Herschel GASPS spectral observations of T Tauri stars in Taurus: unraveling far-infrared line emission from jets and discs
M. Alonso-Martinez, P. Riviere-Marichalar, G. Meeus, I. Kamp, M. Fang, L. Podio, W. R. F. Dent, C. Eiroa
aa r X i v : . [ a s t r o - ph . S R ] A p r Astronomy&Astrophysicsmanuscript no. malonso˙herschel˙V3 © ESO 2018July 23, 2018
Herschel GASPS spectral observations of T Tauri stars in Taurus ⋆ Unraveling far-infrared line emission from jets and discs
M. Alonso-Mart´ınez1 , , , , , Dpto. de F´ısica Te´orica, Fac. de Ciencias, UAM Campus Cantoblanco, 28049 Madrid, Spain Astro-UAM, UAM, Unidad Asociada CSIC Dpto. de Astrof´ısica, Centro de Astrobiolog´ıa, ESAC Campus, P.O. Box 78, E-28691 Villanueva de la Ca˜nada, Madrid, Spain Kapteyn Astronomical Institute, University of Groningen, Postbus 800, 9700 AV Groningen, The Netherlands Department of Astronomy, University of Arizona, 933 North Cherry Avenue, Tucson, AZ 85721, USA INAF - Osservatorio Astrofisico di Arcetri, Largo E. Fermi 5, 50125, Firenze, Italy ALMA SCO, Alonso de Cordova 3107, Vitacura 763-0355, Santiago, ChileReceived: / Accepted:
ABSTRACT
Context.
At early stages of stellar evolution young stars show powerful jets and / or outflows that interact with protoplanetary discs andtheir surroundings. Despite the scarce knowledge about the interaction of jets and / or outflows with discs, spectroscopic studies basedon Herschel and ISO data suggests that gas shocked by jets and / or outflows can be traced by far-IR (FIR) emission in certain sources. Aims.
We want to provide a consistent catalogue of selected atomic ([OI] and [CII]) and molecular (CO, H O, and OH) line fluxesobserved in the FIR, separate and characterize the contribution from the jet and the disc to the observed line emission, and place theobservations in an evolutionary picture.
Methods.
The atomic and molecular FIR (60–190 µ m) line emission of protoplanetary discs around 76 T Tauri stars located in Taurusare analysed. The observations were carried out within the Herschel key programme Gas in Protoplanetary Systems (GASPS). Thespectra were obtained with the Photodetector Array Camera and Spectrometer (PACS). The sample is first divided in outflow and non-outflow sources according to literature tabulations. With the aid of archival stellar / disc and jet / outflow tracers and model predictions(PDRs and shocks), correlations are explored to constrain the physical mechanisms behind the observed line emission. Results.
Outflow sources exhibit brighter atomic and molecular emission lines and higher detection rates than non-outflow sources.The line detection fractions decrease with SED evolutionary status (from Class I to Class III). We find correlations between [OI] 63.18 µ m and [OI] 6300 Å, o–H O 78.74 µ m, CO 144.78 µ m, OH 79.12 + µ m, and the continuum flux at 24 µ m. The atomic line ratioscan be explain either by fast ( V shock >
50 km s − ) dissociative J-shocks at low densities ( n ∼ cm − ) occurring along the jet and / orPDR emission ( G > , n ∼ − cm − ). To account for the [CII] absolute fluxes, PDR emission or UV irradiation of shocksis needed. In comparison, the molecular emission is more compact and the line ratios are better explained with slow ( V shock < − ) C-type shocks with high pre-shock densities (10 –10 cm − ), with the exception of OH lines, that are better described byJ-type shocks. Disc models alone fail to reproduce the observed molecular line fluxes, but a contribution to the line fluxes from UV-illuminated discs and / or outflow cavities is expected. Far-IR lines dominate disc cooling at early stages and weaken as the star + discsystem evolves from Class I to Class III, with an increasing relative disc contribution to the line fluxes. Conclusions.
Models which take into account jets, discs, and their mutual interaction are needed to disentangle the di ff erent compo-nents and study their evolution. The much higher detection rate of emission lines in outflow sources and the compatibility of line ratioswith shock model predictions supports the idea of a dominant contribution from the jet / outflow to the line emission, in particular atearlier stages of the stellar evolution as the brightness of FIR lines depends in large part on the specific evolutionary stage. Key words.
Stars: formation, circumstellar matter, protoplanetry discs, evolution, astrochemistry, jets
1. Introduction
Protoplanetary discs are ubiquitously found around young starsand are the birth sites of planets. They are initially composed ofwell-mixed gas and dust (e.g. Williams & Cieza 2011, and ref-erences therein) and are in continuous evolution (e.g. Semenov2011). Although gas constitutes the bulk of the disc mass, be-fore the advent of ALMA our knowledge of protoplanetary discs
Send o ff print requests to : M. Alonso-Mart´ıneze-mail: [email protected] ⋆ Herschel is an ESA space observatory with science instrumentsprovided by European-led Principal Investigator consortia and with im-portant participation from NASA. was mainly based on dust studies (e.g. Beckwith et al. 1990;Andrews & Williams 2005; Hartmann 2008).Di ff erent molecular transitions probe a diversity of gas ki-netic temperatures and densities; for example, CO ro-vibrationaltransitions can be excited in hot (T ∼ n > cm − ) gas located at ∼ and CO) andforbidden atomic and ionized lines (S and Fe) trace warm gas( T ∼ + , H CO, HCN, and
CN (e.g. ¨Oberg et al. 2011, and references therein), probe coldgas (20 K < T <
50 K) at radii >
20 au. The incident radiationfield, depth in the disc, and distance from the central star, etc.,govern the chemical reactions and temperature structure of thegas in protoplanetary discs (Dutrey et al. 2014).Young stars produce X-ray and far-ultraviolet (FUV) ra-diation (Calvet et al. 2004; Ingleby et al. 2013) either bychromospheric activity (Robrade et al. 2007) or by accretion(G¨udel et al. 2007b). This radiation shapes the structure ofthe disc and its temperature distribution. In the inner 50 auof the disc surface, the temperature can be up to ∼ K(Jonkheid et al. 2004; Kamp & Dullemond 2004), favouring arich ion–atomic chemistry, while in deeper and colder ( ∼
100 K)regions where the UV / X-ray photons can still penetrate, thechemistry is more ion–molecule rich. Determination of wherethe di ff erent lines arise gives insight on accretion, photoevapo-ration, and planet formation mechanisms (Frank et al. 2014, andreferences therein).Jets, outflows, and winds associated with young stellar ob-jects have been observed from X-ray to radio wavelengths(Hartigan et al. 1995; Reipurth & Bally 2001; Bally et al. 2007;Schneider et al. 2013; Lynch et al. 2013) in scales that rangefrom tens of au (Agra-Amboage et al. 2011) up to several par-secs (McGroarty et al. 2004), and can persist for millions ofyears (Cabrit et al. 2011). The Herschel Space Observatory(HSO; Pilbratt et al. 2010) has revealed that, on average, far-IR(FIR) emission lines are more frequently seen and are strongerin systems with jets and / or outflows (e.g. Podio et al. 2012;Howard et al. 2013; Lee et al. 2014) with temperatures of ∼ / outflows and protoplanetary discs in stellar evolution istreated separately, although shocks produced by a jet are im-portant contributors to emission; hence, they a ff ect the chemicalproperties of the disc.The Gas in Protoplanetary Systems (GASPS; Mathews et al.2010; Dent et al. 2013) programme observed 240 stars in dif-ferent star forming regions in order to probe the evolutionof gas and dust in protoplanetary discs. The Herschel / PACS(Photodetector Array Camera and Spectrometer, Poglitsch et al.2010) was used to observe 76 T Tauri stars in the Taurus region.Howard et al. (2013) concentrated on [OI] 63.18 µ m, but alsolisted line intensities for [OI] 145.53 µ m and [CII] 157.74 µ mand identified other lines in the spectra. Riviere-Marichalar et al.(2012) analysed the o–H O 63.32 µ m line; Keane et al. (2014)focused on [OI] 63.18 µ m and o–H O 63.32 µ m lines in tran-sitional discs; and Podio et al. (2012) focused on the analysisof six well-known jet sources showing extended [OI] 63.18 µ memission.In this work, we make an inventory of atomic and molecularspecies covered with PACS in the Taurus sample, and present aconsistent line flux catalogue. We assess whether the observedemission is dominated by the jet or the disc, and how this de-pends on the evolutionary status of the source. Observations in-clude atomic [OI] and [CII], and molecular H O, CO, and OH.These lines have been attributed to arise in discs in TW Hya(Kamp et al. 2013), HD 163296 (Tilling et al. 2012) and HD100546 (Thi et al. 2011). Indeed disc models (e.g. the DENTgrid Woitke et al. 2010; Pinte et al. 2010; Kamp et al. 2011) canreproduce the line ratios but fail to explain high line fluxes. Inaddition, FIR line emission with a jet / outflow origin has beenspatially resolved for several Class 0 / I protostars clearly showingthat the line emission is more extended than continuum emission(Herczeg et al. 2012). The structure of the paper is as follows. Section 2 describesthe sample and observations. The data reduction is explained inSect. 3, and the main results are described in Sect. 4. Relationsbetween FIR lines are explored in Sect. 5, and its possible originsand excitation mechanisms are discussed in Sect. 6. The mainconclusions are summarized in Sect. 7.
