The jet and the disk of the HH 212 low-mass protostar
imaged by ALMA: SO and SO emission
Key Words.:Stars: formation – ISM: jets & outflows – ISM: molecules – ISM: individual: HH 212
Context:To investigate the disk formation and jet launching mechanism in protostars is crucial to comprehend the earliest stages of star and planet formation.
Aims:We aim to constrain the physical and dynamical properties of the molecular jet and the disk of the HH 212 protostellar system at unprecedented angular scales exploiting the capabilities of the Atacama Large Millimeter Array (ALMA).
Methods:ALMA observations of HH 212 in emission lines from sulfur-bearing molecules, SO , SO , SO , are compared with simultaneous CO , SiO data. The molecules column density and abundance are estimated using simple radiative transfer models.
Results:SO and SO show broad velocity profiles. At systemic velocity they probe the circumstellar gas and the cavity walls. Going from low to high blue-/red-shifted velocities the emission traces the wide-angle outflow and the fast ( km s) and collimated ( AU) molecular jet revealing the inner knots with timescales years. The jet transports a mass loss rate M yr, implying high ejection efficiency (). The SO and SO abundances in the jet are . SO emission is compact and shows small-scale velocity gradients indicating that it originates partly from the rotating disk previously seen in HCO and CO, and partly from the base of the jet. The disk mass is M and the SO abundance in the disk is .
Conclusions:SO and SO are effective tracers of the molecular jet in the inner few hundreds AU from the protostar. Their abundances indicate that of sulfur is in SO and SO due to shocks in the jet/outflow and/or to ambipolar diffusion at the wind base. The SO abundance in the disk is orders of magnitude larger than in evolved protoplanetary disks. This may be due to an SO enhancement in the accretion shock at the envelope-disk interface or in spiral shocks if the disk is partly gravitationally unstable.
The first steps of the formation of a low-mass star are regulated by the simultaneous effects of the mass accretion onto the star and the ejection of matter from the stellar-disk system. As a result, protostars (the so-called Class 0 objects) are characterised by the occurrence of fast bipolar jets flowing perpendicular to the plane of accretion disks. In practice, although the precise launch region (star, inner disk edge at 0.1 AU, outer disk at 0.1–10 AU) remains unknown (e.g., Ferreira et al. 2006), jets are thought to remove excess angular momentum from the star-disk system, thus allowing disk accretion onto the central object. Unfortunately, the observations of the jet-disk pristine systems in deeply embedded protostars are very difficult to perform given the small scales involved as well as the occurrence of numerous other kinematical components involved in the star formation recipe (cavities of swept-up material, infalling envelope, static ambient cloud).
The HH 212 region in Orion (at 450 pc) can be considered as an ideal laboratory to investigate the interplay of infall, outflow and rotation in the earliest evolutionary phases of the star forming process. HH 212 is low-mass Class 0 source driving a symmetric and bipolar jet extensively observed in typical molecular tracers such as H, SiO, and CO (e.g., Zinnecker et al. 1998). High-spatial resolution observations (down to ) performed with the SubMillimeter Array (SMA) (Lee et al. 2006, 2007a, 2008), the IRAM Plateau de Bure (PdB) interferometer (Codella et al. 2007; Cabrit et al. 2007, 2012), and the Atacama Large Millimeter Array (ALMA; see Lee et al. 2014; Codella et al. 2014a) reveal the inner AU collimated jet (width 100 AU) close to the protostar. HH 212 is also associated with a flattened rotating envelope in the equator perpendicular to the jet axis observed firstly with the NRAO Very Large Array (VLA) in NH emission by Wiseman et al. (2001) on 6000 AU scales. More recent SMA and ALMA observations in the CO isotopologues and HCO on , AU scales indicates that the flattened envelope is not only rotating but also infalling onto the central source, and can therefore be identified as a pseudo-disk according to magnetized core collapse models (Lee et al. 2006, 2014). In addition, a compact ( AU), optically thick dust peak is observed by Codella et al. (2007); Lee et al. (2006, 2014) and attributed to an edge-on disk rather than the inner envelope. This seems to be confirmed by HCO and CO emission showing signatures of a compact disk of radius 90 AU keplerian rotating around a source of M (Lee et al. 2014; Codella et al. 2014a). Keplerian rotating disks had previously been observed towards only other three Class 0 objects: IRAS 4A2 (Choi et al. 2010), L1527 (Tobin et al. 2012; Sakai et al. 2014), and VLA1623A (Murillo et al. 2013) but HH 212 can be considered as the only object clearly revealing both a disk and a fast collimated jet, calling for further observations aimed to characterise its inner regions.
