Linking low- to high-mass young stellar objects with Herschel-HIFI observations of water⋆
1 Leiden Observatory, Leiden University, PO Box 9513, 2300 RA Leiden, The Netherlands
2 Max Planck Institut für Extraterrestrische Physik, Giessenbachstrasse 2, 85478 Garching, Germany
3 Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA
4 SRON Netherlands Institute for Space Research, PO Box 800, 9700 AV Groningen, The Netherlands
5 Kapteyn Astronomical Institute, University of Groningen, PO Box 800, 9700 AV Groningen, The Netherlands
6 Université de Bordeaux, Observatoire Aquitain des Sciences de l’Univers, 2 rue de l’Observatoire, BP 89, 33270 Floirac Cedex, France
7 CNRS, LAB, UMR 5804, Laboratoire d’Astrophysique de Bordeaux, 2 rue de l’Observatoire, BP 89, 33270 Floirac Cedex, France
8 National Research Council Canada, Herzberg Institute of Astrophysics, 5071 West Saanich Road, Victoria, BC V9E 2E7, Canada
9 Department of Physics and Astronomy, University of Victoria, Victoria, BC V8P 1A1, Canada
10 Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, 53121 Bonn, Germany
Received: 21 January 2015
Accepted: 5 October 2015
Context. Water probes the dynamics in young stellar objects (YSOs) effectively, especially shocks in molecular outflows. It is therefore a key molecule for exploring whether the physical properties of low-mass protostars can be extrapolated to massive YSOs, an important step in understanding the fundamental mechanisms regulating star formation.
Aims. As part of the WISH key programme, we investigate excited water line properties as a function of source luminosity, in particular the dynamics and the excitation conditions of shocks along the outflow cavity wall.
Methods. Velocity-resolved Herschel-HIFI spectra of the H2O 202–111 (988 GHz), 211–202 (752 GHz) and 312–303 (1097 GHz) lines were analysed, together with 12CO J = 10–9 and 16–15, for 52 YSOs with bolometric luminosities ranging from <1 to >105 L⊙. The H2O and 12CO line profiles were decomposed into multiple Gaussian components which are related to the different physical structures of the protostellar system. The non-LTE radiative transfer code radex was used to constrain the excitation conditions of the shocks along the outflow cavity.
Results. The profiles of the three excited water lines are similar, indicating that they probe the same gas. Two main emission components are seen in all YSOs: a broad component associated with non-dissociative shocks in the outflow cavity wall (“cavity shocks”) and a narrow component associated with the quiescent envelope material. More than 60% of the total integrated intensity in the excited water lines comes from the broad cavity shock component, while the remaining emission comes mostly from the envelope for low-mass Class I, intermediate- and high-mass objects, and dissociative “spot shocks” for low-mass Class 0 protostars. The widths of the water lines are surprisingly similar from low- to high-mass YSOs, whereas 12CO J = 10–9 line widths increase slightly with Lbol. The excitation analysis of the cavity shock component shows stronger 752 GHz emission for high-mass YSOs, most likely due to pumping by an infrared radiation field. Finally, a strong correlation with slope unity is measured between the logarithms of the total H2O line luminosity, LH2O, and Lbol, which can be extrapolated to extragalactic sources. This linear correlation, also found for CO, implies that both species primarily trace dense gas directly related to star formation activity.
Conclusions. The water emission probed by spectrally unresolved data is largely due to shocks. Broad water and high-J CO lines originate in shocks in the outflow cavity walls for both low- and high-mass YSOs, whereas lower-J CO transitions mostly trace entrained outflow gas. The higher UV field and turbulent motions in high-mass objects compared to their low-mass counterparts may explain the slightly different kinematical properties of 12CO J = 10–9 and H2O lines from low- to high-mass YSOs.
Key words: stars: formation / stars: protostars / ISM: molecules / ISM: kinematics and dynamics / line: profiles
© ESO, 2015
The physical and chemical conditions present in low- and high-mass star-forming regions differ significantly. Massive star-forming regions are found to have higher UV radiation fields and levels of turbulence than their low-mass counterparts (see Stäuber et al. 2007; Herpin et al. 2012). The temperatures, feedback mechanisms, magnetic fields, accretion rates, and outflow forces are different between low- and high-mass young stellar objects (for more details see Bontemps et al. 1996; Behrend & Maeder 2001; Beuther et al. 2002, 2007; Palla & Stahler 1993; Cesaroni et al. 2005; Zinnecker & Yorke 2007).
