Far-infrared molecular lines from Low- to High-Mass Star Forming Regions observed with Herschel
A. Karska, F. Herpin, S. Bruderer, J.R. Goicoechea, G.J. Herczeg, E.F. van Dishoeck, I. San José-García, A. Contursi, H. Feuchtgruber, D. Fedele, A. Baudry, J. Braine, L. Chavarría, J. Cernicharo, F.F.S. van der Tak, F. Wyrowski
aa r X i v : . [ a s t r o - ph . S R ] N ov Astronomy&Astrophysicsmanuscript no. karska˙highmass c (cid:13)
ESO 2018September 28, 2018
Far-infrared molecular lines from Low- to High-MassStar Forming Regions observed with Herschel
A. Karska , , F. Herpin , , S. Bruderer , J.R. Goicoechea , G.J. Herczeg , E.F. van Dishoeck , , I. San Jos´e-Garc´ıa ,A. Contursi , H. Feuchtgruber , D. Fedele , A. Baudry , , J. Braine , , L. Chavarr´ıa , J. Cernicharo , F.F.S. van derTak , , and F. Wyrowski Max-Planck Institut f¨ur Extraterrestrische Physik (MPE), Giessenbachstr. 1, D-85748 Garching, Germany Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA Leiden, The Netherlands Universit´e de Bordeaux, Observatoire Aquitain des Sciences de l’Univers, 2 rue de l’Observatoire, BP 89, F-33271 Floirac Cedex,France CNRS, LAB, UMR 5804, F-33271 Floirac Cedex, France Centro de Astrobiolog´ıa. Departamento de Astrof´ısica. CSIC-INTA. Carretera de Ajalvir, Km 4, Torrej´on de Ardoz. 28850, Madrid,Spain Kavli Institut for Astronomy and Astrophysics, Yi He Yuan Lu 5, HaiDian Qu, Peking University, Beijing, 100871, PR China SRON Netherlands Institute for Space Research, PO Box 800, 9700 AV, Groningen, The Netherlands Kapteyn Astronomical Institute, University of Groningen, PO Box 800, 9700 AV, Groningen, The Netherlands Max-Planck-Institut f¨ur Radioastronomie, Auf dem H¨ugel 69, 53121 Bonn, Germanye-mail: [email protected]
Received May 24, 2013; accepted November 26, 2013
ABSTRACT
Aims.
Our aim is to study the response of the gas to energetic processes associated with high-mass star formation and compare it withpreviously published studies on low- and intermediate-mass young stellar objects (YSOs) using the same methods. The quantifiedfar-infrared line emission and absorption of CO, H O, OH, and [O i ] reveals the excitation and the relative contribution of di ff erentatomic and molecular species to the gas cooling budget. Methods.
Herschel-PACS spectra covering 55–190 µ m are analyzed for ten high-mass star forming regions of luminosities L bol ∼ − L ⊙ and various evolutionary stages at spatial scales of ∼ AU. Radiative transfer models are used to determine thecontribution of the quiescent envelope to the far-IR CO emission.
Results.
The close environments of high-mass protostars show strong far-infrared emission from molecules, atoms, and ions. Water isdetected in all 10 objects even up to high excitation lines, often in absorption at the shorter wavelengths and in emission at the longerwavelengths. CO transitions from J = −
13 up to typically 29 −
28 ( E u / k B ∼ − T rot ∼
300 K. Typical H O excitation temperatures are T rot ∼
250 K, while OH has T rot ∼
80 K. Far-IRline cooling is dominated by CO ( ∼ i ] ( ∼
20 %), which becomes more important for the mostevolved sources. H O is less important as a coolant for high-mass sources due to the fact that many lines are in absorption.
Conclusions.