Recent observations show that SO can be used to image high-velocity protostellar jets, similarly to a standard tracer such as SiO (Lee et al. 2010; Tafalla et al. 2010; Codella et al. 2014b). However, only two protostellar jets have been so far clearly mapped in SO lines: HH 211 and NGC1333-IRAS2A, and estimates of the SO abundance have been derived only for HH 211. On the other hand, sulfur bearing species have been long searched in protoplanetary disks associated with evolved young stellar objects (i.e. Class I-II) but only CS has been routinely observed while SO is hard to detect: Dutrey et al. (2011) reported no detection of SO (and HS) from three prototypical protoplanetary disks, Fuente et al. (2010) reported a detection in the disk of the T Tauri star AB Aur, and Guilloteau et al. (2013) showed that SO is exceptionally observed in disks with only one definite detection on a sample of 42 T Tauri and Herbig Ae stars towards 04302+2247. The statistics regarding disks around Class 0 protostars is even poorer. The only case is represented by L1527, where Sakai et al. (2014) have shown that SO as observed by ALMA originates from the outer disk near the centrifugal barrier, and its emission is argued to be enhanced by an accretion shock. Observations of HH 212 in the SO line obtained with the SMA by Lee et al. (2007a) suggest that SO may have two components, a low-velocity one originating in the inner rotating envelope/pesudo-disk and an high-velocity one from the jet. However, a detailed analysis of these two components was prevented by the lack of angular resolution and sensitivity.
In this paper we exploit the unprecedented combination of high-spatial resolution and high-sensitivity of ALMA to image and characterise both the molecular jet and the disk around the HH 212 protostar through SO and SO lines. In order to evaluate the reliability of SO and SO lines as tracers of shock chemistry and/or high-density gas in the disk, the SO and SO spatio-kinematical properties are compared with a well-know shock chemistry tracer (SiO), with a universal outflow tracer not sensitive to density or chemistry (CO), and with a disk tracer (CO). The ALMA observations of SO and SO emission presented in Sect. 2 are at higher angular resolution than previous SMA ones by Lee et al. (2007a) (HPBW , i.e. around a factor 4 better than at SMA, HPBW ) and times more sensitive ( mJy/beam/0.43 km s with ALMA and mJy/beam/km s with SMA). This allows us to disentangle for the first time the origin of the different velocity components in SO and SO and to probe the different structures in the compact circumstellar region, i.e. the cavity walls, the outflow, the molecular jet, the envelope, and the disk (see Sect. 3). The simultaneous observation of the SiO and CO/CO lines presented in Codella et al. (2014a) allows determining the molecules abundances and the physical and dynamical properties of the jet and the disk (see Sect. 4). Finally, our conclusions are summarized in Sect. 5.
2 Observations and data reduction
from the JPL molecular database (Pickett et al. 1998)
HH 212 was observed in Band 7 in the extended configuration of the ALMA Early Science Cycle 0 operations on December 1 2012 using 24 antennas of 12-m. The shortest baseline was about 20 m and the longest one 360 m, hence the maximum unfiltred scale is of 3 at 850m. The properties of the observed SO 9-8, SO , and SO transitions are summarized in Table 1 (frequency in MHz, upper level energy in K, in D, coefficient for radiative decay in s). We also report the properties of the SiO , CO , and CO transitions observed during the same run and presented by Codella et al. (2014a), which are also analysed to compare with the SO and SO emission. The obtained datacubes have a spectral resolution of 488 kHz ( km s), a typical beam FWHM of at PA, and an rms noise of mJy/beam in the 0.43 km s channel. The calibration was carried out following standard procedures and using quasars J0538â440, J0607â085, as well as Callisto and Ganymede. Spectral line imaging and data analysis were performed using the CASA111http://casa.nrao.edu and the GILDAS222http://www.iram.fr/IRAMFR/GILDAS packages. Offsets are given with respect to the MM1 protostar position as determined from the dust continuum peak by Codella et al. (2014a), i.e. (J2000) = 05 43 51.41, (J2000) = 02 5317. These values are in excellent agreement with those determined through previous SMA and PdBI observations. Velocities are given with respect to the systemic velocity, , which is estimated following the same method as in Codella et al. (2014a). They show that the emission in the CO and CS lines is most extended and most symmetric in the 0.43 km s wide velocity bin centered at km s and km s, respectively, and assume that is the average of these values. For SO 9-8 and SO the largest extension is observed at km s and km s. The average of the central velocity of CO, CS, SO, and SO gives km s, in agreement with the value adopted by Codella et al. (2014a). Note that despite the coarse velocity binning, all the lines indicate a systemic velocity significantly smaller than what determined by Wiseman et al. (2001) ( km s) and the value adopted by Lee et al. (2006, 2007a, 2014) ( km s). As argued by Codella et al. (2014a) this discrepancy can be due to the fact that the CO, CS, SO, and SO emission seen by ALMA probes the gas motion on much smaller spatial scales ( AU) than the ammonia observed by Wiseman et al. (2001) ( AU). Moreover, SO and SO line peaks towards the MM1 protostar are redshifted by km s with respect to the velocity where they show the maximum spatial extent (see Sect. 3.4). This may due to optical depth effets, and suggests that adopting the velocity where the ambient gas show the maximum extent may be more accurate than using the line peak velocity to define . Note, however, that the results presented in the paper are not dependent on the km s velocity difference between the different estimates.