However, many studies have shown that high-mass YSOs behave in certain aspects as scaled-up versions of their low-mass counterparts (van der Tak et al. 1999, 2000; Shepherd 2003; Johnston et al. 2012; San José-García et al. 2013, 2015; Karska et al. 2014a). In addition, the lifetime of the embedded phase of high-mass YSOs (0.07−0.4 Myr, Mottram et al. 2011) is comparable to that of low-mass YSOs (0.15 Myr for Class 0, 0.5 Myr for Class 0+I, Dunham et al. 2014), even if massive objects evolve more in this period. The line luminosity of molecules like CO, HCO+, and OH scales with bolometric luminosity and envelope mass, as well as the degree of turbulence in the warmer inner regions of protostellar envelopes (San José-García et al. 2013; van der Tak et al. 2013; Wampfler et al. 2013; Benz et al. 2015). Moreover, the kinematics of outflows and envelopes seem to be linked independently of the mass of the forming star (San José-García et al. 2013). Therefore, a further characterisation of the physical conditions and dynamics of these components will help to identify the differences and similarities between low- and high-mass YSOs and better understand the fundamental processes in the formation of stars.
Water is unquestionably a key molecule for studying the energetics and dynamical properties of protostellar environments (van Dishoeck et al. 2011). In particular, the analysis of the velocity-resolved water data provided by the Heterodyne Instrument for the Far-Infrared (HIFI; de Graauw et al. 2010) on board of Herschel Space Observatory (Pilbratt et al. 2010) allows us to characterise the emission from molecular outflows, which play a crucial role in the formation of young stars and in the feedback on their surroundings (Kristensen et al. 2012; van der Tak et al. 2013; Mottram et al. 2014). Given that the bulk of the water data on extragalactic sources out to the highest redshifts and the data provided by the other Herschel instruments (the Photodetector Array Camera and Spectrometer, PACS, Poglitsch et al. 2010; and the the Spectral and Photometric Imaging Receiver, SPIRE, Griffin et al. 2010) on galactic sources are spectrally unresolved, it is important to quantify the different components that make up the observed lines.
Outflows remove angular momentum effectively, which is necessary for the formation of a disk and mass accretion onto the forming star (see review by Lada & Kylafis 1999). The power agent of these structures (either jets or winds from the star/disk system) triggers not only the formation of the outflows, but also extreme and complex physical and chemical processes across the protostellar environment. In particular, different types of shocks take place in the outflow cavity wall at the interface between the cavity and the envelope. Non-dissociative outflow-cavity shocks are localised along the outflow cavity wall (Visser et al. 2012; Kristensen et al. 2012; Mottram et al. 2014). On the other hand, dissociative shocks take place either along the jet, revealed through perturbations known as extremely high velocity components (EHV; Bachiller et al. 1990; Tafalla et al. 2010), or at the base of the outflow cavity wall where the jet or wind impacts directly (Kristensen et al. 2012, 2013; Mottram et al. 2014). These shocks are also called spot shocks. Therefore, shocks and turbulent motions injected into the cavity wall propagating within this physical structure are products linked to the activity of the molecular outflow (Arce et al. 2007). The dynamical nature of these two phenomena (turbulence and shocks) is different, and they also differ from that characterising the entraining gas in classical outflows. To comprehensibly interpret the molecular emission of H2O and 12CO within the outflow-cavity system, it is important to investigate whether the dynamical properties of low-mass objects can be extrapolated to more massive YSOs.
