Gas and dust in the star-forming region rho Oph A: II. The gas in the PDR and in the dense cores
Key Words.:interstellar medium (ISM): general – interstellar medium: individual objects: – interstellar medium: molecules – interstellar medium: abundances – interstellar medium: photon-dominated region (PDR) – Stars: formation
Context:The evolution of interstellar clouds of gas and dust establishes the prerequisites for star formation. The pathway to the formation of stars can be studied in regions that have formed stars, but which at the same time also display the earliest phases of stellar evolution, i.e. pre-collapse/collapsing cores (Class -1), protostars (Class 0), and young stellar objects (Class I, II, III).
Aims:We investigate to what degree local physical and chemical conditions are related to the evolutionary status of various objects in star-forming media.
Methods: displays the entire sequence of low-mass star formation in a small volume of space. Using spectrophotometric line maps of H, HO, NH, , O, O I, CO, and CS, we examine the distribution of the atomic and molecular gas in this dense molecular core. The physical parameters of these species are derived, as are their relative abundances in . Using radiative transfer models, we examine the infall status of the cold dense cores from their resolved line profiles of the ground state lines of HO and NH, where for the latter no contamination from the VLA 1623 outflow is observed and line overlap of the hyperfine components is explicitly taken into account.
Results:The stratified structure of this photon dominated region (PDR), seen edge-on, is clearly displayed. Polycyclic aromatic hydrocarbons (PAHs) and O I are seen throughout the region around the exciting star S 1. At the interface to the molecular core 0.05 pc away, atomic hydrogen is rapidly converted into H, whereas O I protrudes further into the molecular core. This provides oxygen atoms for the gas-phase formation of O in the core SM 1, where (O). There, the ratio of the O to HO abundance [(HO)] is significantly higher than unity. Away from the core, O experiences a dramatic decrease due to increasing HO formation. Outside the molecular core , on the far side as seen from S 1, the intense radiation from the 0.5 pc distant early B-type star HD 147889 destroys the molecules.
Conclusions:Towards the dark core SM 1, the observed abundance ratio (O)/(HO), which suggests that this object is extremely young, which would explain why O is such an elusive molecule outside the solar system.
|Species and||Frequency||Wavelength||Energy||Beam width||Main beam|
|[O I] ( - )||4744.738586||63||228||7||0.73|
|( - )||2060.052199||145||326||12||0.54|
PCalSpectrometer_Beam_v6: for 62, 45 and 168/187 m spatial averages are shown.
Mueller et al. (2014), Release Note #1, HIFI-ICC-RP-2014-001, v1.1.
The physics and chemistry of star-forming clouds are governed by the competition between the energy input into the cloud and the energy lost through the radiation from the gas and the dust. This ability to cool efficiently is paramount during the gravitational formation of stars and planets and, on a grander scale, of galaxies. The abundance of the major coolants is determined by feedback processes between the atoms, molecules, and dust grains that are present in the media. However, many of these processes are only partially understood and are the focus of intense contemporary research. In this paper, we attempt to disentangle the various interrelations among key atomic and molecular species by analysing their distribution within a star-forming cloud.
We chose as a prime example the nearest region of star formation, i.e. the molecular cloud L 1688 at a distance of 120 pc (Loinard et al., 2008). Major properties of the dark clouds in the constellation of Ophiuchus have been reviewed by Wilking et al. (2008) and, more recently, by White et al. (2015). A dense stellar cluster has been forming over the past one to two million years (Bontemps et al., 2001; Evans et al., 2009; Ducourant et al., 2017), transforming cloud matter into a young stellar population at an overall efficiency of a few percent, but with local excursions to a few tens of percent (Wilking & Lada, 1983; Liseau et al., 1995, 1999; Bontemps et al., 2001; Evans et al., 2009). Several intensity maxima in DCO line emission were named ‘cores’ and assigned the capital letters A through F by Loren et al. (1990). Since then, this scheme has been extended (e.g. White et al., 2015; Punanova et al., 2016).
The dense core attracts our particular attention as it harbours several manifestations of star formation within a relatively small volume (length scale approximately 0.1 pc), namely a number of pre-stellar cores, a Class 0 source with its highly collimated molecular jet, and Class I/II sources with their outflows and HH-objects (Fig. 1 and Liseau et al., 2015). In addition, is prominent for its unique chemistry, including that of O, HO, and HO. These were first detected outside the solar system by our group (Larsson et al., 2007; Liseau et al., 2012; Bergman et al., 2011a; Parise et al., 2012).
We have studied from space and the ground in a variety of spectral lines and in the continuum. The first results, focussing on the dust properties and the gas-to-dust relationship, are presented in Paper I of this series (Liseau et al., 2015). The subregions of that were observed with various instruments aboard Herschel (Pilbratt et al., 2010) are outlined in Figure 1, superposed onto an image in the CO (3-2) line; in this paper, we investigate further the complex physical and chemical interplay between the gas and the dust in , and what their relation is with regard to star formation per se.
In Sect. 2, the observations and the reduction of the data are described. In Sect. 3, the spatially resolved low-() photon dominated region (PDR) is introduced. In Sect. 4 these data are compared with theoretical PDR models, and the O formation in and possible scenarios of observed protostellar mass infall are discussed. In Sect. 5 our main conclusions of this paper are summarized, and in Sect. 6 a synergy of the results and conclusions of Papers I and II are put into the wider context of star formation. Finally, our Odin mapping observations in the HO line of , B1, B2, C, D, E, and F are presented in Appendix A.