In order to constrain the origin of the detected SO and SO lines we define four velocity intervals over which the emission show different morphologies and kinematic properties333the defined velocity intervals are the same as in Cabrit et al. (2007):
systemic velocity: km s;
low-velocity (LV hereafter): km s;
intermediate-velocity (IV hereafter): km s;
high-velocity (HV hereafter): km s.
Figure 1 shows channel maps of the the SO 9-8 and SO emission in the four velocity intervals defined above. In the following sections we discuss the origin and the properties of the different SO and SO velocity components.
3.1 The cavity walls
At systemic velocity SO 9-8 and SO peak at the source position (see left panel of Figure 1). This may be due to a density enhancement in the circumstellar region, as suggested by the modeling of HCO emission (Lee et al. 2014), and/or to an enhancement of the SO and SO abundances due to dust grain mantles sublimation and release of sulphur-bearing species in gas-phase.
The SO 9-8 emission also extends in a wide-angle biconical structure around the direction of the H/SiO jet which extends up to distance from source and up to km s with respect to systemic velocity. Also the SO emission shows a similar morphology even though only southern to the source. Given the high critical density of the observed SO and SO transitions ( cm for gas temperatures of K) and the similarity with the CS emission detected by Codella et al. (2014a), this emission is believed to originate in the compressed, swept-up gas in the outflow cavity walls.
3.2 The outflow
At low velocities the SO 9-8 and SO emission is bipolar and collimated along the jet direction. Blue- and red-shifted LV emission largely overlap in the two lobes similarly to SiO and (Codella et al. 2007, 2014a; Lee et al. 2007a). The emission has a transverse FWHM of , which implies an intrinsic width of AU after correction for the ALMA HPBW ( in the transverse direction), i.e. larger than the SiO jet width measured by Cabrit et al. (2007) ( AU for all velocity components). This indicates that the LV component, with large blue/red overlap, is probing a slower wider-angle outflow surrounding the narrow HV jet (see the following section).
3.3 The fast and collimated molecular jet
|(K)||(km s)||(km s)||(km s)||(km s)||(K)||(K km s)||(km s)||(K km s)|
|B2||0.3||7.3||24.1||-||33.9 0.3||370.0 1.1||82.6 0.5|
|B3||0.3||7.7||24.5||-||73.9 0.3||732.5 0.9||46.7 0.4|
|MM1||0.3||7.3||20.2||-||62.9 0.3||424.3 1.0||-||-|
|R2||0.5||16.8||26.2||-||71.1 0.5||567.5 1.6||43.3 0.7|
|B2||0.5||8.6||25.4||27.6 0.5||350.2 1.6||87.0 0.7|
|B3||0.4||11.2||27.1||25.2 0.4||294.2 1.3||30.5 0.6|
|MM1||0.3||15.9||35.3||3.4||11.2 0.3||163.2 1.1||-||-|
|R2||0.4||14.2||21.5||9.0||31.8 0.4||252.6 1.1||32.2 0.5|
|B2||0.1||6.9||23.2||4.4 0.1||49.4 0.4||18.5 0.2|
|B3||0.1||9.9||24.9||13.1 0.1||140.2 0.5||14.6 0.2|
|MM1||0.1||11.6||24.5||1.1||26.0 0.1||238.1 0.4||-||-|
|R2||0.1||12.9||23.2||7.7||11.6 0.1||118.8 0.4||2.2 0.2|
|B2||0.1||1.8||15.9||0.4 0.1||3.7 0.2||0.9 0.1|
|B3||0.1||10.9||22.4||1.2 0.1||11.6 0.3||0.8 0.1|
|R2||0.1||10.0||17.2||0.6 0.1||5.9 0.2||0.9 0.1|
|MM1||0.1||4.7||9.9||-||0.7 0.1||4.2 0.3||1.4 0.1|
|MM1||0.1||10.8||18.5||7.7 0.1||30.7 0.3||5.0 0.1|
velocities are in terms of and the error on the velocity values is half of the spectral element, i.e. km s
maximum blue- and red-shifted velocity defined as the velocity where the intensity becomes lower than rms noise
full width at zero intensity
velocity at the emission peak. No value is reported for lines affected by self-absorption, i.e. CO and SO
main-beam temperature at the emission peak. For lines affected by self-absorption could be a lower limit
line intensity integrated over the whole velocity profile (i.e. over the FWZI)
high velocity range where jet and disk emission are detected
line intensity integrated over HV. For SO and CO , (HV is measured for both the blue- and the red-shifted line wings
Figure 1 shows that as we move to intermediate and high velocities SO and SO blue- and red-shifted emission only slightly overlaps and become more and more collimated. The FWHM is which translate in an intrinsic width of , i.e. AU, after correction for the beam HPBW across the jet (). Assuming an inclination angle (Claussen et al. 1998) the de-projected gas velocity is km s.