This was one of the goals of the “Water In Star-forming regions with Herschel” key programme (WISH; van Dishoeck et al. 2011), which observed several water lines, as well as high-J CO and isotopologue transitions, for a large sample of YSOs covering early evolutionary stages over a wide range of luminosities. An extensive study has been performed on all HIFI water observations for low-mass protostars (Mottram et al. 2014) and for low-lying water transitions within the high-mass sub-sample (van der Tak et al. 2013). The line profiles of the water transitions were analysed and decomposed into multiple velocity components, which are associated to different physical structures of the protostellar system. These studies investigated trends with luminosity, mass, and evolution and explored the dynamical and excitation conditions probed by these lines. In addition, observations with PACS reveal that 12CO J> 20 transitions originate mostly in shocks for both low- and high-mass YSOs (Herczeg et al. 2012; Manoj et al. 2013; Green et al. 2013; Karska et al. 2013, 2014a,b). The excitation of warm CO is also similar across the luminosity range, but rotational temperatures in high-mass objects are higher than in their low-mass counterparts in the case of H2O, due to their higher densities (Karska et al. 2014a).
In order to link these two studies, this paper focuses on the analysis of the excited water lines across the entire WISH sample of YSOs. The ground-state water transitions of high-mass sources show absorption features from foreground clouds which complicates the extraction of velocity information from these lines (van der Tak et al. 2013), a reason why these lines are not considered in this study. Results from the line profile, line luminosity and excitation condition analysis are connected from low- to high-mass YSOs and interpreted together with those obtained from high-J12CO observations (J ≥ 10). In addition, the obtained results may help to interpret and understand those of extragalactic sources. The aim is to better constrain the dynamical properties of molecular outflows across a wide range of luminosities and complement the study presented in San José-García et al. (2013) based on the analysis of high-J CO and isotopologue transitions for the same sample of YSOs.
We start by introducing the selected sample, the studied H2O and 12CO observations and the reduction and decomposition methods in Sect. 2. The results from the water line profile and line luminosity analysis are presented in Sect. 3 and compared to those obtained for CO. In this section, the excitation conditions are also derived from the line intensity ratios. The interpretation of these results are discussed in Sect. 4, and extrapolated to extragalactic sources. Finally, in Sect. 5 we summarise the main conclusions of this work.
The sample of 51 YSOs is drawn from the WISH survey and is composed of 26 low-mass, six intermediate-mass and 19 high-mass YSOs. In addition, the intermediate-mass object OMC-2 FIR 4 (Kama et al. 2013) taken from the “Chemical HErschel Surveys of Star forming regions” key programme (CHESS; Ceccarelli et al. 2010) is added to enlarge the number of sources of this sub-group. The intrinsic properties of each source such as coordinates, source velocity (νLSR), bolometric luminosity, distance (d), and envelope mass (Menv) can be found in Mottram et al. (2014), Wampfler et al. (2013) and van der Tak et al. (2013) for the low-, intermediate- and high-mass YSOs respectively.
The sample covers a wide interval of luminosity and each sub-group of YSOs contains a mix of different evolutionary stages: both low- and intermediate-mass objects range from Class 0 to Class I; and the high-mass YSOs from mid-IR-quiet/mid-IR-bright massive young stellar objects (MYSOs) to ultra-compact H ii regions (UC H ii). The focus of this paper is to analyse different physical properties across the luminosity range; trends within the low-mass sample are discussed in Mottram et al. (2014); the intermediate- and high-mass samples are too small to search for trends within their several evolutionary stages.
The targeted para-H2O 202–111 (988 GHz) and 211–202 (752 GHz) lines and the ortho-H2O 312–303 (1097 GHz) line were observed with the HIFI instrument. The upper energy level (Eu), frequency, HIFI-band, beam efficiency (ηMB), beam size (θ) and spectral resolution of each water transition are given in Table 1. The beam efficiencies of the different HIFI-bands have been recently updated1 and in general the values decrease by 15−20% for the band considered here. The presented observations have not been corrected with the new ηMB parameters because the analysis in this paper was completed before the new numbers were available and for consistency with our previous CO study. For completeness, the new beam correction factor of each HIFI-band are listed in Table 1.
The H2O 202–111 line was observed for the entire WISH sample; the 211–202 line for 24 out of the 26 studied low-mass protostars and for all intermediate- and high-mass YSOs; and the 312–303 transition was observed for only 10 low-mass protostars, two out of six intermediate-mass sources and all high-mass YSOs. These water lines are detected for all observed intermediate- and high-mass YSOs and for 75% of the observed low-mass protostars (see Mottram et al. 2014).
Overview of the main properties of the observed water lines with HIFI.