2 Observations and data reduction
2.1 HO, Nh, and
The cloud was observed in two HO lines with the Heterodyne Instrument for the Far Infrared (HIFI, de Graauw et al., 2010) on board the ESA spacecraft Herschel (Pilbratt et al., 2010) carrying a 3.5 m telescope. The 557 GHz data were acquired during OD 475, 1226, and 1231 in mixer band 1b and the 1670 GHz (179.5 m) data on OD 1221 in band 6b. Data were taken simultaneously in the H- and V-polarizations using both the Wide Band Spectrometer (WBS, accousto-optical) with 1.1 MHz resolution and the High Resolution Spectrometer (HRS, correlator) with 0.25 MHz resolution. The Band 1 receiver was tuned such that rotational transitions of ammonia and diazenylium, i.e. both the NH 572 GHz line and the (6-5) 560 GHz transition, were also simultaneously admitted. The map and observational details for diazenylium can be found in Paper I.
For the programme GT2_rliseau_1, the mapping was done on the fly (OTF) with a reference position at (, ) on OD 475, 1226, and 1231. At 557 GHz, the map size was 8 65 with samplings in 16 steps. At 1670 GHz (OD 1221), the map was 25 1 with a sampling rate of 20. The observing times were 7 and 3 hours, respectively. All data have been reduced with the HIFI pipeline HIPE version 14.1.0. Instrument characteristics are provided by Roelfsema et al. (2012) and some of the observational parameters can be found in Table 1.
The NH and (6-5) lines were simultaneously observed with the HO () transition. These rectangular maps were all observed at an angle of 55 with respect to the equatorial coordinate system and had to be rotated for proper alignment. The rotation and regridding data can be found in Bjerkeli et al. (2012) and in Paper I.
Following up on our initial O observations with HIFI (Liseau et al., 2012), we obtained data in the O () 487 GHz line towards 12 new positions on OD 1357, 1358, 1367, and 1371 (OT2_rliseau_2, 38.5 hr integration time). In addition, the O () 773 GHz line in HIFI Band 2 was observed towards HH 313 A, an optically visible shocked region of the VLA 1623 outflow (8.5 hr integration on OD 1384). The reduction of these new data also followed the procedures outlined in the 2012 paper.
All in all, in we observed O 487 GHz towards 15 positions and O 773 GHz towards five highly oversampled positions, each corresponding in size to that of the 487 GHz data. The 773 GHz Herschel beam was 27 and the grid spacing 10.
2.3 O I
Observations of the fine structure components of O I at 63 m (-) and 145 m (-) of the inverted ground state with the ISO-LWS (Long Wavelength Spectrometer, Clegg et al., 1996) towards the VLA 1623 outflow have been presented earlier by Liseau & Justtanont (2009). On OD 1202, the Photodetector Array Camera and Spectrometer (PACS, Poglitsch et al., 2010) was used to map the fine structure line emission north of the outflow (see Fig. 1) as part of the programme GT2_rliseau_2 and with an observing time of 2 hr 16 min. The map consists of PACS rasters separated by 38. An individual PACS raster covers 47 47 and is composed of spaxels (spatial pixels), providing a spatial sampling of 94 per spaxel. At 63 m, the spectral resolution is about 3500, and at 145 m 1000, corresponding roughly to 100 and 300 km s, respectively. The lines in are therefore expected to be unresolved, except perhaps towards the outflow.
The observing mode was unchopped grating scan and a position about 15 NE of the map was used as a reference. To estimate the absolute accuracy of the fluxes we compared data from one map position with those that we previously obtained with the ISO-LWS (e.g. Liseau & Justtanont, 2009). However, we refer here to re-reduced data using the final version of the pipeline (OLP 10). The LWS line and continuum fluxes were both measured using two methods: the default point-source calibration and by exploiting extended source corrections. To compare the PACS measurements with those with the LWS, the PACS data were convolved with a flattened top Gaussian beam (Larsson et al., 2000) and with a pure equivalent aperture measurement. The results show that these two different methods lead to intra-instrumental agreement within 10-30%.
2.4 H and polycyclic acromatic hydrocarbons
Towards , the Infrared Space Observatory (ISO) archive contains three frames of spectrophotometric data using the ISOCAM-CVF (Circular Variable Filter, Cesarsky et al., 1996). For one frame, the pixel field of view (FOV) is 300 and it images the region around the star GSS 30. The other two, with a pixel FOV of 600, cover the VLA 1623 outflow (Liseau & Justtanont, 2009) and the PDR of S 1, respectively. Together, these two frames cover most of the region. In addition, Spitzer data of the NW part of the VLA 1623 outflow have been presented by Neufeld et al. (2009).
Towards the S 1 PDR region, PAH emission totally dominates the spectrum over the wavelength range of the ISO-CVF (see also Justtanont et al., 2008). However, three pure rotational H lines, S(2), S(3), and S(5), could also be extracted. The S(5) line at 6.9 m is relatively unaffected by other spectral features, while the S(2) 12.3 m line is blended with a PAH feature and the S(3) 9.3m line is situated inside the deep silicate absorption. Consequently, the fluxes for these two lines are more uncertain. The observations and procedures to reduce the CVF data are described in detail by Liseau & Justtanont (2009). There, numerical model fits to the spectral PAH features were used to estimate the fluxes also of other H lines. Comparing these results with the present ones, where possible, demonstrates excellent agreement.