2. Sample and observations
The sample consists of 76 T Tauri stars of the Taurus regionobserved by GASPS. Spectral types, as given by Luhman et al.(2010) and Herczeg & Hillenbrand (2014), range from K0 toM6, except for three earlier type stars: RY Tau (G0), SU Aur(G4), and HD283573 (G4). More than one-third of the starsin the sample ( ∼ / PACS results since the pixel size is 9.4 arcsec, corre-sponding to a separation of ∼ / or Rebull et al. (2010). Objects not observed by these au-thors and with no sign of infrared excess are classified asClass III. The sample is divided in outflow and non-outflowsources motivated by the correlation between the 63 / µ m con-tinuum emission and the [OI] 63.18 µ m line flux in Taurusand Chamaelon II stars found by Howard et al. (2013) andRiviere-Marichalar et al. (2014). The outflow sources are thoseshowing blue-shifted [OI] 6300 Å emission in Hartigan et al.(1995). A more detailed description of the sample is given inTable A.1, Appendix A, including stellar temperatures, mass ac-cretion rates, stellar X-ray and accretion luminosities, ages, anddisc masses. Spectroscopic observations were performed between February2010 and March 2012. PACS covers the wavelength range51–220 µ m in two channels (blue: 51–105 µ m and red:102–220 µ m). The spatial resolution of the PACS spectrometeris 9.4” at 62-100 µ m, 11.4” at 150 µ m and 13.1” at 180 µ m. Theintegral field unit (IFU) images a 47” ×
47” field of view (FOV)in 5 × × PACS Observers Manual) with a small throw (1.5’) toremove telescope and background emission. The observationswere performed in one (1152 s on source) or two (3184 s onsource) nod cycles with total integration times in the range ∼ ∼ µ m) targeting the [OI] 63.18 µ mand o-H O 63.32 µ m lines and the adjacent continuum. TheRangeSpec observations cover a larger wavelength range, de-fined by the observer. The lines observed in this mode include Table 1.
Observed wavelength ranges, resolution, instrument configuration, amount of sources observed, and properties of the maintransitions covered.
Channel λλ [ µ m] R Mode No. Sources Sp Transition T ex [K] λ [ µ m]Blue 62.93–63.43 3150 LineSpec 76 [OI] P – P
228 63.18o–H O 8 –7 + J = O 7 –6
685 71.94CH + J = J = O 4 –3
396 144.52CO J = P – P
326 145.52Blue 78.37–79.73 1900 RangeSpec 38 o–H O 4 –3
432 78.74p–H O 6 –5
396 78.92OH Π / , / – Π / , /
182 79.12 + J = P / – P /
91 157.74p–H O 3 –4
410 158.31Blue 89.29–90.72 2500 RangeSpec 30 p–H O 3 –2
297 89.99CO J = O 2 –1
115 179.53o–H O 2 –2
194 180.49 several transitions of o-H O (at 71.94, 78.74, 179.53, and 180.49 µ m), p-H O (at 78.92, 89.99, and 144.52 µ m), CO (at 72.74,79.36, 90.16, and 144.78 µ m), and a OH doublet (at 79.12 and79.18 µ m). The entire sample (76 sources) was observed inLineSpec mode, while for the RangeSpec mode the number oftargets observed varies. None of the RangeSpec observations in-cludes Class III objects. Details of the instrument channel, cov-erage, spectral resolution, observing mode, number of observedsources, and emission lines covered are summarized in Table 1.Identifiers (OBSIDs) and exposure times of the spectroscopicobservations are summarized in Table B.1.
3. Data reduction
The data were reduced using HIPEv10 (Ott 2013). The PACSpipeline removes saturated and bad pixels, subtracts the chop onand the chop o ff nod positions, applies a correction for the spec-tral response function and flat field, and re-bins at half the instru-mental resolution (oversample =
2, upsample = ± σ regions aroundeach line present in the spectral range of interest. Then, a first-order polynomial fit is applied. The line fluxes are obtained fromthe continuum subtracted spectra by Gaussian fits to the lines,and considered as real when the signal-to-noise ratio of the emis-sion peak is > σ . The errors in line fluxes are computed as theintegral of a Gaussian with width equal to the fitted value, andpeak equal to the RMS noise of the continuum. In case of non-detections, we report 3 σ upper limits computed as the integralof a Gaussian with a FWHM equal to the instrumental FWHM at the wavelength of interest, and amplitude three times the stan-dard deviation of the continuum. The line fluxes and upper limitsin the 60–80 µ m and 90–190 µ m ranges are given in Tables C.1and C.2, respectively.There are a few problems with our approach of only extract-ing the spaxel with the highest continuum level, as describedbelow. In most cases, the fluxes were extracted from the cen-tral spaxel at the location of the star. However, some Taurus ob- servations su ff er from large pointing errors, which means thatthe star lies between two or more spaxels; in these cases the re-ported fluxes are lower limits to the real flux. Previous papershave tried to solve this problem either by reconstructing the PSFto recover the on-source emission (Howard et al. 2013) or by in-tegrating all the spaxels (5 ×
5) to recover the extended emission(Podio et al. 2012). An intermediate solution that we apply is toderive the flux by summing the 3 × ff erence in flux is larger than threetimes the quadratic sum of the errors, the 3 × × µ m extended emission, we obtained lower line fluxes thanthose given in Podio et al. (2012).Figure 1 compares the [OI] 63.18 µ m line flux with thosepublished in Howard et al. (2013) and Keane et al. (2014) andthe o–H O 63.32 µ m in Riviere-Marichalar et al. (2012). Thesethree studies used di ff erent HIPE versions to reduce the data,but similar line fitting algorithms to estimate line fluxes. Themain discrepancies are towards extended, misaligned objects ortowards objects displaying very high [OI] fluxes (above ∼ − W m − ). For these observations the median di ff erences are be-tween 11% and 27%, compatible with the PACS absolute fluxaccuracy (see pages 40–44 of PACS Observers’ Manual ). Morerecent pipeline versions (HIPEv14) aim to recover the emissionfrom mispointed sources. A comparative test yields that fluxesfrom the di ff erent HIPE versions are compatible within errors.
4. Results
Atomic ([OI] and [CII]) and molecular (CO, H O, and OH)emission lines are seen in a large number of Taurus sources.Table D.1 gives the detection fractions for the entire sam-ple, as well as the split in outflow and non-outflow sources.Uncertainties were estimated assuming binomial distributions(see Burgasser et al. 2003). A clear result (Fig. 2) is that out-flow sources are richer in emission lines, and show system-atically higher fluxes (on average ∼ − W m − compared to Fig. 1.
Comparison of our [OI] 63.18 µ m fluxes with thosein Howard et al. (2013) ( top ) and Keane et al. (2014) ( mid-dle ), and comparison of o-H O 63.32 µ m fluxes with those inRiviere-Marichalar et al. (2012) ( bottom ). In all the panels thered circles represent the detections, while black down-facingand left-facing triangles are upper limits in the y-axis and x-axis, respectively. Arrows represent upper limits in both axes. ∆ F ( F old − F new F new ) is the fractional di ff erence between the two sets ofmeasurements. ∼ − W m − ) and detection fractions (on average 42% com-pared to 16%) than non-outflow sources.This suggests that jets and / or outflows are important con-tributors to the line emission and that they dominate in sourcesshowing extended emission (Podio et al. 2012). However, a (par-tial) disc origin cannot be ruled out. In the following we discussthe atomic and molecular line detections in more detail accord-ing to outflow activity, evolutionary status, and spectral types. Fig. 2.
Line emission detection fractions for the di ff erent speciesobserved within PACS range. Objects with ( top ) and without( bottom ) an outflow. Each species has a di ff erent colour: [OI]( purple ), [CII] ( yellow ), H O ( blue ), CO ( green ), and OH ( red ).The numbers on top of the bars refer to the total detectionsover the total targets observed. The atomic / molecular speciesand central wavelengths for the transitions are also indicated. Figures 3 and 4 show spectra centred on the [OI] fine-structurelines at 63.18 and 145.53 µ m. The [OI] 63.18 µ m is detected in42 out of 76 stars (55%); the line is detected towards all outflowsources, but for non-outflow sources the detection rate drops to31%. Line fluxes vary between 6 × − and 2 × − W m − ;outflows show stronger lines ( ∼ × − – 2 × − Wm − ) than non-outflows ( ∼ × − – 4 × − W m − ). UnlikeHoward et al. (2013), we did not detect [OI] 63.18 µ m in CY Tauand Haro 6–37, probably due to di ff erent reduction pipelines andcalibration files used. The profiles (see e.g. Fig. 3) are mainly Fig. 3.
Continuum subtracted spectra at 63 µ m from the spaxelshowing the highest continuum level. The scale is the same forall panels. The spectra were divided by the factor indicated inthe upper left corner of each panel. The red lines indicate theposition of the [OI] 63.18 µ m and o–H O 63.32 µ m lines. Theo ff sets seen in the [OI] 63.18 µ m line can be due to incorrectpointing or to the presence of an additional physical component. Fig. 4.