The data were observed simultaneously by the Wide Band Spectrometer (WBS) and the High Resolution Spectrometer (HRS), in both vertical (V) and horizontal (H) polarisation (more details in Roelfsema et al. 2012). We present the WBS data because the baseline subtraction for the HRS data becomes less reliable due to the width of the water lines, which is comparable to the bandwidth of the HRS setting. Single pointing observations were performed for all targets in dual-beam-switch (DBS) mode with a chop throw of 3′. Contamination from emission by the off-position has only been identified in the H2O 202–111 spectrum of the low-mass protostar BHR71 (further information in Mottram et al. 2014). The allocated observation numbers for each source and line, designated with the initial obsIDs, are indicated in Table A.2 of Mottram et al. (2014) for the low-mass protostars, and in Table A.1 of this manuscript for the intermediate- and high-mass YSOs.
Complementing the water observations, 12CO J = 10–9 and 16−15 spectra are included in this study to extend the comparison to other components of the protostellar system traced by this molecule and set a reference for abundance studies. The J = 10–9 transition was observed as part of the WISH key programme for the entire low- and intermediate-mass sample and for the high-mass object W3-IRS5 (see San José-García et al. 2013). The J = 10–9 spectrum was obtained for AFGL 2591 from the work of Kaźmierczak-Barthel et al. (2014). For the other high-mass sources, 12CO J = 3–2 spectra are used as a proxy (San José-García et al. 2013).
12CO J = 16–15 observations of 13 low-mass Class 0 protostars were observed within the OT2_lkrist01_2 Herschel programme (Kristensen et al., in prep.). Finally, this transition was obtained for three high-mass YSOs: W3-IRS5 (OT2_swampfle_2 Herschel programme; Wampfler et al. 2014), and for AFGL 2591 (Kaźmierczak-Barthel et al. 2014) and NGC6334-I (Zernickel et al. 2012) as part of the CHESS key programme (Ceccarelli et al. 2010).
The calibration process of the water observations was performed in the Herschel Interactive Processing Environment (HIPE2; Ott 2010) using version 8.2 or higher. The intensity was first converted to the antenna temperature scale and velocity calibrated with a νLSR precision of a few m s-1. Further reduction was performed with the GILDAS-CLASS3 package. The spectra observed in the H and V polarisations were averaged together to improve the signal-to-noise ratio (S/N) and the intensity scale converted to main-beam brightness temperature scale, TMB, using the specific beam efficiencies for each band (Roelfsema et al. 2012). Finally, a constant or linear baseline was subtracted.
All data were then resampled to 0.27 km s-1 in order to compare the results among the water lines and to those from the high-J CO lines (San José-García et al. 2013) in a systematic manner. The rms noise of the spectra at that resolution, the maximum peak brightness temperature, , the integrated intensity, W = ∫TMBdν, and the full width at zero intensity, FWZI, are presented in Tables A.2 to A.4. The latter parameter was measured as in Mottram et al. (2014): first by resampling all spectra to 3 km s-1 to improve the S/N, then re-calculating the rms and finally considering the “zero intensity” where the intensity of the spectrum drops below 1σ of that rms. The velocity range constrained by the FWZI is used to define the limits over which the integrated intensity of the line is calculated.
Since the spectra have not been corrected with the recently released beam efficiencies of the different HIFI-bands1 (Sect. 2.2), the results presented in Tables A.2 to A.4, as well as those shown in Figs. 3 and 9, should be divided by the correction factor indicated in Table 1. The line profiles do not change if the new ηMB values are applied and the variation of the line ratios is of the order of 1%.
Finally, the C18O J = 10–9 emission line is detected in the line wing of the H2O 312–303 spectrum for five low-mass protostars (NGC 1333 IRAS2A and IRAS4B, Ser SMM1, GSS30 and Elias 29) and four high-mass YSOs (G5.89-0.39, W3-IRS3, NGC 6334-I and W51N-e1). Therefore, a Gaussian profile with the same FWHM, position and amplitude as those constrained in San José-García et al. (2013) was used to remove the contribution of C18O J = 10–9 line from the reduced H2O 312–303 spectrum for each of these sources. The data of these specific YSOs are then analysed and plotted after subtracting the C18O line.