2.5 CS and CO
CS (10-9) data were obtained simultaneously with the O 487 GHz maps with HIFI. Complementary data had previously been collected with the 15 m Swedish ESO Submillimetre Telescope (SEST, Booth et al., 1989) in the CS transitions of (2-1), (3-2), and (5-4), in addition to pointed observations in the isotopologues CS and CS (see Appendix C).
Simultaneously with the O 773 GHz line, maps at 10 sampling of CO (7-6) and CO (7-6) were obtained (see Fig. 8 below).
3 Results: A spatially resolved PDR
The mapping in atomic and molecular line emission revealed a spatially resolved PDR, seen from the side (‘edge-on’).
3.1 Hydrogen, H
As can be seen in Figure 4, the H S(2) emission outlines a clear spherical shell structure of warm PDR gas around the stellar object S 1. In the figure this structure is also seen in the S(3) and S(5) lines, but at a slightly more fragmented level. This spherical structure motivates the radial averaging of the H fluxes and these are plotted in the lower rightmost frame of Fig. 4, together with a model fit using the Meudon PDR code
The observed data are consistent with gas at a temperature of about 10 K that has a total column density of a few times 10 cm, values that are identical to those determined by Liseau & Justtanont (2009) towards the VLA 1623 outflow, and where an H ortho-to-para ratio of 2 was derived.
3.2 Oxygen: O I and O
Atomic oxygen, O I
As in the H image, the VLA 1623 outflow is also seen in the upper part of the leftmost panel of Fig. 4, showing the distribution of the fine structure line emission of [O I] (-) 63 m. In the upper middle panel, the distribution of the (-) 145 m line, originating from a higher upper level
The upper right panel in Fig. 4 displays the line flux ratio . In PDRs, observed values are generally around 10 to 20 (Hollenbach & Tielens, 1999). It is therefore remarkable that some regions show ratios ten times lower, and even lower than unity. The region in the southeast where this occurs, i.e. N 6, is otherwise conspicuous only in emission (Di Francesco et al., 2004, and Paper I).
Line ratios of the magnitude observed here are characteristic of gas that is neither optically thin nor in thermodynamic equilibrium (e.g. Liseau et al., 2006; Canning et al., 2016). Oxygen fine structure lines from cold gas ( 20 K), which are optically thick in both ground state transitions, can show ratios below unity. A gas kinetic temperature of 17 K in this particular region was recently determined in Paper I. However, the N 6 column density, cm, is too low by a factor of 300 for the 145 m optical depth to reach unity (see Liseau et al., 2006) and must therefore be dismissed as an explanation. Instead, as advocated by these authors, a cold foreground cloud that absorbs much of the 63 m radiation but leaves the higher level 145 m line unaffected appears to be the only viable alternative to explain the very small 63 m to 145 m flux ratios of N 6.
Molecular oxygen, O
In Fig. 4 it can be seen that the atomic oxygen, O I, appears to be penetrating quite deeply into the dense parts of the dark core , outlined by the contours of the emission. It feeds the O production region (Fig. 3) with atomic oxygen. There, the molecular oxygen is produced in the gas phase, with additional contribution from the ice-crusted dust grains (Hollenbach et al., 2009). The models of Melnick & Kaufman (2015) for Orion A, which invoke shock chemistry, do not likely apply to as the O emission regions and the outflow/HH-shocks are spatially well separated. In , these shocks enhance the HO abundance, but not that of O.
Unfortunately, our O map is not complete and maximum emission occurs near the edge of the map. To remedy this, complementary observations would be required. However, after the decommissioning of Herschel, there is no other facility available or planned any time soon. The factor of two higher intensity at the local O emission maximum is likely an effect of the increased temperature (25 K) rather than a higher column density of O molecules.
3.3 Water, HO
Spectral components of the line profiles
The emission due to water is seen essentially everywhere in the molecular core , except in the easternmost parts towards the early-type stellar object S 1, where the water molecules have been destroyed by the radiation from the B-star.
The lines of the two ground state transitions
The HIFI observation of the higher excitation line HO () 1670 GHz is a completely new result, but with its three times smaller beam the map is not as extended as that for the 557 GHz line. However, our strip map of the dense core region of (Fig. 1) is still useful for a comparison of the two line profiles (see Figure 5). The 557 GHz line displays several components, whereas the 1670 GHz line is mostly in absorption with weak wing emission towards the dark cores SM 1 and SM 1N. This may seem surprising, as these positions are not anywhere near the strong outflow from VLA 1623 (Fig. 5), but could be due to the weak CO outflows of Kamazaki et al. (2003).
The most significant part of the line profile, though, is the deep absorption feature, reaching down to close the zero level. The same depth is also seen in the line, where the signal-to-noise (S/N) is significantly higher. There, however, the emission core varies appreciably in intensity with position. This is not an effect caused by a foreground absorber drifting over the cloud, thereby changing the strength of the emission peak, because the centre of the absorption dip in both lines is completely stable with respect to the systemic Velocity Standard of Rest ().