The inferred velocity and width are in perfect agreement with those estimated from SiO emission (see, Codella et al. 2007; Cabrit et al. 2007) and indicate that for radial velocities larger than 5 km s SO and SO trace the molecular jet.
To further investigate the effectiveness of SO and SO as tracers of the molecular jet, position-velocity diagrams (PVs) of SO and SO emission along and perpendicular to the jet direction (PA, PA) are extracted and compared with the PVs of CO and SiO . Figure 2 shows that the considered tracers have different spatio-kinematical properties. While SiO traces only jet emission with no contamination from ambient gas, SO and SO are dominated by cloud and outflow emission at systemic and low velocities. Moreover, the LV and IV SO emission shows accelerating “arms”, from systemic velocity at the source position to km s at distance from source, which are not seen in SiO. These arms suggest that some SO/SO-rich gas from the circumstellar envelope or disk may be continuously accelerated to high velocities away from the source, either by interaction with the fast jet or by magneto-centrifugal forces operating on envelope/disk scales (e.g., Ciardi & Hennebelle 2010; Panoglou et al. 2012). Also CO 3–2 traces multiple components depending on velocity (the envelope, the outflow, the jet). Moreover, the CO PVs show a strong absorption feature at sligthly red-shifted velocity ( km s). Despite the different behaviours of the considered tracers, Figure 3 shows that for velocities larger than km s the PV diagram of SO along the jet is very similar to that of CO 3–2 and SiO 8–7 suggesting that at high velocities all tracers probe the same jet component. This in turn allows us to compare the emission from SO and CO at high velocities and to estimate the abundances of SO and SO in the jet by comparing their column densities with that of CO 3–2 (see Sect. 4.2).
Figure 1 shows that the SO and SO HV emission has three emission peaks, or knots, marked by black triangles on the SO and SO channel maps, two along the blue lobe (B3 and B2) and one along the red lobe (R2). The knots are located at , , and respectively, i.e. at a de-projected distance from the MM1 protostar of around 300, 700, and 650 AU. The B2 and R2 SO/SO knots are roughly coincident with the B2, R2 knots observed in SiO emission by Codella et al. (2007) and the SN, SS knots in SiO and CO emission observed by Lee et al. (2007a). 1d spectra of SO and SO are extracted at the position of the knots using the same synthetized beam and the obtained intensity profiles ( in K versus in km s) are shown in Figure 4. From the spectra we estimate the emission line properties in the knots, i.e.: rms noise in K, maximum blue- and red-shifted velocity ( and in km s), full width at zero intensity ( in km s), velocity at the emission peak ( in km s), main-beam temperature peak ( in K), and integrated intensity ( in K km s) (see Table 2).