Continuum subtracted spectra centred at 145 µ m fromthe spaxel showing the highest continuum level. The scale is thesame for all the panels. The spectra were divided by the factorindicated in the upper left corner of each panel. The red linesindicate the position of the p–H O 144.52 µ m, CO 144.78 µ m,and [OI] 145.53 µ m lines.Gaussians with some skewness in a few cases (e.g. XZ Tau). Arecent and detailed study of [OI] 63.18 µ m line profiles of youngstellar objects (YSOs) by Riviere-Marichalar et al. (2016) sug-gests that such line profiles can be explained as a combination ofdisc, jet, and envelope emission.Table D.2 shows the fraction of sources where the observedspectral lines are detected as a function of SED class. In the fol-lowing, the statistics of Class II sources do not include transi-tional discs (TD). The atomic detection fractions decrease as thesources evolve from Class I down to Class III. The [OI] 63.18 µ m line is detected in 100% of Class I objects, 70% of Class IIobjects, 60 % of transitional discs, and 0% of Class III objects,with average fluxes decreasing from ∼ × − W m − , through ∼ × − , to 3 × − W m − for Class I, Class II, and transi-tional discs respectively.Table D.3 gives the line detection fractions as a function ofspectral types. The bins were selected so that they have a similarnumber of targets observed at 63 µ m. For the [OI] 63.18 µ mline, a decrease in the detection fraction with later spectral typeis observed.The [OI] 145.52 µ m line is harder to excite than [OI] 63.18 µ m because of its higher energy level (see T ex in Table 1).Consequently, the line flux is on average ten times weaker than[OI] 63.18 µ m. It is detected in 19 out of 39 objects (49%)with line fluxes between ∼ − W m − (the detection limit) and ∼ × − W m − . The detection rate is 75% (18 out of 24) foroutflow sources, while it is only 7% (1 out of 15, DE Tau) fornon-outflow objects. The fluxes decrease according to the evo-lutionary status of the source from ∼ × − W m − for Class Idown to ∼ × − W m − for transitional discs. Some spectra centred at the [CII] 157.74 µ m fine-structure lineare shown in Fig. 5. We note that the [CII] line critical densityis two orders of magnitude lower than for [OI] lines, making iteasily excited in the surrounding cloud as well. In most of ourtargets the line is also detected in the o ff - source positions. The[CII] fluxes reported in Table C.2 are for the on- minus the o ff - positions to be consistent with the procedure followed for therest of the lines.The [CII] 157.74 µ m emission line has a detection fractionof 34% (13 out of 38); only detected in 54% (13 out of 24) ofthe outflow sources with average line flux ∼ × − W m − . InDG Tau, DG Tau B, FS Tau, T Tau, UY Aur, and XZ Tau [CII] isobserved as extended with an average line flux of ∼ × − Wm − . RY Tau (TD) is associated with a jet, mapped in [OI] 6300Å (Agra-Amboage et al. 2009), and embedded in di ff use and ex-tended nebulosity (see Fig. 1 in St-Onge & Bastien 2008), sug-gesting that its [CII] line arises in its immediate surroundings.We see a clear decline in the detection fractions from Class I(50%), through Class II (36%), to transitional discs (17%), andaverage line fluxes of ∼ × − , ∼ × − , and ∼ × − Wm − , respectively. This trend is similar to that observed for [OI]63.18 µ m. Spectra of mid- to high- J CO transitions (see Table 1) are shownin Figs. 4, 6, 7, and 8. The CO J = transition is most oftendetected in the stars observed (22 out of 39). Outflow sourceshave a detection rate of 79% (19 out of 24) with an average fluxof ∼ × − W m − , while the line is only detected in 3 non-outflows sources (CI Tau, DE Tau, and HK Tau), i.e. a rate of20% (3 out of 15) and an average flux of ∼ × − W m − . Withrespect to the SED classes, the detection fractions are 100% (4out of 4) for Class I objects, decreasing to 59% (17 out of 29) forClass II objects, and 17% (1 out of 6) for transitional disc sourceswith average line fluxes of ∼ × − , ∼ × − , and ∼ × − W m − , respectively.The CO J = J = J = ∼ × − , ∼ × − , and ∼ × − Fig. 5.
Continuum subtracted spectra centred at 158 µ m fromthe spaxel showing the highest continuum level. The scale is thesame for all the panels. The spectra were divided by the factorin the upper left corner of each panel. The red lines indicate theposition of the [CII] 157.74 µ m and p–H O 158.31 µ m lines.W m − , respectively. The J = J = / II), andare known to drive powerful bipolar jets (e.g. Eisl¨o ff el & Mundt1998). None of the CO lines shows a trend with spectral type. O emission
Several transitions of water were observed (see Figs. 3, 4, 7, 8,and 9). The o–H O lines at 78.74, 179.53, and 180.49 µ m, andthe p–H O lines at 78.92, 89.99, and 144.52 µ m are only de-tected in outflow sources. The highest detection fraction is foro–H O 78.74 µ m (50%) with an average ∼ × − W m − , fol-lowed by p–H O 89.99 µ m (41%) and average flux of ∼ × − W m − . The p-H O 144.52 µ m, 78.92 µ m, o–H O 179.53,and 180.49 µ m lines have respectively a detection fraction of38%, 29%, 23%, and 14%; and average fluxes are of ∼ × − , ∼ × − , ∼ × − , and ∼ × − W m − , respectively. Theo–H O 63.32 µ m line is detected in 10 out of 27 (37%) of theoutflow sources with an average flux of ∼ × − W m − . It isthe only water line detected in non-outflow sources, seen in 3of them (6%), namely BP Tau, GI / GK Tau, and IQ Tau. This( warm ) water line was first reported in Riviere-Marichalar et al.(2012); Fedele et al. (2013). The p–H O 158.31 µ m line is un-detected in all targets (even in T Tau). Table 2 lists the averagewater line fluxes of Class I, Class II, and TD sources. Water linesare brighter and more often detected (higher detection fractions)towards Class I objects than towards Class II and TD sources. Selected spectra of the hydroxyl doublet at 79.11 and 79.18 µ mare shown in Fig. 8. The line is only detected in outflow sources(12 out of 38) with line flux values ∼ − W m − , on aver-age. It is detected in 50% of the Class I sources (average flux ∼ × − W m − ), decreasing to 38% in Class II (average flux ∼ × − W m − ), and is undetected in transitional discs. Wepoint out that DO Tau and DL Tau show peculiar OH detec-tions: DO Tau shows only the 79.11 µ m component, while DLTau only shows the 79.18 µ m component. We tested whether Table 2.
Average water line fluxes at di ff erent wavelengths (60–180 µ m) in units of 10 − W m − for Class I, Class II, and TDobjects. Source o-H O63 µ m 78 µ m 179 µ m 180 µ mClass I 1.05 2.50 6.35 2.05Class II 0.10 0.14 0.20 0.07TD 0.20 0.13 . . . . . .Source p-H O78 µ m 90 µ m 145 µ m 158 µ mClass I 0.38 1.83 0.45 . . .Class II 0.06 0.08 0.08 . . .TD . . . 0.10 0.04 . . . it could be due to significant pointing errors that translate intowavelength shifts which yield negative results (see Sect. 3.2.2in Howard et al. 2013). Thus, the detection of only one compo-nent of the OH doublet appears to be real. Such asymmetriesof OH lines in a doublet have already been noticed in ISO data(Goicoechea et al. 2011), in Class 0 / I sources (Wampfler et al.2013), and are discussed by Fedele et al. (2015) for HD100546. + emission The CH + feature remains undetected in almost all targets. Only TTau shows CH + emission at 72.14 µ m (see Fig. 6) with a line fluxvalue of 1.77 ± × − W m − . There could also be blends ofCH + with H O lines at 89.99 and 179.53 µ m. Fig. 6.
Continuum subtracted spectra at 72 µ m. The scale is thesame for all the panels. The spectra were divided by the factorin the upper left corner of each panel. The red lines indicate theposition of the CH + µ m and CO 72.84 µ m lines.
5. Observational trends of far-IR lines
For the full sample of YSOs discussed here, the fraction ofsources with [OI] 63.18 µ m detections (55 + − %) is higher thanfor older ( > ∼ ∼ ∼ η Cha (8%, age ∼ Fig. 7.
Continuum subtracted spectra at 90 µ m from the spaxelshowing the highest continuum level. The scale is the same forall the panels. The spectra were divided by the factor in the upperleft corner of each panel. The red lines indicate the position ofthe p–H O 89.99 µ m and CO 90.16 µ m lines. Fig. 8.
Continuum subtracted spectra at 79 µ m from the spaxelshowing the highest continuum level. The scale is the same forall the panels. The spectra were multiplied by the factor in theupper left corner of each panel. The red lines indicate the po-sition of the o–H O 78.74 µ m line, p–H O 78.92 µ m line, OH79.11 / µ m doublet, and CO 79.36 µ m lines.Riviere-Marichalar et al. 2015). Such a decrease with age is aclear indication of evolution.Figure 10 shows the average fluxes of the [OI] 63.18 µ m,OH 79.12 + µ m, CO 144.78 µ m, o-H O 78.74 µ m, ando-H O 63.32 µ m lines of Class I and II sources with jets, andfor Class II objects with no jets, for which the OH 79.12 + µ m and o-H O at 78.74 µ m lines are below the detection limit.The atomic and molecular line fluxes decrease rapidly, approxi-mately by an order of magnitude at each stage, suggesting that aphysical mechanism related to evolution is the most likely sce-nario.One possible explanation could be that FIR line emission isdue to a combination of jets shocking the surrounding materialand UV radiation. In general, molecular emission from shocksis expected and is very important at the earliest stages, whereasphotodissociation is more e ff ective when the envelope dissipates Fig. 9.