3.4 Ammonia, NH
The spatial distribution of the NH line is similar to that of HO (), both being of the ortho flavour. As for water and diazenylium (), ammonia peaks close to SM 1N. A clear difference is displayed by deuterated species, which all peak at the SM 1 position. For instance, formaldehyde has its emission maximum towards SM 1N, whereas its doubly deuterated form peaks towards SM 1 (Bergman et al., 2011b), i.e. DCO has essentially vanished towards SM 1N where the HCO emission is strongest. This difference may have its explanation in the fact that temperatures are lower at SM 1 (Paper I), promoting the exothermic process of molecular deuteration. However, Friesen et al. (2014) detected HD towards SM 1N, but not towards SM 1.
The ground state lines of both HO and NH have double-peaked line profiles. However, a remarkable difference between the ammonia and water lines are the high-velocity wings that are entirely absent in the ammonia data (Fig. 12). Consequently, low-amplitude velocity fields like those due to gravitational infall are potentially better traced in the NH lines.
3.5 Carbon monoxide, CO
In an adjacent shell, next to H, low- CO is distributed along a dense ridge that is broken up into several high-density clumps, of which SM 1 and SM 1N are the most conspicuous ones (Fig. 1).
Profile maps of CO (7-6) and CO (7-6) have been secured simultaneously with the observations of O () 773 GHz as prime. The positions are identified in Fig. 3 by the red rectangles. The profiles are fit by a two-component model, where Gaussians were fitted to the observations. Westward of the ridge, the isotopologues of CO decline rapidly.
3.6 Methylidine, CH, and ethynyl, CH
At the very edge of the PDR-interface, the () line of the cation CH has been mapped simultaneously with the HO () line at 179.5 m. In the left frame of Fig. 9, it can be clearly seen that the emission arises in front of the cores SM 1N, SM 1, and SM 2 (red dots from north to south) as seen from the B star S 1.
In contrast, the neutral CH radical is detected further in and closer to the cores, where the UV penetration into the cloud has diminished due to the increased extinction by the dust (Fig. 9, right-hand panel).
3.7 Carbon sulfide, CS
The CS maps, up to (), were obtained at the SEST (Appendix C) and the () data with HIFI aboard Herschel. These maps indicate, following the excitation gradient of the PDR, that CS appears behind CO and not predominantly in the densest parts of the cloud core.
4.1 The - Pdr
For the stellar parameters given in Table 5 of Paper I, and assuming solar elemental abundances, we use ATLAS 9 model atmospheres for the energy distribution of the stars (Castelli & Kurucz, 2004). At 0.05 pc from S 1, the FUV field
The extinction towards is well described by a law with = 5.5, so that the gas-dust relation reads (H I + H)(H I)(H) cm (Bohlin et al., 1978, see also Paper I).
4.2 The physical conditions in
A gaseous sphere around S 1
At distances greater than cm from the centrally positioned S 1, radiation pressure on the hydrogen gas is not very important. Due to the efficient braking by the 600 to 10 cm gas (seen in [O I] throughout this region, Fig. 4), radiation pressure in the Lyman lines can result in terminal velocities of only a few times 10 km s at the 0.05 pc interface
Without any significant dynamical pressure gradients, the nearly perfectly spherical shape of this region is kept in place by a delicate pressure balance, where the higher temperature of the H region is offset by the higher density of the molecular core, i.e. . In high-() PDRs the situation is different, where an H II region drives an ionization front into the molecular cloud (e.g. Tielens, 2005). A clumpy structure would further alter both molecule formation depths and, in particular, the radiative transfer.
Abundant O production in
Of considerable interest is to understand the formation history of O in as this molecule, in spite of numerous attempts (Goldsmith et al., 2000, 2002; Pagani et al., 2003; Yıldız et al., 2013; Sandqvist et al., 2015; Wirström et al., 2016), has hardly been found anywhere else outside the solar system. Taking into account new lab results for the oxygen binding energy on dust grains (He et al., 2015), Taquet et al. (2016) recently discussed a number of formation scenarios and argued in favour of a particular one that could fit the physical conditions in .
Taquet et al. (2016) suggested that the high abundance of O seen in the gaseous comae of two solar system comets, with (O)/(HO) = 0.01-0.1 (Bieler et al., 2015; Rubin et al., 2015), points towards a primordial origin, i.e. that this O was initially produced in the protosolar cloud core and subsequently transported through the viscous protoplanetary disc to its inner regions. It is argued that chemistry on icy dust grains would produce observed levels of related species, i.e., (HO)/(O) and (HO)/(O), see Bergman et al. (2011b) and Parise et al. (2012), respectively, and with the accompanying theoretical chemistry models presented by Du et al. (2012a, b).
However, we find it very difficult to reconcile the comet results for O and HO with our abundance determinations for , i.e. (O) = (Larsson et al., 2007; Liseau et al., 2012, this paper) and (HO) = (see Sect. 4.3.2), hence (O)/(HO) = 10.
In a model specifically designed for , Taquet et al. (2016, see their Fig. 6) find that (O)/(HO) during the time of about 5000 - 30 000 yr on their chemical clock, with a free-fall time of 16 000 yr (see below) falling right into this interval. If this chemical model is basically correct (there might be issues with their HO and HO abundances), the data would suggest the dense clumps in to be (chemically) very young indeed. As discussed by Liseau et al. (2012), the relative brevity of the period of abundant O in the gas phase would readily explain the elusiveness of this molecule in the interstellar medium.