From the velocity at the SO emission peak reported in Table 2 we estimate the knot de-projected velocity, , and derive the knots dynamical timescale, years for B3 and years for the symmetrically located blue- and red-shifted knots B2 and R2. The position, distance, velocity and dynamical timescale of the detected SO knots are summarized in Table 4. The inferred dynamical timescales are very low suggesting that the observed SO emission probes very recent ejection events. Note, however, that these estimates are affected by a large uncertainty due to the assumed inclination angle, the choice of the tracer ( varies up to a factor 3 depending on the tracer, either SO, SO, or SiO), and the broad line profiles (up to km s). The line broadening may be caused by the jet geometry and/or by internal jet shocks. In the first case, if the jet is conical and propagates with constant velocity at an angle to the plane of the sky, the observed maximum blue- and red-shifted velocities (see Table 2) imply a jet half-opening angle 444the jet half-opening angle is , where R is the algebrical ratio of maximum to minimum radial velocities and the inclination with respect to the plane of the sky (see Codella et al. 2007).. This in turn implies a knot deprojected velocity which is about a factor 2 smaller than the deprojected , implying a factor 2 larger timescales than those reported in Table 4. Note, however, that the observed line profiles are strongly asymmetric and show broad line wings, unlike the sinthetic profile computed by Kwan & Tademaru (1988) for a conical jet of constant speed seen close to edge-on (see their Fig. 1, dashed curve for and ). Alternatively, the observed line broadening may be due to an “internal jet shock” where fast jet material catches up and shocks slower previous ejecta. In this scenario, over-pressured shocked material is squeezed sideways out of the jet beam, producing an expanding bowshock structure in its wake, which cause the observed line broadening. The bowshock scenario is favored by Codella et al. (2007) and further supported by the SiO maps and transverse PVs of Lee et al. (2008). In that case, traces the bow wings while the knot velocity along the jet axis is , i.e. around a factor 2 larger than the deprojected . This would imply a factor two smaller timescales than those reported in Table 4. We conclude that the estimated knot timescales are affected by an uncertainty of a factor 2-3 depending on the origin of the line broadening and the selected knot tracer.
In Figure 4 the obtained SO and SO intensity profiles are also compared with the profiles of a typical jet tracer, i.e. SiO (Codella et al. 2007; Lee et al. 2007a), and a jet/outflow (at high/low velocities respectively) tracer, i.e. CO , extracted at the same positions and with very similar synthetized beam (differences between the beams are ). The comparison shows that even though the SiO and SO knots are roughly co-spatial SO and SO peak at higher velocities than SiO in the blue lobe and at lower velocities in the red one. This stratification in velocity is a typical signature of shock-excited emission as both the abundances of the observed molecular species and the gas physical conditions (density and temperature) undergo strong gradients in the post-shock cooling region, giving rise to chemical segregation. The different profiles and molecular stratification shown by the symmetrically located knots B2 and R2 suggest different shock conditions in the blue and red lobes within 2″ from source. This asymmetry is in contrast with the symmetry observed on larger scales in H (Zinnecker et al. 1998), and suggests that asymmetries may wash out on large timescales. A detailed comparison of the observed intensity profiles with shock models predictions is out of the scope of this paper and will be presented in Gusdorf et al. (in preparation).
Finally, it should be noted that, even though SO, SO, SiO, and CO show different behaviours at low velocity (i.e. different spatio-kinematical distribution in the PVs and different line profiles and peak velocities), the maximum blue and red velocities are remarkably similar. This indicates that, as also suggested by the comparison of the line PVs in Figure 3, all the considered lines probe the same jet component at high velocities.
3.4 The flattened rotating envelope
As suggested by Lee et al. (2007a) while the high-velocity SO 9-8 emission is clearly associated with the molecular jet, the low velocity one could arise in the inner rotating envelope and/or in a pesudo-disk. The bottom panels of Fig. 2 shows the PVs of the observed emission lines obtained cutting the datacube perpendicular to the jet PA, i.e. along the disk direction (PA). With the exception of SiO, which shows poor emission along the disk, the other lines show a velocity gradient perpendicular to the jet PA of a few km s over a scale (i.e. 450 AU), which suggests rotation in a flattened envelope and/or a disk with the same rotation sense as observed in HCO by Lee et al. (2014) and in CO by Codella et al. (2014a).
The SO and SO intensity profiles at the position of the MM1 protostar (Fig. 4), peak at sligthly redshifted velocity, i.e. km s, where CO shows a deep absorption feature. This could be a signature of infalling gas. We also note that the SO line profile shows two spectral “bumps” at blue- and red-shifted velocities, which are not detected in the other tracers. The PV diagram in Fig. 2 shows that these spectral features are seen only in a spatially unresolved region towards the MM1 position, and not along the outflow. This suggests that SO is blended with two unidentified lines emitted in the inner few 100 AU of the protostellar envelope. These lines are likely produced by Complex Organic Molecules (COMs) evaporated from the icy mantles of dust grains in the warm circumstellar region (T K), as observed towards many other low-luminosity protostars (e.g., Maury et al. 2014). The identification and analysis of these lines is out of the scope of this paper.
3.5 So 10-10: the compact rotating disk ?
Channel maps and position-velocity diagrams along and perpendicular to the jet direction are also obtained for the fainter SO line as shown in Figures 5 and 6. At difference with SO 9-8 and SO , the emission in the SO 10-10 line is compact ( AU scales) and narrow in velocity extending only up to km s with respect to systemic. As shown in Fig. 5 the red- and blue-shifted emission peaks are displaced roughly along the jet direction for velocities km s (see left panel). However, in the velocity range where CO probes a compact ( AU) disk rotating around a M star ( km s, Codella et al. 2014a) the SO 10-10 emission peaks are consistently displaced along the disk axis (see middle and right panels). This suggests that the SO emission at those velocities may also originate in the compact disk.