As shown by Kristensen et al. (2010, 2012), van der Tak et al. (2013) and Mottram et al. (2014), the water line profiles are complex and can be decomposed into multiple velocity components. The purpose of decomposing the line profile is to disentangle the different regions probed within the protostar, which are characterised by specific physical conditions and kinematics.
Generally, these velocity components can be well reproduced by Gaussian-like profiles; other types of profiles do not give improved fits (Mottram et al. 2014). Depending on the water transition and luminosity of the source, the number of components needed to fit the profile varies. For most of the low-mass protostars, the spectra can be decomposed into a maximum of four different Gaussian components, as shown in Kristensen et al. (2010, 2012, 2013) and Mottram et al. (2014). In order to determine the number of velocity components of the water lines for the intermediate- and high-mass YSOs, these spectra were initially fit with one Gaussian profile using the idl function mpfitfun. Then, the residual from this fit was analysed and since it was larger than 3 sigma rms for all lines, an extra Gaussian component was added to the decomposition method to improve the fitting. The procedure is the same but now considering two independent Gaussian profiles. A self-absorption feature at the source velocity was detected in the H2O 202–111 line for 9 out of 19 high-mass objects, so for those objects an extra Gaussian component in absorption was added in the decomposition method. In some high-mass sources this component is weaker or non-detected in the other studied transitions and it can be negligible (for an example, see the DR21(OH) observations). Therefore, the number of components is determined by the spectrum itself and its S/N and not by the assumed method.
As in Mottram et al. (2014), we force the FWHM and central position of each component to be the same for all H2O transitions of a given sources. This procedure is adopted because, as for the water observations of the low-mass protostars, the width of the line profiles does not change significantly between the different observed transitions of a source (see also Fig. 2 of Kristensen et al. 2013, for several low-mass protostars), suggesting that in each case the emission from the excited water lines comes from the same parcels of gas. While these two parameters are constrained simultaneously by all available spectra of a given YSO, the intensities of each Gaussian component are free parameters that can vary for each transition. The resulting fits were examined visually as a sanity check. The values of the FWHM, Tpeak, νpeak, and integrated intensity of each Gaussian component are summarised in Tables A.5 to A.7.
The multiple velocity components needed to reproduce the H2O line profiles can be related to physical components in protostellar systems (Kristensen et al. 2011, 2012; van der Tak et al. 2013; Mottram et al. 2014).
Quiescent inner envelope gas produces a Gaussian profile in emission with the smallest FWHM centred at the source velocity (see Sect. 3.2.2 of Mottram et al. 2014). In previous studies these velocity components were called narrow components. The cold outer protostellar envelope can cause a self-absorption, which is more common in ground-state H2O lines and in objects with massive and cold envelopes (e.g. van der Tak et al. 2013; Mottram et al. 2013).
The chemical and physical conditions present in shocks increase the abundance of water molecules by sputtering from the grain mantles (Codella et al. 2010; Van Loo et al. 2013; Neufeld et al. 2014) and/or by the action of the high-temperature water formation route in the warm post-shock gas (van Dishoeck et al. 2013; Suutarinen et al. 2014). The line profile resulting from shocked water emission depends on the nature and kinematical properties of the shocks generating it, which translate into velocity components with different features (see Table 3 and Sect. 3.2 of Mottram et al. 2014).
The emission from non-dissociative shocks in layers along the outflow cavity wall produces velocity components with the largest FWHM (>20 km s-1) and are generally centred near the source velocity (Kristensen et al. 2010, 2013; van Kempen et al. 2010; Nisini et al. 2010; Suutarinen et al. 2014; Santangelo et al. 2014). However, these broad Gaussian-like profiles, named cavity shock components (Mottram et al. 2014) or simply broad components, should be differentiated from the broad velocity component identified in low- and mid-J (J< 10) CO spectra, even if shape and width are similar. The reason is that the water emission originates in shocks in the cavity while the CO emission comes from cooler material deeper in the wall and closer to the quiescent envelope (Raga et al. 1995; Yıldız et al. 2013).