4.3 Theoretical models of infall: HO
The HIFI instrument aboard Herschel is based on the heterodyne technique that implies high spectral resolution, . At the operating frequencies of HIFI, radial velocity resolutions are fractions of a km s, sufficient to spectrally resolve narrow molecular lines that originate in the cold interstellar medium.
Line profiles: HO ()
Due to their expected high optical depths (), the ground state lines of water are particularly suited as tracers of protostellar infall, as these large optical depths lead to large contrasts between receding and approaching emission. The high optical depths, in combination with the complex system of its many energy levels of the water molecule, imply however that the radiative transfer is difficult and the proper computation is therefore often avoided. A widely used method is to fit the observed profiles, essentially ‘perfectly’, with multi-component Gaussian functions; however, this procedure does not appropriately recover the physics and so we turn to other options.
We compute the line profiles from a physical model and account for the radiative transfer using an Accelerated Lambda Iteration (ALI) code; at these high optical depths other possibilities such as Monte Carlo methods frequently run into conversion problems. The benchmarking of the ALI code has been described by Maercker et al. (2008). In the present work, the collisional rate constants of Faure et al. (2007) are used.
Profiles of the infall centre
The observed 557 GHz line profile, with its blue-red asymmetry and deep central absorption, is highly reminiscent of theoretical profiles of spherical infall, i.e. from a protostar at the onset of gravitational collapse prior to observable disc formation (e.g. Ashby et al., 2000).
Because of its conceptual and computational simplicity (see e.g. Foster & Chevalier, 1993), our protostar model is a Bonnor-Ebert sphere (BE, Bonnor, 1957; Ebert, 1957); based on observed criteria (Paper I), it is found to be gravitationally unstable (dimensionless mass
The line profiles to be compared to the observed ones are computed self-consistently, based on observation, and with the following ALI parameters: to 15 K, to 0.6 km s, gas-to-dust mass ratio of 17.5, and a dust-. We found an average water abundance (HO) that resulted in a central optical depth of 160 in the 557 GHz ground state line. This model correctly reproduces the observations of the spectra at the central position in both the ground-state lines of ortho-HO, i.e. the 557 GHz and the 1670 GHz transitions.
A subregion of our HO 557 GHz spectrum map is shown in Fig. 10, where the protostar SM 1N is shown by the red circle. Apparently, the basic feature of the characteristic line profile extends over a much larger region than that. A radial cut at half-beam spacing (about Nyquist sampled) is shown below the map, from which it is clear that the BE sphere model reproduces the observed line reasonably well at the infall centre. However, as one goes farther away from the centre, observation and model diverge increasingly (see the lower part of Fig. 10). Evidently, a one solar mass collapsing protostar alone cannot account for the observed widespread inverse P Cygni profiles. It should also be clear that the observation of the spectrum towards merely one single position will not be sufficient to provide unambiguous evidence of protostellar mass infall.
4.4 Theoretical models of infall: NH
Line profiles: NH
In contrast to the HO lines, the observed line profiles of NH are free from contaminating emission from the high-velocity gas in the outflow (Fig. 11). The ammonia ground-state line can therefore be assumed to be potentially better suited to tracing protostellar infall.
In Fig. 12 b, an infall model whose physical parameters are similar to those for water has been computed for the NH line. The radiative transfer takes the overlap of the hyperfine structure (hfs) components explicitly into account. The highest optical depth of 240 is found at the frequency of the transition (see Appendix B), consistent with the observed absorption feature in the spectrum. Next to is the line that fits the observed red wing of the line. However, the theoretical emission peak is too low. This is in contrast to the blue peak that is well fitted in intensity, but there the wing is overestimated by the model. This appears to be a persistent feature of the computed profiles. Models that fit both the observed two-peak shape and the intensities have necessarily very high optical depths and the lines become broader.
We have also tried to fit a static BE configuration (Fig. 12 a). As for the infall case, the ortho-NH abundance is , i.e. a factor of 35 higher than that found on the basis of our Odin observations with a 2 beam (Liseau et al., 2003). Both the red wing and peak are well fit with this model; however, the blue peak intensity is much too low, and at the same time the blue wing is overestimated. The theoretical profile is too broad on the blue side, even though the turbulent speed is merely 0.3 km s.
As for the HO lines, little compelling evidence for gravitational infall is provided by the ammonia line observations. Essentially perfect fits can be obtained by introducing two velocity components on an ad hoc basis (Fig. 12c). Good examples would be hypothetical components at =+2.4 and +4.1 km s, and with FWHM = 0.8 km s and 0.9 km s, respectively. However, on the basis of the analysis of spectral line maps Liseau et al. (2015) found no evidence for different velocity components greater than 0.2 km s. We therefore abandon the results from Gaussian profile fitting as unreliable indicators of the velocity fields in the source.
In the densest and coldest regions of , the abundance ratio of O to HO is of the order of ten. According to theoretical models of grain chemistry, (O)/(HO) occurs only during brief periods of time, i.e. during intense O production in dense cores. This would limit the (chemical) age of SM 1 to less than 30 000 years and explain the elusiveness of O outside the solar system.