The PV diagram perpendicular to the jet axis in Fig. 6 supports this scenario as it shows a blue- and a red- lobe which are consistent with emission from a rotating structure. However, when cutting the datacube along the jet PA the blue- and red-shifted lobes are consistent with jet emission. This indicates that SO 10-10 traces partly rotating gas in a flattened envelope and/or in a disk, and partly the jet.
In Figure 7 the SO 10-10 line intensity profile extracted at the source position is compared with CO analysed by Codella et al. (2014a). In the velocity range the SO 10-10 and CO line profiles are very similar suggesting that at those velocities they are tracing the same disk component. However, at low velocities the SO 10-10 line intensity profile shows a red-shifted absorption feature at km s. This may be due to the SO 10-10 higher excitation energy which makes it less affected by the emission from the circumstellar gas and more sensitive to absorption by the infalling gas. As the observed SO 10-10 is compact, the detected absorption feature could probe the infall of the outer disk regions, on much smaller scales than the infalling gas traced by CO .
4.1 Searching for jet rotation
Magneto hydro-dynamical (MHD) models predict that the jet is launched by magneto-centrifugal forces which lift the gas from the disk along the open magnetic field line, accelerating and collimating it. According to these models the jet rotates around its axis carrying away the excess angular momentum from the disk, thus allowing accretion onto the central star. Many recent studies have been devoted to search for systematic velocity asymmetries across the jet axis as a signature of jet rotation, first at optical wavelengths (e.g., Bacciotti et al. 2002; Coffey et al. 2004) and more recently in the sub-mm and mm range (e.g., Codella et al. 2007; Lee et al. 2007a, b, 2008, 2009; Launhardt et al. 2009).
Tentative rotation measurements have been obtained for HH 212 by Davis et al. (2000), Codella et al. (2007), and Lee et al. (2008). Davis et al. (2000) reported a rotation speed of km s at a distance of AU from the jet axis in the southern H knot SK1 located at 2300 AU distance from the source. However, the velocity asymmetry in the northern lobe would indicate rotation in the opposite sense. Moreover, at such large distances from the driving source and from the jet axis the velocity pattern is likely dominated by effects other than rotation, i.e. by interaction with the surrounding cloud in the bow-shock wings, and/or jet wiggling and precession (e.g., Correia et al. 2009). More recent high angular resolution measurements in the inner B2 and R2 knots ( from source) gives contradictory results: Codella et al. (2007) do not observe transverse velocity shifts above 1 km s in the SiO emission with HPBW resolution, while higher angular resolution SMA observations of SiO (HPBW ) indicate a velocity shift at the tip of the two bow-shocks (Lee et al. 2008).
To test the validity of these tentative rotation measurements we search for velocity asymmetries across the collimated jet by extracting PV diagrams of SO 9-8 and SO emission perpendicular to the jet axis at the positions of knots B2, B3, and R2 (see Fig. 8). PV diagrams of SiO are not shown as these are at lower resolution with respect to those by Lee et al. (2008). The figure shows no evidence of a velocity gradient across the jet axis, consistent with the high-velocity SO/SO jet being practically unresolved transversally. Higher angular resolution () which will be available with the full ALMA array is needed to properly sample the velocity pattern across the jet width and to provide a reliable test of rotation in the HH 212 jet.
In the PV across the B3 knot there is a hint of rotation for the gas close to systemic velocity, which could indicate rotation of the outflow cavity as seen in CS (see Codella et al. 2014a).
4.2 SO and SO abundance in the jet
|knot||(km s)||( cm)||( cm)||()||(km s)||( cm)||( cm)||()|
As CO 3–2 is optically thick in knot B2, the estimated CO column density, , is a lower limit, and the SO and SO abundances, and , are upper limits.