In contrast, spot shocks occur in small localised regions and are associated to hotter and more energetic dissociative shocks (Mottram et al. 2014). This emission may originate in extremely-high velocity (EHV) gas along the jets (Bachiller et al. 1990; Tafalla et al. 2010; Kristensen et al. 2011) or at the base of the outflow cavity (previously referred to as either the medium or the offset component; Kristensen et al. 2013). These Gaussian profiles show smaller FWHMs than those measured for the cavity shock component and are in general more offset from the source velocity. A more detailed characterisation and discussion can be found in Mottram et al. (2014), van der Tak et al. (2013) and Kristensen et al. (in prep.).
The contribution of the cavity shock and envelope components with respect to the total integrated intensity of the water lines for the low-, intermediate- and high-mass YSOs are presented in Table 2 together with the analogous contribution from the entrained outflowing material (broad) and envelope gas (narrow) components for the 12CO J = 10–9 line (ratios derived from San José-García et al. 2013). The values for the low-mass Class 0 and Class I protostars were calculated by Mottram et al. (2014) for different water transitions as well as the fraction corresponding to the spot shock component (see Table 4 of that manuscript).
Left figure: averaged and normalised spectrum calculated for the low-mass Class 0 (LM0) protostars, the low-mass Class I (LMI), intermediate-mass YSOs (IM) and high-mass objects (HM) for the H2O 202–111 988 GHz (left panel), 211–202 752 GHz (middle-left panel), 312–303 1097 GHz (middle-right panel) transitions and the 12CO J = 10–9 (right panel) spectra (see San José-García et al. 2013). The spectra of each sub-group of YSOs have been shifted vertically for visualisation purposes. The low intensity features on the blue wing of the H2O 211–202 high-mass profile are due to methanol emission. Right figure: H2O 202–111, 211–202 and 312–303 spectra plotted in red, blue and purple respectively for NGC 1333 IRAS4B (LM0), GSS 30 (LMI), NGC 2071 (IM) and W33A (HM). The horizontal light green lines in both figures represent the baseline level.
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Average fraction of the integrated intensity that the narrow (envelope) and broad (cavity shock or entrained outflowing material) components contribute to the total integrated intensity of the H2O 202–111 (988 GHz), 211–202 (752 GHz), 312–303 (1097 GHz) and 12CO J = 10–9 lines.
In the following, the different velocity components of the water and CO lines are distinguished according to the terminology based on the probable physical origin of the emission and the width of the profile.
The basic properties of the spectra and their decomposition are introduced in Sect. 3.1. In Sect. 3.2 the results of the line profile decomposition are compared to those obtained for the high-J CO observations described in San José-García et al. (2013). The water line luminosity properties are also compared to those of CO in Sect. 3.3. In Sect. 3.4 integrated intensity ratios calculated for different pairs of water transitions are presented and these line ratios are further studied as across the line profiles in Sect. 3.5. Finally, the excitation conditions of the water lines are analysed with the non-LTE radiative transfer code radex in Sect. 3.6.
The observed H2O 202–111, 211–202, and 312–303 spectra for the intermediate- and high-mass YSOs are presented in Figs. A.1−A.3, respectively. The Gaussian profile fitting the broad (cavity shock) component of each water transition and source is indicated with a pink line. The spectra of the low-mass protostars are shown in Appendix A of Mottram et al. (2014).
In order to easily compare all the data, normalised averaged spectra of the H2O 202–111, 211–202 and 312–303 transitions are computed for the low-mass Class 0 and Class I protostars, the intermediate-mass sources and high-mass YSOs, as shown in the three first panels of Fig. 1. These spectra are calculated for each transition by shifting each spectrum to zero velocity, normalising it to its peak intensity and averaging it together with the observations of the corresponding sub-group of objects. The presence of self-absorption features for some of these sources, which are stronger for certain water transitions, prevents the normalised averaged spectra to reach unity for several of these lines. Independently of this issue, the normalised averaged spectra obtained for the three H2O transitions are similar for each sub-type of YSOs. Only the H2O 202–111 high-mass averaged spectrum shows a slightly different profile with respect to the other two water lines because a larger number of YSOs show strong and deep self-absorption features. Except for the Class I protostars, the averaged spectrum for a given transition seems to be broader for the low-mass objects, but at the base of the spectra the widths are similar, independent of the luminosity.