Continuum subtracted spectra centred at 180 µ m fromthe spaxel showing the highest continuum level. The scale is thesame for all the panels. The spectra were divided by the factorin the upper left corner of each panel. The red lines indicate theposition of the o–H O 179.53 and 180.49 µ m lines.(Nisini et al. 2002). In particular, H O in shocks is abundantbecause neutral-neutral reactions switch on at high tempera-ture (see e.g. van Dishoeck et al. 2013). Processes involving dustgrains are also important. Owing to photodesorption, sputtering,and grain-grain collisions, likely triggered by shocks, H O isalso removed from icy dust grains. The progressive dissipationof gaseous and dusty envelopes from Class 0 / I to Class III allowsstellar or interstellar FUV fields to penetrate deeper and to dis-sociate more H O and OH to produce O. This scenario proposedby Nisini et al. (2002) was followed by Karska et al. (2013) toexplain the FIR line weakening of CO and H O observed fromClass 0 to Class I objects. When the mass accretion and outflowrates drop as the source evolves, the FIR emission originatingin shock gas decrease because the strength of the FIR lines isrelated to the amount of shocked gas (Manoj et al. 2016). Thisdoes not hold for the more evolved Class II sources (‘only disc’in Fig. 10), in which FIR line emission is coming from illumi-nated discs by UV (France et al. 2014) and X-ray (G¨udel et al.2007a). It is expected that as the disc is accreted and / or dis-persed, the strength of the FIR lines will decrease too. This isalso suggested by the non detections in Class III objects. We performed an extensive search for correlations to addressthe possible origins of the FIR lines discussed here and to seehow they are related. Only those atomic and molecular lineswith high detection fractions were selected. Correlation factors ρ x , y (see Appendix A in Marseille et al. 2010), where x , y isthe pair of lines considered, are used to validate any possibletrends. The 3 σ correlation corresponds to the threshold coe ffi -cient ρ thres = / √ N −
1, where N is the number of detectionsused in the calculation. Only those trends with correlation fac-tors above the confidence threshold ( ρ thres ) are taken as statisti-cally real: | ρ | < . . < | ρ | < .
9a weak (3 σ ) correlation, and | ρ | > . Fig. 10.
Average line flux of [OI] 63.18 µ m ( red ), OH79.12 + green ), CO 144.78 µ m ( black ), o-H O 78.74 µ m( dark blue ), and o-H O 63.32 µ m ( blue ) lines of Class I andII objects with jets and Class II sources with no jets. Most ofthe TD sources are included within Class II (only disc) objects.The exception is RY Tau (see Sect. 4.1), included in the Class II(disc + jet) group.Figure 11 shows the most promising correlations. Bothatomic and molecular lines are observed to correlate. Two newtight ( | ρ | > .
9) correlations between FIR lines are identified:one between [OI] 63.18 µ m and CO 144.78 µ m, and the otherbetween [OI] 63.18 µ m and o-H O 78.74 µ m. Perhaps statisti-cally less meaningful because of T Tau, we also find correlationsbetween CO 144.78 µ m and o-H O 63.32 µ m, and between OH79.12 + µ m and o-H O 78.74 µ m. The correlation between[OI] 63.18 µ m and o-H O 63.18 µ m has already been observedin T Tauri stars (Riviere-Marichalar et al. 2012).
6. Discussion
The [OI] 63.18 µ m line can arise in the surface of discs de-pending on disc size and spectral type (Gorti & Hollenbach2008). It can be produced in photodissociation regions (PDRs;Tielens & Hollenbach 1985), in shocks (Neufeld & Hollenbach1994), or in the envelopes of Class I sources (Ceccarelli et al.1996). Nothing precludes all mechanisms from contributing si-multaneously. Given that the majority of our objects are ClassII (see Sect. 2), whose envelopes are likely already dissipated,we dedicate the following sections to a discussion of the mostprobable ones, i.e. shocks and discs. Each scenario is consideredseparately; line fluxes and their line ratios (see Appendix C) arecompared with shock and disc model predictions. We stress thatin both scenarios a contribution from PDRs is also expected. Jets, outflows, and winds associated with PMS (pre-main se-quence) stars can be traced with forbidden lines, e.g. [OI]6300 Å (e.g. Appenzeller & Mundt 1989; Edwards et al. 1993;Hirth et al. 1997). A correlation between these jet / outflow trac-ers and FIR lines would suggest a similar origin. Figure 12shows the [OI] 63.18 µ m luminosity as a function of the [OI]6300 Å line luminosity from Hartigan et al. (1995) integratedover the entire profile. The lines correlate ( ρ ∼ − with respect to the stellarvelocity, tracing collimated jets, and a low velocity component(LVC) shifted by 5–20 km s − , whose origin is possibly due to aphotoevaporative wind (Rigliaco et al. 2013; Simon et al. 2016).Unfortunately, the PACS spectral resolution at 63 µ m is ∼ − , not high enough to resolve velocity components, but wenote that in several cases the wings are broad. Such broadeningis observed in the red wing of the [OI] 63.18 µ m line profile (seeFig. F.1) of CW Tau, DO Tau, DQ Tau, FS Tau, Haro6-5 B, HVTau, and RW Aur. Interestingly, the [OI] 6300 Å and [OI] 63.18 µ m line profiles of RW Aur do not show the same line shape,but have roughly the same 300 km s − broadening. This fur-ther suggests the presence of several components like in HH 46(van Kempen et al. 2010) and DK Cha (Riviere-Marichalar et al.2014). The line ratios of [OI]63 / [OI]145 combined with[CII]158 / [OI]63 can be used as diagnostics of the excita-tion mechanisms (e.g. Nisini et al. 1996; Kaufman et al. 1999).The [OI]63 /
145 ratios of our sample are between ∼
10 and70 with a median of ∼
23, therefore compatible with ISOobservations (Liseau et al. 2006) and the ratios observed inHerbig Ae / Be stars (Meeus et al. 2012; Fedele et al. 2013).There is no statistical di ff erence in terms of [CII] / [OI] line ratiosbetween extended objects (detected in more than one spaxelbut not incorrectly pointed) and compact objects (only detectedin one spaxel) in our sample. We note that in those sourceswith extended outflow emission, the ISO and Herschel absolutefluxes are not expected to be similar, due to the much largerbeam and a lack of background emission subtraction in ISO.Figure 13 shows the observed atomic line ratios compared toshock model predictions by Flower & Pineau des Forˆets (2015).Both C- and J-type shocks can reproduce a [OI]63 / [OI]145line ratio for a wide range of shock parameters ( V shock and n ). However, such models fail to explain the observed[CII]158 / [OI]63 ratios. This is not surprising as [CII] is thoughtto arise in PDRs. To further disentangle the origin of [OI]and [CII] lines, in Fig. 14 we combined [OI]63 / [OI]145 and[CII]158 / [OI]63 atomic line ratios and compared them to thePDR models of Kaufman et al. (1999) and the higher velocity( V shock >
50 km / s) J-shock models by Hollenbach & McKee(1989). Figure 14 indicates that the observed ratios are all com-patible with PDR models with densities between ∼ and 10 cm − and FUV fields G > ; only a few cases are compatiblewith fast J-shocks with low pre-shock densities ( V shock = − , n ∼ cm − ), or both. The sources whose line ratiosare compatible with shocks lie in a region in which the modelsoverlap. Thus, it is impossible to discern which phenomenon isresponsible for the emission. Similar PDR and shock parame-ters were obtained by Podio et al. (2012). Although weak [CII]in shocks is predicted by models (Flower & Pineau Des Forˆets2010), the [CII] 157.74 µ m line is more likely to originatein PDRs, as studies based on SOFIA / GREAT (Heyminck et al.
Fig. 11.
Correlation plots between: a) [OI] 63.18 µ m and [OI] 145.52 µ m; b) [OI] 63.18 µ m and [CII] 157.74 µ m (see Howard et al.2013); c) [OI] 63.18 µ m and CO 144.78 µ m; d) CO 144.78 µ m and o-H O 63.32 µ m; e) CO 144.78 µ m and o-H O 78.74 µ m; f) OH 79.12 + µ m and o-H O 78.74 µ m; g) [OI] 63.18 µ m and o-H O 63.32 µ m (see Riviere-Marichalar et al. 2012); h) [OI]63.18 µ m and o-H O 78.74 µ m; and i) [OI] 63.18 µ m and o–H O 179.53 µ m (tentatively). The red circles represent detections, greytriangles are upper limits in their pointing direction, and black arrows are upper limits in both axes. Filled symbols represent outflowsources. The solid line corresponds to a linear fit for detections.2012) observations suggest (e.g. Sandell et al. 2015; Okada et al.2015).In the few cases where [CII] 157.74 µ m has been spectro-scopically resolved with Herschel / HIFI, it is clear that the line isnot due to a disc, but rather to a remnant envelope or a di ff usecloud (HD 100546 Fedele et al. 2013), or even to to PDR emis-sion in the outflow (DG Tau, Podio et al. 2013). PACS observa-tions of Upper Scorpius, have revealed low [CII] fluxes in two TTauri stars (Mathews et al. 2013). These are early K-type proto-planetary systems without any signature of jet / outflow emission,further suggesting that [CII] 157.74 µ m emission is PDR domi-nated. Figure 15 shows a combination of molecular line ra-tios compared to C-type and J-type shock models fromFlower & Pineau des Forˆets (2015). The excitation conditionsare, within errors, compatible with both C- and J-type shocks,with pre-shock densities between 10 cm − and 10 cm − and V shock = − for most of the sources. The agreementbetween observations and shock model predictions depends on the specific line ratio (see Karska et al. 2014b) and the evolu-tionary stage of each source. Similar conditions are found inthe few cases where individual sources have been compared toshock models (Lee et al. 2013; Dionatos et al. 2013). However,we must stress that C-type shocks are probably the main driverof the molecular emission (see below).Non-outflow sources show low J CO and ‘hot’ ( T ex > O line detections. In addition, outflow sources also showhigh J CO and ‘cold’ ( T ex <
700 K) o-H O lines. This indicatesthat high J CO transitions ( J up ≥
20 in Karska et al. 2014a) areharder to excite in discs (Woitke et al. 2009), whereas shocks canaccount for such emission (van Kempen et al. 2010; Visser et al.2012); in addition, the fact that combinations of hot and coldo-H O lines are compatible with shock models (upper panels ofFig. 15) with similar parameters suggests that H O (and CO) canarise in similar regions.The OH / H O line ratios (see lower right panel of Fig. 15)can only reproduce the excitation conditions of HL Tau, whichshows emission from very hot water lines (Kristensen et al.2016). Karska et al. (2014b) could not reproduce those line ra-tios for their sample of Class 0 / I source in Perseus. The dis-crepancy between observations and models depends on the H O Fig. 12. [OI] 63.18 µ m as a function of [OI] 6300 Å line lumi-nosities in solar units. The fit is only for detections and is indi-cated by the solid line. The red circles represent detections, greytriangles are upper limits in their pointing direction, and blackarrows are upper limits in both axes. Filled symbols representoutflow sources. Fig. 13.