Although the line shapes of the ground state line of HO in are textbook examples of protostellar infall, detailed modelling and radiative transfer calculations of the observed spatially extended emission makes this option unlikely. A similar conclusion is reached on the basis of accompanying NH observations.
Mapping the cores A, B1, B2, C, D, E, and F in the HO line with Odin resulted in clear detections only in . Upper limits are within 10 to 50 mK (rms).
6 Epilogue: star formation in
is a region of active low-mass star formation. It is special in some respects, but also shares common properties in others. Based on a large observational material, a few conclusions that should also be valid in a wider and general context can be drawn. Below we summarise our findings and combine the results of Papers I and II.
6.1 Age of the cloud
Based on dynamical considerations, dense and cold molecular cores are generally believed to be very young. However, a proper calibration of these qualitative age assessments is not readily achievable.
The chronometry of molecular clouds is a difficult task in general, but it may be achievable under special conditions. The hope is for ‘chemical clocks’ that measure brief periods of transient times during the chemical evolution of the cloud.
In the dense core SM 1 in , the highly unusual abundance ratio of O to HO, i.e. (O)/(HO), indicates that the core is very young indeed; the age ranges from 5 000 to 30000 years. Outside this narrow window, abundance ratios equal to or larger than unity are rarely ever encountered in the ISM and rapidly approach observed upper limiting values.
However, given the association of young stellar objects with the cloud would suggest that has produced stars for at least one million years, unless the stellar sources that happen to move in front of the cloud are interlopers (at a rate of 1 pc per km s) from other regions of the . Proper motion data for the surroundings of could help to settle this issue.
For the more distant F core ( 0.8 pc, see Fig. 14), and partially also for E, Ducourant et al. (2017) determined a mean proper motion that would correspond to an average lateral speed of 8 km s. If also representative for the neighbourhood, young objects (Class II to III) could have spilled in on time scales of less than 10 yr, providing an age estimate for the cloud core of that order.
6.2 Mass of the cloud
It seems that it has generally been accepted that cloud masses are best determined from far-IR/submillimetre continuum observations, where the emission is likely optically thin and falls onto the Rayleigh-Jeans tail. Thus, to the first order, the only parameter to be determined appears to be the temperature of the dust.
In reality, however, the dust itself complicates things. In particular, the chemical composition and the size distribution of the grains determine the dust mass absorption coefficient , which, as indicated by the subscript, is frequency dependent. This potentially distorts the spectrum of the cloud from being a pure black body. In Paper I, we examine this parameter in detail to determine its impact on the mass estimates. The most frequently exploited values in the literature differ by factors of more than five.
Gas-to-dust mass ratio
To transform the estimated dust mass into total cloud mass requires the application of a gas-dust relationship, and an is commonly assumed, often already implicitly in the value of .
However, in cold cores, a constant value of one hundred will most likely not be true everywhere since much of the gas-phase molecules will locally be frozen onto the dust grains, decreasing this ratio. Thus, potential freeze-out will invalidate commonly assumed calibrations, e.g. (CO)/(H). In addition, gas and dust tracers will not measure the same parcels of material. For instance, in the cold core SM 1 in , the gas-to-dust mass ratio is down by one order of magnitude from its canonical value of one hundred, i.e. ; in addition, the projected distributions of the gas and dust do not coincide.
6.3 Star formation efficiency
Counting the stars that are associated with is, in principle, a way to determine a limit to the cloud mass already converted into stars. Within a radius of 20 (0.07 pc) around VLA 1623 A, the SIMBAD database lists 30 young stellar objects (4 TT + 26 classified as YSO). Only a few have an assigned spectral type, which means that proper mass assignments are basically impossible for the majority of objects. T Tauri stars have statistically about half a solar mass and assuming that to be the case for all objects probably provides an upper limit to the stellar mass ( ) in . Similarly, assuming that the objects listed as YSOs have a mass of 0.1 yields a strict lower limit, so the total stellar mass probably falls within the range .
The upper limit is roughly of the same order of magnitude as estimates of the present mass ( 5 to 35 , Liseau et al., 2015). The star formation efficiency, , in is likely less than 50%; however, if all 30 stellar objects have actually formed in , it could be as high as 20%. For the entire complex (L 1688), Evans et al. (2009) determined an efficiency of 6%, much less than the % of Wilking et al. (1989), but in line with the value estimated by Liseau et al. (1995, 5-7%).
6.4 Evolution of the dust
In , there is a clear trend of grain growth from the earliest to the latest phases of star formation (Fig. 13): big opacity exponents mean small grains, whereas small mean big grains. The dust grows from 0.001 m during the initial dark cloud phase to 100 m towards the later T Tauri and ZAMS phases. At intermediate times, i.e. during the protostellar and protoplanetary stages, grain sizes vary between 0.1 and 10 m. Respective grain size distribution parameters , where , are (Class and 0), (Class I and II), and 3 (Class III).
Further quantitative assessment is hampered by the poorly constrained early evolutionary time scales. However, if we assume currently adopted ages (e.g. Evans et al., 2009), the data can be fit by an exponential, as shown in Fig. 13, from which we can infer that the observed dust is dominated by a population of very small grains at the earliest times (starless cores and dynamical collapse) up to 50 - 100 thousand years. After that, grains grow predominantly in circumstellar discs, as the relative contribution to the far-IR/submm emission from these grains continues to increase.