In the right-hand panel of Fig. 1 the three H2O transitions (202–111 in red, 211–202 in blue and 312–303 in purple) are over-plotted for four different sources: a low-mass Class 0 (NGC 1333 IRAS4B), a low-mass Class I (GSS 30 IRS), an intermediate-mass objects (NGC 2071) and a high-mass YSO (W33A). For each source, the shapes of the three water line profiles are similar but scaled-up in intensity, a result that is confirmed from the visual inspection of the water line profiles of all YSOs. In particular, the line wings are very similar. This indicates that the three water transitions are probing the same dynamical properties in each source.
Moving to the outcomes from the line decomposition explained in Sect. 2.5, the analysis suggests that the quiescent envelope and cavity shock components are the only two physical components consistently present in the H2O 202–111, 211–202 and 312–303 spectra of all low-, intermediate- and high-mass YSOs. The spot shock components are not detected for the excited water lines presented here towards high-mass YSOs and six out of seven intermediate-mass objects, though they have been seen in absorption against the outflow in some ground-state H2O lines for some high-mass sources (for more information see van der Tak et al. 2013).
As shown in Table 2, for a given sub-sample of YSOs the averaged contribution of the cavity shock component with respect to the total integrated intensity of the line is the same for the three water transitions. This fraction seems to decrease from low- to high-mass, but no statistically significant trend with Lbol can be claimed because the specific contribution of the cavity shock emission varies from source to source. The remaining emission comes from the envelope in the case of the low-mass Class I, intermediate- and high-mass YSOs and from spot shock components for low-mass Class 0 protostars (Mottram et al. 2014). This picture is consistent with the average spectra presented in Fig. 1, where the envelope component of the water lines is more prominent for the high-mass sources than for their low-mass Class 0 counterparts. In addition, this narrow component associated with the envelope is also less prominent in the H2O 312–303 transition regardless of the YSO mass, as expected since the envelope is presumably composed of cool quiescent gas.
Independently of these numbers, in this paper we focus on characterising the physical conditions causing the line-wing emission in the water line profiles, i.e., the broader velocity component associated to the shock emission along the outflow cavity.
Top: FWZI of the H2O 202–111 (988 GHz) emission line versus the bolometric luminosity. Middle: same as top panel but for the 12CO J = 10–9 and J = 3–2 observations. Bottom: ratio of the 12CO and H2O 202–111 FWZI values as a function of Lbol. The blue plus symbols correspond to the low-mass Class 0 protostars, the black triangles the low-mass Class I, the green asterisks the intermediate-mass objects, the pink crosses the high-mass YSOs for which the 12CO J = 3–2 spectra are used, and the red cross symbols the high-mass object for which 12CO J = 10–9 data are available (see San José-García et al. 2013). The low- and intermediate-mass sources with detected EHV components are surrounded by a box, as well as the high-mass YSO with triangular water line profiles.
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Averaged rms value in 0.27 km s-1 bin and mean (dash) values of the FWHMb and FWZI for the three water lines and the 12CO J = 10–9 spectra for each sub-type of YSO.
The fourth panel in Fig. 1 includes the normalised averaged 12CO J = 10–9 spectrum of each sub-sample of YSOs. The procedure followed to obtain these spectra is the same as that used for the water data. These averaged profiles are clearly narrower than those of water, especially for the low-mass Class 0 protostars, and the width of the spectra seems to increase from low- to high-mass. Therefore, just from a basic visual inspection of the water and the high-J CO normalised averaged spectra we can point to differences in the shape of the line profiles of these two molecules and a different trend in the width from low- to high-mass.
To consistently compare the dynamical conditions of the entrained outflowing material traced by CO and the shocked gas along the outflow cavity, the line-wing emission of these two species is analysed using two parameters: the FWHM and the FWZI (see Sects. 2.4 and 2.5). Both variables are used because FWHM characterises the average extent of emission from the source velocity while FWZI characterises the fastest material. For simplification, the FWHM of the Gaussian profile reproducing the cavity shock component is differentiated from the other velocity components by using the sub-script b to indicate that this is the broader velocity component obtained from the line decomposition.