Observed atomic line ratios compared to C-type (reddashed lines) and J-type (blue lines) shock model predictionsfrom Flower & Pineau des Forˆets (2015) for shock velocities( V shock ) between 10 and 40 km / s. The circles, diamonds, trian-gles, and squares correspond to pre-shock densities ( n ) of 10 cm − , 10 cm − , 10 cm − , and 10 cm − , respectively. The rangeof the observed atomic ratios is represented by the shaded region.transition that are used because dissociative radiation (Lyman α photons) may have an impact on the composition of the pre-shocked gas (Flower & Pineau des Forˆets 2015), hence a ff ectingthe abundance of H O (Melnick & Kaufman 2015). Spitzer mid-IR observations of OH in DG Tau by Carr & Najita (2014) sup-port the idea of hot OH emission induced by dissociation of H Oby FUV radiation.
Herschel observations of OH in the range 70–160 µ m indicate that the OH emission originates from dissocia-tive shocks in young stellar objects (Class 0 / I in Wampfler et al.2010). Further modelling of the OH radical by Wampfler et al.(2013) with radiative transfer codes of spherical symmetric en-velopes could not reproduce the OH line fluxes nor the linewidths, strongly suggesting that the OH is coming from shockedgas. There is evidence pointing to C-type rather than to J-typeshocks as the main mechanisms driving the excitation of molec-ular FIR lines (see Karska et al. 2014b, for a discussion). The[OI] / H O line ratios can be used to discern between shocktypes. We followed the criteria established by Lee et al. (2014).The division between J- and irradiated C-type shocks occurswhen [OI] / H O ∼ / H O ∼
1. Table 3 lists the [OI] / H O of our sample.The ratios of the outflow sources are compatible with irradiatedC- shocks; and only in XZ Tau ([OI] / H O > / GK Tau, and IQ Tau are low(1 < [OI] / H O < Table 3. [OI] / o–H O line ratios.
Target 63 /
63 63 /
78 63 /
179 63 / / GK Tau 1.84 – – –Haro6-5B 2.87 3.66 – –Haro6-13 – 10.42 – –HL Tau 8.86 5.79 – –HN Tau 9.19 – – –IQ Tau 1.60 – – –IRAS04385 + Notes.
All the targets listed in Col. (1) have [OI] 63.18 µ m detected.Columns (2) to (5) show the line ratios between the [OI] line at 63.18 µ m and o–H O lines at 63.32, 78.74, 179.53, and 180.49 µ m, respec-tively; otherwise is indicated by ‘–’. Table 4.
Diameter (in au) of the emitting area for H O, CO, andOH estimated from C- and J- type shock models. Source D (H O) D (CO) D (OH)AA Tau – – 102–109DP Tau – 56–122 –FS Tau A 33–99 184–395 261–277HL Tau – 339–821 174–183T Tau 160–364 307–610 1613–1701UY Aur – 53–115 212–225XZ Tau 33–97 225–730 170–181 To further test the shock scenario we follow theFlower & Pineau des Forˆets (2015) models to estimate the size( D ) of the emitting areas (see Table 4) necessary to reproducethe observed molecular line fluxes. If the emission is indeed as-sociated with shocked gas, the emitting areas have to be com-patible with the observed scales (10 arcsec level) of molecular Fig. 14.
Observed [OI]63 / [OI]145 and [OI]63 / [CII]158 line flux ratios plotted as a function of di ff erent emitting conditions. ThePDR models (grey) are from Kaufman et al. (1999); and J-shock models (black) from Hollenbach & McKee (1989). For the PDRmodels, the labels indicate the gas density ( n ) and the intensity of the FUV field ( G ), respectively. For the shock models, the labelsdenote pre-shock densities ( n ) and shock speeds ( V shock ). The data is plotted according to their SED class: blue circles are Class I,red triangles are Class II, and green stars are TD.gas in T Tauri stars. The OH areas are computed assuming thatthe same physical conditions ( V shock and n ) as for HL Tau holdsfor all detected objects. Details of the derivation are in AppendixE. We find that the emitting areas range from tens to hundredsof au, consistent with molecular emission being compact andunresolved with PACS at the distance of Taurus. The CO emit-ting areas are larger than those of H O, found to be betweentens of au and a few hundred. In particular, the size of the H Oemitting area for T Tau is comparable with previous estimates(Spinoglio et al. 2000; Podio et al. 2012), and compatible withthose obtained by Mottram et al. (2015) for Class 0 / I sources.Concerning the spatial extent of H O compared to [OI] alongthe outflow, in the maps presented in Nisini et al. (2015) the twospecies (o-H O 179.53 µ m and [OI] 63.18 µ m) show a simi-lar extent. However, this depends on the H O line consideredbecause various transitions require di ff erent physical conditionsfor excitation. Even in the case of jet / outflow emission, the waterlines can originate in much denser – and probably more confined– regions compared to [OI]. We now compare the observed fluxes with dust tracers, i.e. in-frared continuum, to try to determine the contribution of the discto the line emission.Figure 16 represents the flux of the [OI] 63.18 µ m line asa function of Spitzer / IRAC 3.6, 8.0 µ m, Spitzer / MIPS 24 µ m,and PACS 70 µ m (Rebull et al. 2010; Luhman et al. 2010;Howard et al. 2013). The outflow sources are clearly brighterthan non-outflow ones. While there is no obvious trend with 3.6 µ m and the scatter is large at 8.0 µ m, there is a clear trend at 24and 70 µ m. Longer continuum wavelengths are associated withcolder dust, probing deeper in and further out disc regions. Atlonger wavelengths it is more likely that continuum and FIR lineemission both come from the same radial zone.The distribution of Taurus sources in the [OI]–infrareddiagrams is connected with the evolutionary status of thesources and the presence of outflows. The di ff erent behaviour ofoutflow / non-outflow sources in such diagrams has already beenpointed out by Howard et al. (2013). The correlation between theline flux and the dust emission for non-outflow sources suggeststhat both arise in the disc. The contribution of the disc to the gas Fig. 15.
Observed molecular line ratios compared with C-type (red dashed lines) and J-type (blue lines) model predictions fromFlower & Pineau des Forˆets (2015). The circles, diamonds, triangles, and squares correspond to pre-shock densities of 10 cm − ,10 cm − , 10 cm − , and 10 cm − , respectively, with shock velocities ( V shock ) from 10 km s − to 40 km s − (C-type) or from 10km s − to 30 km s − (J-type). Filled symbols indicate the position of the lowest velocity ( V shock =
10 km s − ). Black filled circlesrepresent detections and arrows are upper limits. Solid error bars are the intrinsic errors of the line fluxes and dotted error bars areobtained by adding a 30% error due to the PACS flux calibration.emission in the outflow sources can be estimated assuming thatthe [OI]–70 micron correlation for non-outflow sources holdsfor all discs. In the case that outflow and non-outflow sourcesshow similar trends with continuum emission, this test cannotdistinguish clearly between the two origins or whether outflowemission does not have a relevant e ff ect. A linear fit to our datafor such correlation is given bylog ( F [OI]63 µ m ) = − . + . × log ( F µ m ) , (1)where F [OI]63 µ m is the [OI] 63.18 µ m line flux in W m − and F µ m is the continuum flux at 70 µ m in Jy. Table 5 shows thedisc contribution in terms of the SED Classes. The relative con-tribution from the disc increases as the system evolves. In ClassI sources the disc contributes ∼
20% and it keeps increasing untiloutflow activity dissipates. This is clear when comparing ClassII sources with and without outflows (38% compared to 100%).This is in agreement with Podio et al. (2012) who obtained adisc contribution between 3% and 15% for Class I and ClassII sources with outflows and showing [OI] 63.18 µ m extendedemission. In the case of T Tau the disc contribution ( < Table 5.