Emission from the quiescent envelope is responsible for ∼ −
85 % of the total CO luminosity in high-mass sourcescompared with only ∼
10% for low-mass YSOs. The highest − J lines ( J up ≥
20) originate most likely from shocks, based on the strongcorrelation of CO and H O with physical parameters ( L bol , M env ) of the sources from low- to high-mass YSOs. Excitation of warmCO described by T rot ∼
300 K is very similar for all mass regimes, whereas H O temperatures are ∼
100 K higher for high-masssources compared with low-mass YSOs. The total far-IR cooling in lines correlates strongly with bolometric luminosity, consistentwith previous studies restricted to low-mass YSOs. Molecular cooling (CO, H O, and OH) is ∼ − and 10 − times lower than the dust luminosityfor the low- and high-mass star forming regions, respectively. Key words. astrochemistry stars: formation stars: –ISM: outflows, shocks
1. Introduction
High-mass stars ( M > ⊙ ) play a central role in the energybudget, the shaping, and the evolution of galaxies (see reviewby Zinnecker & Yorke 2007). They are the main source of UVradiation in galaxy disks. Massive outflows and H ii regions arepowered by massive stars and are responsible for generating tur-bulence and heating of the interstellar medium (ISM). At the endof their lives, they inject heavy elements into the ISM that formthe next generation of molecules and dust grains. These atomsand molecules are the main cooling channels of the ISM. Themodels of high-mass star formation are still strongly debated; the two competing scenarios are turbulent core accretion and‘competitive accretion’ (e.g. Cesaroni 2005). Molecular line ob-servations are crucial to determine the impact of UV radiation,outflows, infall and turbulence on the formation and evolutionof the high-mass protostars and ultimately distinguish betweenthose models.Based on observations, the ‘embedded phase’ of high-massstar formation may empirically be divided into several stages(e.g. Helmich & van Dishoeck 1997; van der Tak et al. 2000;Beuther et al. 2007): (i) massive pre-stellar cores (PSC); (ii)high-mass protostellar objects (HMPOs); (iii) hot molecularcores (HMC); and (iv) ultra-compact H ii regions (UCH ii ). The
1. Karska et al. 2013: Far-infrared molecular lines from Low- to High-Mass Star Forming Regions observed with Herschel pre-stellar core stage represents initial conditions of high-massstar formation, with no signatures of outflow / infall or maseractivity. During the high-mass protostellar objects stage, infallof a massive envelope onto the central star and strong outflowsindicate the presence of an active protostar. In the hot molecularcore stage, large amounts of warm and dense gas and dust areseen. The temperature of T >
100 K in subregions < . ff the grains. In thefinal, ultra-compact H ii regions stage, a considerable amount ofionized gas is detected surrounding the central protostar.The above scenario is still debated (Beuther et al. 2007), inparticular whether stages (ii) and (iii) are indeed intrinsically dif-ferent. The equivalent sequence for the low-mass Young StellarObjects (hereafter YSO) is better established (Shu et al. 1987;Andr´e et al. 1993, 2000). The ‘embedded phase’ of low-massprotostars consists of: (i) pre-stellar core phase; (ii) Class 0; and(iii) Class I phase. The Class 0 YSOs are surrounded by a mas-sive envelope and drive collimated jets / outflows. In the moreevolved Class I objects the envelope is mostly dispersed andmore transparent for UV radiation; the outflows are less pow-erful and have larger opening angles.Low mass sources can be probed at high spatial resolutiondue to a factor of 10 smaller distances, which allow us to studywell-isolated sources and avoid much of the confusion due toclouds along the line of sight. The line emission is less a ff ectedby foreground extinction and therefore provides a good tool tostudy the gas physical conditions and chemistry in the region.The slower evolutionary timescale results in a larger numberof low-mass YSOs compared to their high-mass counterparts,which is also consistent with observed stellar / core mass func-tions.While low-mass YSOs are extensively studied in thefar-infrared, first with the Infrared Space Observatory (ISO,Kessler et al. 1996) and now with Herschel (Pilbratt et al.2010) , the same is not the case for high-mass sources(see e.g. Helmich & van Dishoeck 1997; Vastel et al. 2001;Boonman & van Dishoeck 2003). For those, the best-studiedcase is the relatively nearby Orion BN-KL