The top panel of Fig. 2 shows the FWZI for the H2O 202–111 transition as a function of bolometric luminosity (in Fig. C.2 the FWZI of the other water lines are also plotted versus Lbol and envelope mass). Similarly, the constrained FWHMb from the cavity shock component (same for all three lines) versus Lbol and Menv are presented in Fig. C.1. The FWZI values vary from 15 to 189 km s-1, while the FWHMb range from 13 to 52 km s-1. Due to the large scatter and dispersion of the data points, no trend or correlation with luminosity can be claimed in either case. The smaller FWZI and FWHMb values are those of the low-mass Class I protostars, consistently lying at the bottom of these figures. In addition, the low- and intermediate-mass YSOs which show EHV components are marked with squares in Fig. 2 to indicate that their FWZI was calculated including the spot shock emission and to investigate if there is any particular trend for these objects. The spectra of the marked high-mass object do not have EHV components but their line profiles are characterised with broad and triangular shapes. More information about these specific sources can be found in Appendix B.
As indicated in Sect. 2.3, the 12CO J = 10–9 transition was not observed for most of the high-mass YSOs. However, those sources for which both 12CO J = 10–9 and J = 3–2 transitions were available, the values of the constrained FWHMb and also FWZI are similar within the uncertainty (see San José-García et al. 2013). Therefore, the J = 3–2 transition is used as a proxy of the J = 10–9 spectra for the study of the kinematical structure of the outflowing gas. The FWZI and FWHMb for the 12CO observations (middle panels of Figs. 2 and C.1 respectively) are spread across a smaller velocity range than that for water. In addition, the FWZI shows less scatter than the FWHMb. There is a statistically significant trend of larger FWZI for more luminous sources (4.7σ4) with a Pearson correlation coefficient r = 0.72, which is also seen to a lesser extent for the FWHMb.
Table 3 presents the mean FWHMb and FWZI values for H2O and 12CO and the averaged rms in a 0.27 km s-1 bin, σrms. In the case of the high-mass YSOs, the derived values are not affected by the higher σrms in those data (as left panels of Fig. 1 already show) since the actual signal to noise, S/N, on the water spectra themselves given by the peak intensity relative to the rms are higher (averaged S/N value of ~60) than those of their low-mass counterparts (averaged S/N of ~20).
Without considering the low-mass Class I protostars, which are more evolved and therefore have weaker, less powerful outflows (Mottram et al. 2014), the average FWHMb and FWZI values derived for H2O are similar from low- to high-mass (Table 3). A decrease of the mean FWZI values with increasing luminosity is only hinted at for the H2O 202–111 transition. Combining the results from both FWHMb and FWZI we conclude that the extent of the water line emission is similar for the entire sample. In contrast, and as suggested by the middle panel of Fig. 2, the averaged values of both FWZI and FWHMb for the 12CO observations seem to increase with luminosity.
The dispersion observed for the FWZI and FWHMb in both H2O and CO could be related to the intrinsic properties of the source, such as its inclination, evolutionary stage, clustering, etc. In order to minimise possible effects caused by these inherent characteristics, the ratio of the FWZI derived for the 12CO observations and the FWZI of the water lines is plotted versus the bolometric luminosity in the bottom panel of Fig. 2. The same procedure is followed for FWHMb of the 12CO and water spectra (see bottom panels of Fig. C.1).
Line luminosity of the H2O 202–111 (988 GHz) line emission (top), the H2O 211–202 (752 GHz) data (middle), and H2O 312–303 (1097 GHz) spectra (bottom) versus the bolometric luminosity of the source. The blue plusses correspond to the low-mass Class 0 protostars, the black triangles the low-mass Class I, the green asterisks the intermediate-mass objects and the red cross symbols the high-mass YSOs. The solid line indicates the linear correlation of the logarithm of the total line luminosity, log(LH2O), and log(Lbol). The dashed line shows the log-log correlation of the luminosity measured for the broader Gaussian velocity component only (cavity shock emission; Lbroad H2O) and log(Lbol).
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Independently of the use of FWZI or FWHMb, a correlation between these ratios and Lbol is measured for each of the three water lines with statistical significance between 3.3σ and 5.0σ (Pearson correlation coefficients between 0.50 and 0.75). While 12CO J = 3–2 and 10–9 may trace different layers in the outflow (Yıldız et al. 2013; Santangelo et al. 2013), the FWZI, i.e., the maximum offset velocity (νmax), of the CO lines does not change with the ction of the water and the high-12CO J = 3–2 and 10̵of the
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