Ranges and mean (median) values of disc contributionin percentage to the [OI] 63.18 µ m line flux for di ff erent SEDclasses. SED Class Range Mean (Median) N sourcesClass I 1%–51% 19% (18%) 4ClassII + Jet 3%–100% 38% (24%) 11TD + Jet 43% 43% (43%) 1ClassII 65%–100% 92% (99%) 7TD 100% 100% (100%) 2 O We showed (see Sect 5.2) that the o–H O 63.32, 78.74 and179.53 µ m line luminosities correlate with [OI] 63.18 µ m ( ρ > ff erent line luminosities.Their models do not match our observed water line luminosi-ties and fail to reproduce the slopes we observe. The predic- Fig. 16. [OI] 63.18 µ m flux versus: IRAC 3.6 µ m ( top left ); IRAC 8.0 µ m ( top right ); MIPS 24 µ m ( bottom left ); and PACS 70 µ m ( bottom right ) continuum fluxes. Class I, Class II, Class III, and transition discs (TD) are represented by circles ( blue ), squares( red ), asterisks ( black ), and stars ( green ), respectively. Filled symbols corresponds to outflow sources. Arrows indicate upper limits.The dashed lines are 1:1 relations and the solid line is a fit to disc sources as explained in the text.tions from the fully parametrized DENT disc grid do reproducethe observed H O / [OI] and CO / [OI] ratios, but fail to explainthe high H O and CO fluxes (Podio et al. 2012), suggesting thatan additional component and / or di ff erent gas-to-dust ratios areneeded to account for these high fluxes. In order to account forthese discrepancies, higher disc masses and / or low dust-to-gasratios, high FUV fluxes, or discs heated by X-rays (Aresu et al.2011, 2012) are needed. Indeed, Podio et al. (2013) showed thata model of DG Tau with a massive gaseous disc associated withstrong UV and X-ray radiation reproduces the H O lines well.Other options include a more complicated inner disc structure,such as a pu ff ed up inner rim (Aresu et al. 2011, 2012), gap, orhole. From our analysis described above, a jet / outflow origin isfavoured for the strong FIR lines because of (1) the observed ex-tended emission in outflow sources (see maps in e.g. Nisini et al.2010; Podio et al. 2012; Nisini et al. 2015), (2) the detectionspredominantly in outflow sources and the correlations between Table 6.
Comparison of the slopes (with errors) between thestrength of the water lines and OI observed, and the predictionsby Aresu et al. (2012) for two di ff erent models. λ ( µ m) Transition Observed UV UV + X-raymodel model63.32 8 –7 –3 –1 emission lines, and (3) the compatibility of line ratios with shockmodels especially for Class I sources for which the disc contri-bution is estimated to be smaller than 20%. Nonetheless, we notethe following: (1) several detections of [OI] and o–H O at 63 µ min non-outflow sources; (2) correlations between the line fluxesand continuum for 24 and 70 µ m , which point to a disc origin;(3) compact molecular emission within the PACS beam also ob-served in Class I sources, which have (small) outflows; and (4)gas line ratios reproduced satisfactorily by the disc models (e.g.Kamp et al. 2011; Aresu et al. 2012). Therefore, the excitationmechanism in discs should not be the main problem. Di ff erent dust-to-gas ratios, larger scale heights, or inner gaps would makethe disc likely hotter than a normal continuous disc model, butother agents may only increase the emitting surface area. Thisparticular issue cannot be fully understood without spectroscop-ically and spatially resolved observations, so that the locationand dynamics of the emitting gas can be pin-pointed; however,for FIR observations this is hard to obtain. A promising alterna-tive to this problem could be to use lines which are likely co-spatial with some of the FIR lines, such as ro–vibrational linesof CO (see e.g. Banzatti & Pontoppidan 2015). This way, onecould determine whether the lines are purely associated witha Keplerian disc or coming from shocks via jets and / or winds.Within a disc, gas kinematics derived from resolved line pro-files may help disentangle various regions, especially at shorterwavelengths where the spatial resolution is higher.
7. Summary and conclusions
We provide a catalogue of line fluxes in the FIR (63–190 µ m) forT Tauri stars located in Taurus, surveyed with Herschel / PACS aspart of the GASPS project. The species observed include [OI],[CII], CO, H O , and OH. The origin of the atomic and molecularemission is investigated by comparing line fluxes and line ratiosto shock and disc models. The main conclusions are as follows: – Outflow sources exhibit brighter atomic and molecularemission lines and higher detection rates than non-outflowsources. In agreement with previous studies, atomic andmolecular FIR line emission in T Tauri systems is observedto decrease with evolutionary stage. – The [OI] line emission is brighter and more often detectedin systems with signs of jet / outflows and in early phases ofstellar evolution (Classes I and II), suggesting a dominantcontribution from shocks in young outflow systems. In thesesystems the emission is spatially extended (10 arcsec level).Slow ( V shock = − ) J- and C-type shocks with den-sities between n = and 10 cm − ) can reproduce the ob-served [OI]63 / [OI]145 ratios. These models do not repro-duce the [CII]158 / [OI]63 line ratios well. Fast ( V shock = − ) J-shocks at low densities ( n = cm − ) and / orPDR models ( n > cm − , G > ) can reproduce thecombined [OI]63 / [OI]145 and [CII]158 / [OI]63 line ratios. – The main contribution to the [CII] 157.74 µ m line most prob-ably come from PDR emission from the disc and its sur-roundings. A (small) contribution from shocks is also ex-pected, as detections only in outflow sources and shock mod-els suggest. The precise origin of the [CII] line has to be bet-ter constrained with spatially and spectroscopically resolvedobservations. – The observed correlations support the interpretation ofjet / outflows (when present) as the dominant contributor tothe FIR line emission, and points to a common excitationmechanism, i.e. shocks. The broad wings of the [OI] 63.18 µ m line and its correlation with [OI] 6300 Å suggest the pres-ence of several velocity components, i.e. a jet origin for theHVC and a combination of disc, envelope, and wind for theLVC. – Molecular line emission (H O, CO, and OH) is also mainlydetected in outflow systems. Slow ( V shock = − ) C-type shocks with densities between n = and 10 cm − canaccount for these fluxes; with emitting areas ranging fromtens to hundreds of au. The OH / H O line ratios are typicallyoverestimated in J-type shocks. – The correlations with photometric bands (24 and 70 µ m) in-dicate that the contribution from the disc to the [OI] 63.18 µ m line flux may be up to ∼
50% for the jet / outflow sources,strongly depending on the evolutionary status of the source.When the jet / outflow activity decreases, the disc contribu-tion relative to the line fluxes increases significantly ( > – Although there are clear indications that the emission isdominated by the outflow, massive discs and / or low dust-to-gas ratios may also explain the observed high molecular linefluxes.The low spectral and spatial resolution from PACS are not suf-ficient to unambiguously determine what fraction of the lineemission comes from the disc, outflow, or the surrounding enve-lope. Spatially and / or velocity resolved observations are neededto pin-point the origin of the emission lines. In this regard, theinstruments on board SOFIA may have the potential to resolvethe brightest lines. Models which include both a disc and ajet / outflow and their interaction are needed to accurately inter-pret the multiwavelength observations of young T Tauri stars. Acknowledgements.
M. Alonso-Martinez, C. Eiroa, and G. Meeus are par-tially supported by AYA2011–26202 and AYA2014–55840–P. G.M. is sup-ported by RyC–2011–07920. P.R.M. acknowledges funding from the ESAResearch Fellowship Programme. I.K. acknowledges funding from the EuropeanUnion Seventh Framework Programme FP7-2011 under grant agreement No.284405. L.P. has received funding from the European Union Seventh FrameworkProgramme (FP7 / References
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In Table A.1 we list the literature properties of the sample includ-ing SED classes, spectral types, e ff ective temperatures, mass ac-cretion rates, stellar luminosities, X-ray luminosities, accretionluminosities, ages, disc masses, and separation if the sources aremultiple systems. The targets are split into outflows and non-outflows depending on their signs of jet / outflow activity as ex-plained in the text (see Sect. 2.1). l on s o– M a r t ´ ı n ez : H e r s c h e l GA SPS s p ec t r a l ob s e r v a ti on s o f TT a u r i s t a r s i n T a u r u s Table A.1.
Sample Properties [1] [2] [3] [4] [5] [6] [7] [8] [9] [10] [11]Target SED Class SpT a T a ˙ M d acc L c ∗ L eX L d acc Age c M f disc Sep. (Pair) g [–] [–] [–] [K] [10 − M ⊙ yr − ] [L ⊙ ] [L ⊙ ] [L ⊙ ] [Myr] [M ⊙ ] [arcsec]OutflowsAA Tau II M0.6 3770 2.51 0.74 1.240 0.21 2.7 0.01 ...CW Tau II K3 4470 5.27 1.35 2.844 1.01 5.8 0.002 ...DF Tau II M2.7 3450 10.05 1.60 ... 0.22 0.1 0.0004 0.09 (A–B)DG Tau II K7 4020 25.25 0.90 ... 1.26 0.6 0.02 ...DG Tau B I M? b < < / Flat M2.4 3510 ... 0.32 3.224 ... 2.8 0.002 0.23 (A–B)GG Tau II K7.5(A) + M5.8(B) 3930(A) + + + c ... 0.02 17.551 ... ... ... 0.23 (A–B)Haro 6–13 II M1.6(E) + K5.5(W) 3670(E) + + M4.8(B) 4470(A) + + c ... 0.41 0.882 ... 0.9 ... ...IRAS 04358 + c ... 0.36 0.401 ... ... ... ...RW Aur II K0(A) + K6.5(B) 5110(A) + / I K0 5250 c + M3.1 3560(A) + + + c ... 0.29 0.178 ... 0.5 < / + M1.7(B) 3960(A) + l on s o– M a r t ´ ı n ez : H e r s c h e l GA SPS s p ec t r a l ob s e r v a ti on s o f TT a u r i s t a r s i n T a u r u s Table A.1. (Continued) [1] [2] [3] [4] [5] [6] [7] [8] [9] [10] [11]Target SED Class SpT a T a ˙ M d acc L c ∗ L eX L d acc Age c M f disc Sep. (Pair) g [–] [–] [–] [K] [10 − M ⊙ yr − ] [L ⊙ ] [L ⊙ ] [L ⊙ ] [Myr] [M ⊙ ] [arcsec]Non-outflowsGH Tau II M2.3 3515 ... 0.79 0.109 ... 0.5 0.0007 0.31 (A–B)GI / GK Tau II M0.4 / K6.5(A) 3770 / / / / / / + M0.9(B) 4050(A) + c ... ... ... ... ... < c ... 0.17 ... 0.01 ... < < c + M3.7(B) 3770(A) + < < < < + M1.9(W) 4870(E) + + + + M1.7(B) 3350(A) + + + / III K8 3990 0.27 0.81 2.205 0.03 2.3 0.0004 ...V836 Tau T M0.8 3740 0.1 0.51 ... 0.01 7.9 ... ...V927 Tau III M4.9 2990 ... 0.33 ... ... 0.3 < c ... 1.30 4.139 ... 0.5 < < Notes. [1] Sample targets, [2] SED Class, [3] spectral type, [4] temperature, [5] accretion rate, [6] stellar luminosities, [7] X-ray luminosities, [8] accretion luminosities, [9] stellar age, [10] disc masses, [11] Separationin arcsec and components. The mass accretion rates were derived from the U band excess as follows: first the U band photometry from Kenyon & Hartmann (1995) was derredened using the Cardelli et al. (1989)extinction curve. Then the U band excess was derived by subtracting the photospheric contribution, and converted to accretion luminosity according to the empirical relation log ( L acc / L ⊙ ) = . × log ( L U / L ⊙ ) + .