region observed withISO’s Long- and Short-Wavelength Spectrometers (Clegg et al.1996; de Graauw et al. 1996). Spectroscopy at long-wavelengths(45-197 µ m) shows numerous and often highly-excited H Olines in emission (Harwit et al. 1998), high- J CO lines (e.g., upto J = O vibration-rotation bands and H pure rotational lines (van Dishoeck et al. 1998; Rosenthal et al.2000). Fabry-Perot (FP) spectroscopy ( λ /∆ λ ∼ − ) data show resolved P-Cygni profiles for selected H O tran-sitions at λ < µ m, with velocities extending up to 100 kms − (Wright et al. 2000; Cernicharo et al. 2006). At the shortestwavelengths ( < µ m) all pure rotational H O lines show ab-sorption (Wright et al. 2000). ISO spectra towards other high-mass star forming regions are dominated by atomic and ioniclines (see review by van Dishoeck 2004), similar to far-IR spec-tra of extragalactic sources (Fischer et al. 1999; Sturm et al.2002).The increased sensitivity, spectral and spatial resolutionof the Photodetector Array Camera and Spectrometer (PACS)(Poglitsch et al. 2010) onboard
Herschel now allows the detailedstudy of the molecular content of a larger sample of high-mass Herschel is an ESA space observatory with science instruments pro-vided by European-led Principal Investigator consortia and with impor-tant participation from NASA. star forming regions. In particular, the more than an order ofmagnitude improvement in the spectral resolution in compari-son with the ISO-LWS grating observing mode allows the rou-tine detections of weak lines against the very strong continuumof high-mass sources with
Herschel , with line-to-continuum ra-tios below 1%.The diagnostic capabilities of far-infrared lines havebeen demonstrated by the recent results on low- andintermediate-mass YSOs and their outflows (Fich et al. 2010;Herczeg et al. 2012; Goicoechea et al. 2012; Manoj et al. 2013;Wampfler et al. 2013; Karska et al. 2013; Green et al. 2013).The CO ladder from J = O lines with a range of excitation energies are de-tected towards the Class 0 sources, NGC1333 IRAS4B andSerpens SMM1 (Herczeg et al. 2012; Goicoechea et al. 2012).The highly-excited H O 8 –7 line at 63.3 µ m ( E u / k B = ∼ ⊙ . Non-dissociative shocks and to a smaller extentUV-heating are suggested to be the dominant physical processesresponsible for the observed line emission (van Kempen et al.2010; Visser et al. 2012; Karska et al. 2013). The contributionfrom the bulk of the quiescent warm protostellar envelope tothe PACS lines is negligible for low-mass sources. Even for theintermediate-mass source NGC7129 FIRS2, where the envelopecontribution is higher, the other processes dominate (Fich et al.2010).In this paper, we present Herschel-PACS spectroscopy of10 sources that cover numerous lines of CO, H O, OH, and[O i ] lines obtained as part of the ‘Water in star forming re-gions with Herschel’ (WISH) key program (van Dishoeck et al.2011). WISH observed in total about 80 protostars at di ff er-ent evolutionary stages (from prestellar cores to circumstellardisks) and masses (low-, intermediate- and high-mass) withboth the Heterodyne Instrument for the Far-Infrared (HIFI;de Graauw et al. 2010) and PACS (Poglitsch et al. 2010). Thispaper focusses only on PACS observations of high-mass YSOs.It complements the work by van der Tak et al. (2013), which de-scribes our source sample and uses HIFI to study spectrally re-solved ground-state H O lines towards all our objects. That pa-per also provides updated physical models of their envelopes.The results on high-mass YSOs will be compared with thosefor low- and intermediate-mass young stellar objects, analyzedin a similar manner (Karska et al. 2013; Wampfler et al. 2013;Fich et al. 2010) in order to answer the following questions: Howdoes far-IR line emission / absorption di ff er for high-mass pro-tostars at di ff erent evolutionary stages? What are the dominantgas cooling channels for those sources? What physical compo-nents do we trace and what gas conditions cause the excitationof the observed lines? Are there any similarities with the low-and intermediate-mass protostars?The paper is organized as follows: § § § § §
2. Observations and data reduction
We present spectroscopy observations of ten high-mass starforming regions collected with the PACS instrument on board
Herschel in the framework of the ‘WISH’ program. The sources
2. Karska et al. 2013: Far-infrared molecular lines from Low- to High-Mass Star Forming Regions observed with Herschel [ O I] [ O III] [ N II] [ O I] [ C II]
Fig. 1.
