)||-0.06(1)||168(11)||35(3)||23(3)||609(61)||126(6)||-0.43||-0.34||b||Probably peak in nebulosity|
|11||0.36(2)||1.89(10)||4.34(31)||-0.62(8)||984(45)||350(67)||73(11)||1996(202)||479(28)||2.07||3.14||a||Very deep absorption features|
|12||-||-||-||-||414(76) ?||0(0) ?||19(5) ?||-||143(16)?||-||-||-||Very dubious measurements:|
|peaks on bad column|
|13||0.04(1)||0.07(2)||1.35(22)||-0.20(3)||348(51)||146(46)||69(11)||2496(291)||185(7)||0.55||1.09||a||Good fit but very faint|
|Cha I 1||0.71(2)||0.65(2)||-0.18(1)||-0.07(1)||0(0)||0(0)||7(1)||548(55)||0(0)||-0.32||-0.28||b||Poor fit, due to large|
|discontinuity at join|
|3||0.14(1)||0.13(1)||-0.37(2)||-0.07(1)||144(10)||36(3)||10(2)||0(0)||0(0)||-0.40||-0.52||b||Noisy, poor fit at long|
|4||0.04(1)||0.02(1)||-1.88(11)||0.13(2)||225(25)||87(9)||0(0)||0(0)||0(0)||-0.65||-1.09||c||Very faint, low S/N|
|5||3.58(6)||3.59(6)||0.00(1)||-0.12(2)||41(2)||0(0)||7(1)||283(28)||0(0)||-0.43||-0.37||b||Strong 11 m emission|
|6||3.18(6)||1.93(4)||-1.32(5)||-0.04(1)||0(0)||0(0)||0(0)||0(0)||0(0)||-0.85||-1.07||c||Silicate fit too narrow|
|] Classification scheme proposed in Sect. 3.2.5.|
|] Short wavelength features fitted using linear continuum, fitted from 5.6-7.9 m, in cases where the automatic fitting procedure produced errors (see Sect. 2.5).|
|] Feature fitted manually as spectrum was incomplete.|
|] Source fell on bad column on array: data from these sources is not considered to be especially accurate.|
|] No K-band data: evaluated using m.|
|Figure A.2: Spectra of RCrA 7, Oph A 1-4 and Oph E 1. Note that source Oph A 4 fell on the bad column on the array.|
|Figure A.4: Spectra of Ser A 2-6 and Ser B 1. The partial spectrum of Ser A 7 is omitted as a complete spectrum of the same object was found as Ser B 11. Note also that both Ser A 3 and 6 were found to be unresolved doubles.|
|Figure A.6: Spectra of Ser B 8-13. Source Ser B 8 fell on the bad column of the array, so the data are somewhat dubious. Ser B 12 peaked on the bad column and the resulting spectrum is very poor, so no fit was possible.|
|Figure A.7: Spectra of Cha I 1-6. Cha I 1 fell on the bad column in CVF1, resulting in the large discontinuity seen at the join. Due to the rotation of the spacecraft Cha I 2 was only observed in CVF2.|