98 inGullbring et al. (1998). The mass accretion rates were finally obtained using the relation M acc = L acc R ⋆ GM ⋆ (1 − R ⋆ / R in ) , following the assumptions in Gullbring et al. (1998). References. a Herczeg & Hillenbrand (2014), b see Howard et al. (2013) for a discussion, c Palla & Stahler (2002), d ˙ M acc and L acc as explained above , e G¨udel et al. (2007b), f Andrews & Williams (2005), g White & Ghez(2001). lonso–Mart´ınez: Herschel GASPS spectral observations of T Tauri stars in Taurus Appendix B: Observations
Table B.1 shows the complete Taurus GASPS list of obser-vations. The object coordinates, identifiers (OBSIDs) of line(LineSpec) and range (RangeSpec) observing modes along withexposure times are given.
Table B.1.
Overview of the spectroscopic OBSIDs that were observed. (D) stands for deeper observations than the regular settings.(D1): observations at 72 and 145 µ m; (D2): observations at 72, 79, 145, and 158 µ m; (D3): observations at 72, 78, 90, 145, 158, and180 µ m. Star RA[h:m:s] DEC[d:m:s] LineSpec ID RangeSpec ID t exp [s]AA Tau 04:34:55.42 + / / + + + / + / + + / + + / + / + / + / + + + / + / + / + / + / + / + / + + + / + / + / ∗ + / + / + + + / ⋆ + / / K Tau † + + / + + / + + + + + + / + / + + + + / + + / + + / + Table B.1. (Continued)
Star RA[h:m:s] DEC[d:m:s] LineSpec ID RangeSpec ID t exp [s]LkCa 1 04:13:14.14 + + + + + + / + / + / + / + / + + / + / + + + + + + + ‡ + / + Notes. + DI Tau was observed in the same FOV as DH Tau ∗ Haro 6–5B was observed in the same FOV as FS Tau ⋆ V807 Tau was observed in the same FOV as GH Tau † HL Tau was observed in the same FOV as XZ Tau ‡ GI / K Tau were spectroscopically unresolved20lonso–Mart´ınez: Herschel GASPS spectral observations of T Tauri stars in Taurus
Appendix C: Line fluxes
Tables C.1 and C.2 give the aperture corrected line fluxes inte-grated over one spaxel in the 60–80 µ m and 90–180 µ m ranges,respectively. The same holds for Tables C.3 and C.4, but for thefluxes integrated over 3 × l on s o– M a r t ´ ı n ez : H e r s c h e l GA SPS s p ec t r a l ob s e r v a ti on s o f TT a u r i s t a r s i n T a u r u s Table C.1.
Line fluxes for species in the 60-80 µ m range integrated over one spaxel. Target SED [OI] 63.18 µ m o–H O 63.32 µ m CO 72.84 µ m o–H O 78.74 µ m p–H O 78.92 µ m OH 79.11 µ m OH 79.18 µ m CO 79.36 µ m– Class (10 − W / m ) (10 − W / m ) (10 − W / m ) (10 − W / m ) (10 − W / m ) (10 − W / m ) (10 − W / m ) (10 − W / m )OutflowsAA Tau II 0.119 (0.038) 0.067 (0.021) < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < / Flat < < < < < < < < < < < < < < < < < < < < < < < < + < < < < < < < + < < < < < < < < < < < < < < < < < < < < < < < < < / I < < < < < < < < < < < < < < < < < < < < < < < / < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < l on s o– M a r t ´ ı n ez : H e r s c h e l GA SPS s p ec t r a l ob s e r v a ti on s o f TT a u r i s t a r s i n T a u r u s Table C.1. (Continued)
Target SED [OI] 63.18 µ m o–H O 63.32 µ m CO 72.84 µ m o–H O 78.74 µ m p–H O 78.92 µ m OH 79.11 µ m OH 79.18 µ m CO 79.36 µ m– Class (10 − W / m ) (10 − W / m ) (10 − W / m ) (10 − W / m ) (10 − W / m ) (10 − W / m ) (10 − W / m ) (10 − W / m )Non-outflowsGH Tau II < < < < < < < < / GK Tau II 0.127 (0.031) 0.069 (0.022) – – – – – –GM Aur T 0.241 (0.045) < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < / III < < < < < < < < < < < < Notes.
Fluxes in bold text represent misaligned and / or extended sources. In these cases the 3x3 fluxes (see Tables C.3 and C.4) are more accurate and are used in the calculations. l on s o– M a r t ´ ı n ez : H e r s c h e l GA SPS s p ec t r a l ob s e r v a ti on s o f TT a u r i s t a r s i n T a u r u s Table C.2.
Line fluxes for species in the 90–180 µ m range integrated over one spaxel. Target SED p–H O 89.99 µ m CO 90.16 µ m p–H O 144.52 µ m CO 144.78 µ m [OI] 145.52 µ m [CII] 157.74 µ m o–H O 179.53 µ m o–H O 180.49 µ m– Class (10 − W / m ) (10 − W / m ) (10 − W / m ) (10 − W / m ) (10 − W / m ) (10 − W / m ) (10 − W / m ) (10 − W / m )OutflowsAA Tau II < < < < < < < < < < < < < < < < < < < < < < < ‡ < < < < < < < < < < < < < < < < ‡ < < < < < / Flat 0.056 (0.014) 0.059 (0.017) < ‡ < < < < < < < < ‡ < < < < < < < < < < < < < < < < < < + < < < < < < + < ‡ – –RW Aur II < < < < < ‡ < < < < < ‡ < < / I 3.408 (0.057)
UY Aur II 0.098 (0.028) 0.103 (0.023) 0.048 (0.010) 0.254 (0.010) 0.171 (0.010) ‡ < < < < < < < < < ‡ < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < Notes.
Fluxes in bold text represent misaligned and / or extended sources. In these cases the 3x3 fluxes (see Tables C.3 and C.4) are more accurate and are used in the calculations. ( ‡ ) Objects shown in Fig. 14 with [CII]contamination in the o ff - source positions. l on s o– M a r t ´ ı n ez : H e r s c h e l GA SPS s p ec t r a l ob s e r v a ti on s o f TT a u r i s t a r s i n T a u r u s Table C.3.
Line fluxes for species in the 60-80 µ m range integrated over 3 × Target SED [OI] 63.18 µ m o-H O 63.32 µ m CO 72.84 µ m o-H O 78.74 µ m p-H O 78.92 µ m OH 79.11 µ m OH 79.18 µ m CO 79.36 µ m– Class (10 − W / m ) (10 − W / m ) (10 − W / m ) (10 − W / m ) (10 − W / m ) (10 − W / m ) (10 − W / m ) (10 − W / m )OutflowsAA Tau II 0.21 (0.06) < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < / Flat 5.34 (0.19) < < < < < < < < < < < < < < < < < < < < < < < < < < < + < < < < < < < + < < < < < < < < < < < < < < < < < < < < < < < < < < < < / I 183.45 (3.87) < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < / < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < Table C.3. (Continued)
Target SED [OI] 63.18 µ m o-H O 63.32 µ m CO 72.84 µ m o-H O 78.74 µ m p-H O 78.92 µ m OH 79.11 µ m OH 79.18 µ m CO 79.36 µ m– Class (10 − W / m ) (10 − W / m ) (10 − W / m ) (10 − W / m ) (10 − W / m ) (10 − W / m ) (10 − W / m ) (10 − W / m )Non-outflowsGH Tau II < < < < < < < < / GK Tau II < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < / III < < < < < < < < < < < < l on s o– M a r t ´ ı n ez : H e r s c h e l GA SPS s p ec t r a l ob s e r v a ti on s o f TT a u r i s t a r s i n T a u r u s Table C.4.