Herschel-PACS continuum-normalized spectrum of W3 IRS5 at the central position. Lines of CO are shown in red, H Oin blue, OH in light blue, CH in orange, and atoms and ions in green. Horizontal magenta lines show spectral regions zoomed inFigure C.1.have an average distance of h D i = ii regions (UC H ii ). The list of sources and their ba-sic properties are given in Table 1. Objects are shown in the se-quence of increasing value of an evolutionary tracer, L . M − ,introduced in Bontemps et al. (1996). The sequence does not al-ways correspond well with the evolutionary stages most com-monly assigned to the sources in the literature (last column ofTable 2), perhaps because multiple objects in di ff erent evolu-tionary stages are probed within our spatial resolution element(see e.g. Wyrowski et al. 2006; Leurini et al. 2013, for the caseof G327-0.6).PACS is an integral field unit with a 5 × spaxels ). Each spaxel covers 9 . ′′ × . ′′
4, providinga total field of view of ∼ ′′ × ′′ . The focus of this work is onthe central spaxel only. The central spaxel probes similar physi-cal scales as the full 5 × × µ m wavelength range with Table 1.
Catalog information and source properties.
Object
D L bol M env L . M − Class(kpc) ( L ⊙ ) ( M ⊙ ) ( L . ⊙ M − ⊙ )G327-0.6 3.3 7.3 10 ii DR21(OH) 1.5 1.3 10
472 0.62 HMPOW33A 2.4 3.0 10
698 0.70 HMPOG34.26 + ii NGC6334-I 1.7 1.1 10
750 1.41 HMCNGC7538-I1 2.7 1.1 10
433 2.45 UCH ii AFGL2591 3.3 1.2 10
373 2.99 HMPOW3-IRS5 2.0 2.1 10
424 3.68 HMPOG5.89-0.39 1.3 4.1 10
140 4.18 UCH ii Notes.
Source coordinates with references and their physical parame-ters are taken from van der Tak et al. (2013).
Nyquist sampling of the spectral elements. The wavelength cov-erage consists of three grating orders (1st: 102-210 µ m, 2nd:71-105 µ m or 3rd: 51-73 µ m), two of which are always ob-served simultanously (one in the blue, λ < µ m, and one in
3. Karska et al. 2013: Far-infrared molecular lines from Low- to High-Mass Star Forming Regions observed with Herschel
Fig. 2.