In order to constrain the physical conditions of the gas and the SO abundance in the inner jet knots (B3, B2, and R2) the observed SO line intensities (corrected for beam dilution) and ratios (independent of beam dilution) are compared with the predictions of the statistical equilibrium, one-dimensional radiative transfer code RADEX adopting plane parallel slab geometry (van der Tak et al. 2007). Figure 9 shows the SO line temperature ratio, T(SO )/T(SO ), versus the SO intrinsic line temperature, T(SO ), predicted by RADEX for a reasonable range of kinetic temperatures ( K) and H densities ( from to cm) (e.g., Cabrit et al. 2007) and for SO column densities () increasing from 10 to 10 cm (km s). The model predictions are compared with the observed line ratios and line intensities corrected for beam dilution. As discussed in Sect. 3.3 for velocities higher than km s with respect to systemic the SO emission is co-spatial with that of CO 3â2 and SiO 8â7 indicating that SO originate in the collimated molecular jet with no contamination from the surrounding outflow and cloud. Therefore, the SO ratios and intensity are measured on the jet high-velocity range defined as: HV km s for knots B3 and B2 and HV km s for knot R2. The different velocity range defined for R2 is to avoid the strong absorption feature around km s when comparing with the CO 3–2 emission. As discussed by Lee et al. (2007a) and Cabrit et al. (2012) this feature is present at all positions and may be caused by an extended foreground component fully resolved-out by interferometers. The SO line is not detected in the jet knots, hence we report the upper limit of the considered SO line ratio. The SO line ratios do not depend on the assumed beam filling factor, as long as both lines trace the same volume of gas. The observed SO line intensity, , is corrected for beam dilution, as the jet of width is broadened by a factor by beam convolution across the jet. The errorbars on the SO line intensity in Fig. 9 accounts for the range of intensity values over HV, and for possible beam dilution along the jet axis (a factor if the knot is circular, see also Cabrit et al. 2007).
Figure 9 shows that the avalaible SO observations do not allow constraining the density and temperature of the emitting gas. However, for the considered temperatures and H densities the observations are in agreement with the model predictions for column densities cm (km s). For these values of the column density the emission in the considered SO lines is optically thin.
If the density is sufficiently high, i.e. for larger than a few 10 cm, the emission is also in local termodynamic equilibrium (LTE). In this case, an estimate of the beam-averaged SO column density in the inner jet knots can be derived from the SO line intensity integrated over the high velocities range HV assuming optically thin and thermalized emission at K. Following Cabrit et al. (2012) also the beam-averaged CO column density is derived from the CO 3–2 intensity integrated on the same velocity range as for SO assuming optically thin, LTE emission. The latter assumption is not valid for knot B2 where the ratio between the CO 3–2 and SiO 8–7 emission is over the HV velocity range. This indicates that both lines are optically thick. Therefore, for knot B2 the derived CO column density is a lower limit. Then, the abundance of SO is derived as , where the abundance of CO with respect to H is assumed to be . The main uncertainty on the estimated column densities and abundances is due to the assumption on the gas temperature, therefore we report a range of values corresponding to the temperature values K. Moreover, if the density in the jet is lower than a few cm, the SO line is sub-thermally excited implying higher SO column densities. In particular, assuming cm, the SO column densities and abundances estimated using RADEX are from a factor 2 (for K) to a factor 4 (for K) larger than in the LTE-optically thin case.
The SO emission is much fainter and covers lower velocities than SO, therefore it is not possible to compare its PV with that of CO 3–2 (see Fig. 2). However, based on the intensity profiles shown in Fig. 4 we tentatively assume that SO and CO 3–2 emission probe the same jet component over the velocity range HV km s in knots B3 and B2 and that both are optically thin and thermalized. Based on these assumptions the column density and abundance of SO are inferred by applying the same method as for SO. The main source of error on the estimated SO column densities and abundances is due to the assumed gas temperature, therefore we report a range of values for K. As for SO, if the density in the jet is lower than a few cm, the SO line is sub-thermally excited. In that case, using RADEX with cm and T K we obtain a factor 2 to 4 higher column densities and abundances.
The SO, SO, and CO line intensities integrated over the high velocities range HV in the inner jet knots (B3, B2, and R2) are reported in the last column of Tab. 2, while the column densities and abundances estimated assuming LTE-optically thin emission are summarized in Tab. 3. The SO abundance in the high velocity jet of HH 212 () is comparable to that estimated in other Class 0 jets, namely HH 211 (Lee et al. 2010), the extremely high velocity (EHV) wings of the L1448-mm and IRAS 04166+2706 outflows (Tafalla et al. 2010), and the B1 and B2 bowshocks at the tip of the L1157 outflow cavity (Bachiller & Perez Gutierrez 1997). On the other hand, the ratio SO/SO appears variable from jet to jet: we find SO/SO in HH 212, while it is around 1 in L1157, and in the L1448-mm and IRAS 04166+270 EHV gas.
Assuming a sulfur elemental abundance (S/H) (Asplund et al. 2005), these values indicate that from 1% up to 40% of the elemental S is in the form of SO and SO in the HH 212 jet. This represents an enhancement by a factor 10â100 with respect to SO and SO abundances in prestellar cores and protostellar envelopes (Bachiller & Perez Gutierrez 1997; Tafalla et al. 2010).