6.5 Protostellar mass infall
Scientists with a talent for poetry once called the discovery of a collapsing protostar ‘the holy grail of star formation’, and indeed, the identification of a true collapsing protostar in appeared initially very promising (Ashby et al., 2000). Like these authors, we used the observed HO 557 GHz line to fit its profile with a protostellar infall model and radiative transfer calculations, disregarding the extended line wings which were attributed to the VLA 1623 outflow by Ashby et al. (2000). In addition, we also applied the same model to the observed NH lines. These are not contaminated with outflow emission in the line wings.
These models did indeed recover the infall signature imprinted in the HO line profile. However, our higher quality data and more advanced theoretical modelling could not uniquely confirm the proposed protostellar scenario as the observations show the infall profile over a much larger region than the model with the slightly larger than 1 potential could account for.
6.6 PDR dynamics
is sandwiched between the far-UV radiation fields from two B-type stars, one to the east and the other to the far west. The PDR parameter ranges from 0.001 to 0.01, considerably lower than for commonly studied PDRs. The spherical shell morphology of the eastern region could suggest that the dark core has been shaped and compressed by the radiation fields. However, currently we do not find convincing observational evidence for a strong pressure gradient at the PDR-border. Rather, it seems that an equilibrium situation has been achieved that has settled the gas into a quasi-static, clumpy configuration.
Acknowledgements.We thank the Swedish National Space Board (SNSB) for its continued support of our Herschel projects. The contributions to this project in its initial phase by Drs. Per Bergman, Alexis Brandeker, and Michael Olberg are acknowledged. Dr. Tim-Oliver Husser and Dr. Emanuele Bertone are thanked for providing high-resolution spectra of stellar model atmospheres.
Appendix A Odin observations of the
Prior to Herschel, we had mapped the in the HO () line at 557 GHz with the spectrometers aboard the Odin spacecraft (Frisk et al. 2003; Nordh et al. 2003; Olberg et al. 2003). Between 2003 and 2006, the was mapped in the HO 557 GHz line during several observing runs (see Table 2). The coordinates of the map centre, i.e. at offset position (0, 0), were R.A. = 1626246 and Dec. = 2354 (J2000). At the frequency of the HO () transition, the beam width of Odin is 2 (HPBW) and for , in particular, the spatial sampling was at about the Nyquist frequency, with relative offsets in 60 steps.
|Year||Date||Orbits||Used Orbits||Integration (hr)||Pointings||Source|
|2003||February 1 - 24||10577 - 10918||120||36.2||12||Oph A|
|2003||August 8 - 18||13372 - 13527||83||20.9||4||Oph A|
|2004||January 31||15992 - 16006||11||3.3||2||Oph A|
|2004||August 1 - 12||18730 - 18901||86||18.9||3 + 5||Oph B1, B2|
|2005||July 31 - October 1||24181 - 25100||394||73.5 + 109.4||51||Oph A to F|
|2006||February 2 - March 2||26956 - 27370||198||75.4||12||Oph A|
|2006||August 26 - September 18||30016 - 30357||142||48.3 + 46.6||3||VLA 16293|
The front- and backends (Frisk et al. 2003; Olberg et al. 2003) were, respectively, the 555 B1 receiver and either of the two digital autocorrelators (ACs) with a bandwidth of 400 MHz. The channel resolution was 0.5 MHz, corresponding to 0.27 km s and implying a velocity resolution of 0.5 km s. However, for the large map of 2005 and the outflow strip scan of 2006, two HO receivers were used and the backends were the other autocorrelator (AC2 in 2005) and the acousto-optical spectrometer (AOS in 2006). The observing mode was always Dicke-switching, with the 44 sky-beam pointing off by (Olberg et al. 2003).
During the mapping procedure in 2005, all 51 positions were observed once per satellite revolution in a raster scan mode. Individual spectral scans were 5 s each. An ‘off-position’, supposedly free from molecular line emission, was also observed once per revolution. The total on-source time spent on every observed point was on average 3.5 hours for the cores C to F. For the B cores, the average observing time was 6 hours on-source. Finally, for , this time was about 12 hours.
The system noise temperature was within the range 3 300 K to 3 400 K (single sideband, SSB). Examination of all the position data reveals that the relative pointing accuracy is better than 2 (rms scatter) for any individual position and better than 5 for the overall final map ( of the Odin beam width).
The basic data reduction procedure is described by Olberg et al. (2003). The appropriate equations are given in that paper and will not be repeated here. The spectrum of the off-position, obtained once per revolution, was used to monitor the reduction of the Dicke-switched spectra along the pipeline. This procedure assured that the final results are reliable. The absolute accuracy of the intensity scale is believed to be better than 10% (Hjalmarson et al. 2003). Comparison with the data obtained by SWAS shows excellent agreement (Fig. 15).
Due to the relatively low orbit of Odin, astronomical sources are generally occulted by the Earth. As a consequence, Odin ‘is observing’ through the Earth’s atmosphere twice per orbital revolution, i.e. during ingress and egress phases of the eclipse (see Fig. 2 in Nordh et al. 2003). This circumstance has an advantage, as the telluric ground state line of HO at 556.936 GHz permits a highly accurate frequency calibration (Larsson et al. 2003). In fact, the final accuracy is better than half a channel width ( km s), which takes into account the averaging of all the scans for the individual map positions, corrected for the motion of the Earth and the satellite in the local standard of rest (LSR) velocity scale.