Line fluxes for species in the 90–180 µ m range integrated over 3 × Target SED p-H O 89.99 µ m CO 90.16 µ m p-H O 144.52 µ m CO 144.78 µ m [OI] 145.52 µ m [CII] 157.74 µ m o-H O 179.53 µ m o-H O 180.49 µ m– Class (10 − W / m ) (10 − W / m ) (10 − W / m ) (10 − W / m ) (10 − W / m ) (10 − W / m ) (10 − W / m ) (10 − W / m )OutflowsAA Tau II < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < / Flat < < < < < < < < < < < < < < < < < < < < < < < < < < < < + < < < < < < < < + < < < < < < < < < < < < < < < < < < < < < < < < < < < / I 3.89 (0.17) 4.24 (0.15) 1.28 (0.13) 11.93 (0.12) 7.93 (0.13) 5.65 (0.17) 6.35 (0.12) 2.05 (0.14)UY Aur II < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < lonso–Mart´ınez: Herschel GASPS spectral observations of T Tauri stars in Taurus Appendix D: Detection fractions
Tables D.1, D.2, and D.3 show respectively the detection frac-tions statistics in terms of outflow / non-outflow, SED classes andspectral types. Table D.1.
Detection fractions for the entire sample, outflow sources only, and non-outflow sources only.
Species λ ( µ m ) All Outflow Non-outflowAtomic[OI] 63.18 55 + − %(42 /
76) 100 + − %(27 /
27) 31 + − %(15 / + − %(19 /
39) 75 + − %(18 /
24) 7 + − %(1 / ⋆ [CII] 157.74 34 + − %(13 /
38) 54 + − %(13 /
24) –(0 / † Molecularo–H O 63.32 17 + − %(13 /
76) 37 + − %(10 /
27) 6 + − %(3 / ‡ o–H O 78.74 32 + − %(12 /
38) 50 + − %(12 /
24) –(0 / O 179.53 17 + − %(5 /
30) 23 + − %(5 /
22) –(0 / O 180.49 10 + − %(3 /
30) 14 + − %(3 /
22) –(0 / O 78.92 18 + − %(7 /
38) 29 + − %(7 /
24) –(0 / O 89.99 30 + − %(9 /
30) 41 + − %(9 /
22) –(0 / O 144.52 23 + − %(9 /
39) 38 + − %(9 /
24) –(0 / + − %(2 /
39) 8 + − %(2 /
24) –(0 / + − %(2 /
38) 8 + − %(2 /
24) –(0 / + − %(8 /
30) 36 + − %(8 /
22) –(0 / + − %(22 /
39) 79 + − %(19 /
24) 20 + − %(3 / + − %(13 /
38) 54 + − %(13 /
24) –(0 / + − %(12 /
38) 50 + − %(12 /
24) –(0 / Notes. ( ⋆ ) Detections in DE Tau. ( ‡ ) Detections in: BP Tau (deeper observations), GI / GK Tau (GK Tau is outflow GI Tau is not), IQ Tau. Theparentheses indicate the ratio of targets where a certain line is detected over the total number of targets observed. The hyphens mean no detections.28lonso–Mart´ınez: Herschel GASPS spectral observations of T Tauri stars in Taurus
Table D.2.
Detection fractions in terms of SED Class.
Species λ ( µ m ) Class I Class II Class TD Class IIIAtomic[OI] 63.18 100 + − %(5 /
5) 70 + − %(31 /
44) 60 + − %(6 /
10) –(0 / + − %(4 /
4) 48 + − %(14 /
29) 17 + − %(1 / + − %(2 /
4) 36 + − %(10 /
28) 17 + − %(1 / O 63.32 40 + − %(2 /
5) 23 + − %(10 /
44) 10 + − %(1 /
10) –(0 / O 78.74 50 + − %(2 /
4) 32 + − %(9 /
28) 17 + − %(1 / O 179.53 –(0 /
4) 24 + − %(5 /
21) –(0 / O 180.49 –(0 /
4) 14 + − %(3 /
21) –(0 / O 78.92 50 + − %(2 /
4) 18 + − %(5 /
28) –(0 / O 89.99 25 + − %(1 /
4) 33 + − %(7 /
21) 20 + − %(1 / O 144.52 50 + − %(2 /
4) 21 + − %(6 /
29) 17 + − %(1 / /
4) 7 + − %(2 /
29) –(0 / /
4) 7 + − %(2 /
28) –(0 / + − %(1 /
4) 33 + − %(7 /
21) –(0 / + − %(4 /
4) 59 + − %(17 /
29) 17 + − %(1 / + − %(2 /
4) 39 + − %(11 /
28) –(0 / + − %(2 /
4) 36 + − %(10 /
28) –(0 / Table D.3.
Detection fractions in terms of SpT.
Species λ ( µ m ) G0-K7 K7-M1 M1-M3 M3-M6Atomic[OI] 63.18 84 + − %(16 /
19) 61 + − %(11 /
18) 35 + − %(7 /
20) 42 + − %(8 / + − %(8 /
13) 50 + − %(4 /
8) 38 + − %(3 /
8) 40 + − %(4 / + − %(6 /
13) 25 + − %(2 /
8) 25 + − %(2 /
8) 33 + − %(3 / O 63.32 32 + − %(6 /
19) 17 + − %(3 /
18) 10 + − %(2 /
20) 11 + − %(2 / O 78.74 31 + − %(4 /
13) 38 + − %(3 /
8) 38 + − %(3 /
8) 22 + − %(2 / O 179.53 23 + − %(3 /
13) –(0 /
5) 20 + − %(1 /
5) 14 + − %(1 / O 180.49 8 + − %(1 /
13) –(0 /
5) 20 + − %(1 /
5) 14 + − %(1 / O 78.92 23 + − %(3 /
13) 13 + − %(1 /
8) 13 + − %(1 /
8) 22 + − %(2 / O 89.99 54 + − %(7 /
13) –(0 /
5) 20 + − %(1 /
5) 14 + − %(1 / O 144.52 31 + − %(4 /
13) –(0 /
8) 25 + − %(2 /
8) 30 + − %(3 / + − %(2 /
13) –(0 /
8) –(0 /
8) –(0 / + − %(2 /
13) –(0 /
8) –(0 /
8) –(0 / + − %(5 /
13) –(0 /
5) 40 + − %(2 /
5) 14 + − %(1 / + − %(9 /
13) 50 + − %(4 /
8) 63 + − %(5 /
8) 40 + − %(4 / + − %(5 /
13) 38 + − %(3 /
8) 38 + − %(3 /
8) 22 + − %(2 / + − %(5 /
13) 25 + − %(2 /
8) 38 + − %(3 /
8) 22 + − %(2 /
9) 29lonso–Mart´ınez: Herschel GASPS spectral observations of T Tauri stars in Taurus
Appendix E: Diameter of emitting regions
As shown in Flower & Pineau Des Forˆets (2010), the relationbetween the emergent flux, F e , in erg cm − s − and T d V in Kkm s − is F e = × π k B λ T dV , (E.1)where λ is the wavelength of the transition in cm, and T d V asprovided in Flower & Pineau des Forˆets (2015).The flux observed at the Earth, F , then is F = F e Ω π , (E.2)where Ω str is the solid angle subtended by the source. Giventhat Ω ∼ D / r , where r and D are the distance to the source andthe diameter of the emitting area in au, respectively, then D ∼ s π r F F e . (E.3)The diameter of the emitting regions provided (see Table 4)are inferred using only the brightest lines among the ratios (i.e.o-H O 78.74 µ m, CO 144.78 µ m, and OH 79.12 µ m), besides be-ing compatible with the shock velocities ( V shock ) and pre-shockdensities ( n ) observed in Figs. E.1, E.2, E.3, and E.4. Fig. E.1.
Detail of the observed o-H O molecular line ratios(78 /
179 and 180 / blue ) and C-type ( red )shock models from Flower & Pineau des Forˆets (2015) for indi-vidual sources. The numbers refer to shock velocities ( V shock ).Di ff erent symbols refer to pre-shock densities ( n ). See legendfor details. Fig. E.2.
Same as Fig. E.1, but for CO and o-H O / CO ratios.
Fig. E.3.
Same as Fig. E.1, but for CO and o-H O / CO line ratios.
Fig. E.4.
Same as Fig. E.1, but for OH / o-H O line ratios.
Appendix F: Spectra 60-190 µ m Figures F.1 to F.7 show the objects in the sample with detectionsin the range 60–190 µ m. Fig. F.1.
Continuum subtracted spectra at 63 µ m for all objects with [OI] detections. The red vertical lines indicate the positions of[OI] 63.18 µ m and o-H O 63.32 µ m. The source GI / GK Tau is counted twice to make a total of 42 [OI] detections.
Fig. F.1. (Continued.)
Fig. F.2.
Continuum subtracted spectra at 72 µ m for all objects with CO detections. The red vertical lines indicate the positions ofCH + µ m and CO 72.84 µ m. Fig. F.3.
Continuum subtracted spectra at 78 µ m for all objects with detections. The red vertical lines indicate the positions of o-H O78.74 µ m, p-H O 78.92 µ m, OH 78.12 + µ m, and CO 79.36 µ m. Fig. F.4.
Continuum subtracted spectra at 90 µ m for all objects with detections. The red vertical lines indicate the positions of p-H O89.99 µ m and CO 90.16 µ m. Fig. F.5.
Continuum subtracted spectra at 145 µ m for all objects with detections. The red vertical lines indicate the positions ofp-H O 144.52 µ m, CO 144.78 µ m, and [OI] 145.52 µ m. Fig. F.6.
Continuum subtracted spectra at 145 µ m for all objects with detections. The red vertical lines indicate the positions of [CII]157.74 µ m and p-H O 158.31 µ m. Fig. F.7.
Continuum subtracted spectra at 180 µ m for all objects with detections. The red vertical lines indicate the positions ofo-H O 179.53 µ m and o-H O 180.49 µ m.m.