SO and SO enhancements up to are predicted by models of magnetized molecular C-shocks, as a result of neutral-neutral reactions in the shock-heated gas (Pineau des Forets et al. 1993) and sputtering of HS-rich icy grain mantles (Flower & Pineau des Forets 1994). The formation of SO by oxydation of SO is predicted to occur late in the post-shock, hence the ratio of SO/SO would be an indicator of shock age (see e.g. Fig 7 of Flower & Pineau des Forets 1994). In the case of HH 212, the ratio SO/SO that we observe seems inconsistent with the very short dynamical timescale of years of the knots. Such a high ratio at early times would require that SO is mainly released from grain mantles rather than formed in the gas phase out of SO as assumed in published shock models. Observations of other sulfur-bearing species, such as CS, OCS, HS, HCS, are crucial to constrain the sulfur chemistry and obtain an accurate estimate of the sulphur released in gas-phase in shocks (e.g., Podio et al. 2014). Alternatively to shocks, high SO and SO abundances up to are predicted at the turbulent interface between the outflow and the surrounding cloud material (Viti et al. 2002). SO and SO is also strongly enhanced in an MHD disk wind due to the thermal sublimation of icy mantles near the source and the gradual heating by ambipolar diffusion during MHD acceleration (Panoglou et al. 2012). Detailed comparison with such models will be the subject of a future paper (Tabone et al., in prep).
4.3 Jet physical and dynamical properties: H density and mass loss rate
|()||()||(AU)||(km s)||(yr)||(km s)||(cm)||(cm)||(M yr)|
As CO 3–2 is optically thick in knot B2, the estimated CO column density, , and hence H column and volume density, and , and mass loss rate, , are lower limits.
From the beam-averaged CO column density estimated in the HV range, (see Tab. 3), one can infer the H column and volume density in the knots along the jet, and , and the jet mass loss rate, . The determination of these quantities relies on the assumption that CO 3–2 emission at high velocity is optically thin and traces all the emitting mass in the beam. As CO 3–2 in knot B2 is optically thick the , , and values inferred for this knot are lower limits. First, the beam-averaged H column density is derived assuming . Then, the H density is derived assuming that the jet is a cylinder uniformely filled by the observed gas. Hence, , where is the jet width ( AU) and the factor accounts for the beam dilution of the AU jet emission in the transverse beam direction. Finally, the mass loss rate is estimated assuming that the mass in the jet flows at constant density and speed along the jet axis over the beam length. The latter assumption is not fulfilled if the gas in the knots is highly compressed by shocks, therefore following Lee et al. (2007a, b) we correct the mass loss rate for a compression factor of . Based on the above assumptions, is calculated as , where is the de-projected jet velocity in the considered HV range ( km s) and is the beam size across the jet width (). Note that the correction for compression is uncertain, as it depends on the unknown shock parameters (magnetic field and shock speed), therefore the values are affected by a factor uncertainty. Moreover, the estimated mass loss rate accounts only for the mass transported by the molecular jet component at high velocity ( km s). The total jet mass loss rate may be up to a factor of a few higher than the estimated value as the low velocity component can considerably contribute to the total mass loss rate if it traces a slow wide angle wind rather than entrained ambient material (see, e.g., Maurri et al. 2014). If the jet is partially ionized the atomic gas may further contribute to the mass loss rate. However, Spitzer and Herschel observations of molecular jets driven by Class 0 sources indicate that the mass flux transported by the atomic component is about 10 times smaller than that carried by the molecular gas (Dionatos et al. 2009, 2010; Nisini et al. 2015). Hence, the atomic component of HH 212, if any, should not significatly contribute to the total jet mass loss rate.
The inferred values of , , and for the B2, B3, and R2 knots are summarized in Tab. 4. These quantities are affected by the uncertainty on the estimates, hence a range of values is given depending on the temperature assumed to derive ( K). As explained above, the values inferred for the molecular high velocity jet component are a lower limit to the total jet mass loss rate. Therefore, the mass flux transported by the jet ( M yr) is at least 3% to 33% of the envelope infall rate ( M yr, Lee et al. 2006), indicating an high jet efficiency (), in agreement with the young age of the source. Note that the mass accretion rate from the disk onto the central source may be higher than the envelope infall rate. However, there are no estimates of the disk mass accretion rate for HH 212 as Class 0 sources like this one are too deeply embedded in their parental envelope to enable the use of well-calibrated accretion tracers such as, e.g., the Br line.
4.4 Disk properties: mass and SO abundance
|lobe||(kms)||( cm)||( cm)||()|