With Odin, the average of the scans for each position results in a 70-point map for ortho-water with 60 spacing of the dense cores A to F. As can be seen in Fig. 14, no water emission was detected towards any of the cores except towards . Furthermore, Table 3 shows that the overall achieved sensitivity is about 40 mK per pointing
Appendix B Hyperfine structure in NH
The rotational levels of NH are hyperfine-split due to quadrupole interaction (see Fig. 16). The relative line strengths in equilibrium are
where the first factor in square brackets is the line strength for rotational transitions (). The relative strengths for the hyperfine transitions of the () line are given in Table 4.
Ammonia comes in two ‘flavors’, namely ortho-NH (or A-state) and para-NH (or E-state). The statistical weights for the rotational levels are , where and where refers to the spin of the three hydrogen nuclei, . These are (Townes & Schawlow 1955)
A-state or ortho-NH:
E-state or para-NH:
so that the equilibrium ratio of the weights of the ortho-to-para states is . At high temperatures, (Faure et al. 2013). For the ortho-states, takes the values , and for the para-states all the others. Therefore, the results for the ortho-line reported here are not necessarily directly related to those for the inversion lines, e.g. (1, 1), (2, 2), etc., that have been commonly cited in the literature.
The statistical weights of the hfs lines are , and the Einstein -values are obtained from
For NH, the dipole moment Debye (Poynter & Kakar 1975). The hfs levels and their weighted Einstein s have been implemented in the ALI programme, which includes 49 levels, up to K, and 104 radiative transitions. When at line centre, the relative contributions of the hfs lines will be different from those in equilibrium. Line overlap is explicitly accounted for by the transfer code.
Collision rate constants for the hfs transitions , have been published by Chen et al. (1998). However, for the sake of consistency, we computed the collision rates for the hfs transition among all J, K levels according to the prescription given by Alexander & Dagdigian (1985). These approximations are valid within a given -ladder (see e.g. Fig. 17). For cross- transitions we simply adopted the data provided by Danby et al. (1988). Included are 1162 transitions for eight temperatures in the range 15 to 300 K.
Adopted centre frequency is MHz and MHz (Cazzoli et al. 2009).
Appendix C CS line maps with the Swedish ESO Submillimetre Telescope
Mapping observations in CS (2-1), (3-2), and (5-4) were performed at the 15 m Swedish ESO Submillimetre Telescope (SEST, Booth et al. 1989) during 1996 and 1997 (Table 5). As frontends we used SIS-mixer receivers and as backend a 1000-channel acousto-optical spectrometer (AOS) with 43 kHz resolution. The observations were performed in frequency switching mode with a frequency throw of 7 MHz for the 100 and 150 GHz observations and of 15 MHz for those at 250 GHz. At these frequencies system noise temperatures were 150, 200, and 400 K, respectively.
The pointing of the telescope was regularly checked using stellar SiO masers and found to be better than 3 rms. The (2-1) and (3-2) maps were obtained simultaneously with two receivers and these covered an area of 5 5. In the (5-4) line, a was obtained. The data are internally chopper-wheel calibrated in the -scale (Ulich & Haas 1976) and as celestial calibrator, we used M 17 SW (R.A. =18h 20m 23.1s, Dec. = 11 43, J2000.0). Repeated observations on different occasions resulted in stable intensities of the lines within 4 to 5%. The values were obtained by applying the main beam efficencies, , provided in Table 5.
Standard data reduction techniques were applied in Class and with the locally available software package xs
Appendix D Synopsis of gas observations: lines and telescopes
In Table 6 we provide an overview of the spectral lines that have been observed and analysed in Papers I and II.
|Atom or||Quantum||Frequency||Upper level||Observing|
|H||, S(2)||1682||ISO-CAM CVF|
|, S(3)||2503||ISO-CAM CVF|
|, S(5)||4586||ISO-CAM CVF|
- thanks: Based on observations with Herschel which is an ESA space observatory with science instruments provided by European-led Principal Investigator consortia and with important participation from NASA.
- thanks: The data cubes of Figures 3 to 10 and 12 and A.1 are only available in electronic form at the CDS via anonymous ftp to cdsarc.u-strasbg.fr (126.96.36.199) or via http://cdsweb.u-strasbg.fr/cgi-bin/qcat?J/A+A/
- Unless explicitly noted, we are using standard values of the code.
- The upper level energy of the 145 m line is 326 K above ground, and for the 63 m transition, K.
- For the lowest rotational levels, an energy diagram can be found in van Dishoeck et al. (2011).
- Submillimeter Wave Astronomy Satellite (SWAS, Melnick et al., 2000).
- For a definition of , see e.g. Hollenbach & Tielens (1999).
- These computations are based on an ATLAS 9 model atmosphere for K, , [M/H] = 0.0 (Castelli & Kurucz, 2004).
- , where and are the dimensionless gravitational potential and the radial variable in the Lane-Emden equation, respectively.
- OB + OC; VLA Outflow: AOS + AC
- The conversion from K to Jy is about 0.02 mK/Jy (Sandqvist et al. 2003).
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