Herschel-PACS profiles of the [O i ] 63.2 µ m line at cen-tral position.the red, λ > µ m, parts of the spectrum). The spectral resolv-ing power is R = λ/ ∆ λ ≈ ∼
75 to 300 km s − ).Two nod positions were used for chopping 3 ′ on each sideof the source. The comparison of the two positions was madeto assess the influence of the o ff -source flux of observed speciesfrom the o ff -source positions, in particular for atoms and ions.The [C ii ] fluxes are strongly a ff ected by the o ff -position fluxand saturated for most sources – we therefore limit our analysisof this species to the two sources with comparable results forboth nods, AFGL2591 and NGC7538-IRS1 (see Table 2).Typical pointing accuracy is better than 2 ′′ . However, twosources (G327-0.6 and W33A) were mispointed by a largeramount as indicated by the location of the peak continuum emis-sion on maps at di ff erent wavelengths (for the observing log seeTable A.1 in the Appendix). In order to account for the non-centric flux distribution on the integral field unit due to mis-pointing and to improve the continuum smoothness, for thesesources two spaxels with maximum continuum levels are used(spaxel 11 and 21 for G327 and 23 and 33 for W33A). Summinga larger number of spaxels was not possible due to a shift of lineprofiles from absorption to emission. The spatial extent of lineemission / absorption will be analyzed in future papers.We performed the basic data reduction with the HerschelInteractive Processing Environment v.10 (HIPE, Ott 2010). Theflux was normalized to the telescopic background and cali-brated using Neptune observations. Spectral flatfielding withinHIPE was used to increase the signal-to-noise (for details, seeHerczeg et al. 2012; Green et al. 2013). The overall flux calibra-tion is accurate to ∼ ff erent programs, cross-calibrations with HIFI and ISO, and continuum photome-try. Custom IDL routines were used to further process the dat-acubes. The line fluxes were extracted from the central spaxel(except G327-0.6 and W33A, see above) using Gaussian fitswith fixed line width (for details, see Herczeg et al. 2012). Next,they were corrected for the wavelength-dependent loss of radi-ation for a point source (see PACS Observer’s Manual ). Thatapproach is not optimal for the cases where emission is extendedbeyond the central spaxel, but that is mostly the case for atomiclines, which will be presented in a companion paper by Kwonet al. (in preparation, hereafter Paper II). The uncertainty intro-duced by using the point-source correction factors for extendedsources depends on the amount of emission in the surroundingring of spaxels. The continuum fluxes are calculated using all 25spaxels, except G327-0.6 where 1 spaxel was excluded due tosaturation. In most cases, the tabulated values are at wavelengthsnear bright lines. They are calculated using spectral regions onboth sides of the lines (but masking any features) and inter-polated linearly to the wavelength of the lines. The fluxes arepresented in Table B.1 in the Appendix. Our continuum fluxesare included in the spectral energy distribution fits presented invan der Tak et al. (2013), who used them to derive physical mod-els for all our sources. Those models and associated envelopemasses are used in this work in Sections 5.1 and 5.3.
3. Results
Figure 1 shows the full normalized PACS spectrum with lineidentifications for W3 IRS5, a high-mass protostellar objectwith the richest molecular emission among our sources. Carbonmonoxide (CO) transitions from J = J = J CO lines inFigures C.1 and C.2). Water vapor (H O) transitions up to E up ∼ − at 66.1 µ m, see blow-ups inFigure C.1). At wavelengths shortwards of ∼ µ m many H Olines are seen in absorption, but those at longer wavelengths areprimarily in emission.Seven hydroxyl (OH) doublets up to E up ≈
618 K are seen .All lines within the Π / ladder (119, 84, and 65 µ m doublets;see Figure 1 in Wampfler et al. 2013) are strong absorptionlines. The Π / ladder lines (163 and 71 µ m) are seen in emis-sion. The cross-ladder transitions at 79 µ m (OH / , / - / , / , E up ≈
180 K) and 96 µ m (OH / , / - / , / , E up ≈
270 K) areabsorption and emission lines, respectively. Only the ground-rotational lines of methylidyne (CH) are detected at 149 µ min absorption. The [O i ] transitions at 63 µ m and 145 µ m areboth strong emission lines in W3 IRS5. That is not always thecase for other sources in our sample. The [O i ] line at 63 µ m,where the velocity resolution of PACS is at its highest ( ∼ − ), shows a variety of profiles (see Figure 2): pure ab-sorption (G327-0.6, W51Ne1, G34.26, W33A), regular P-Cygniprofiles (AFGL2591, NGC6334-I), hints of inverse P-Cygni pro-files in DR21(OH) and pure emission (W3 IRS5, NGC7538-I1,G5.89). The [O i ] line at 145 µ m, however, is always detectedin emission. The P-Cygni profiles resolved at velocity scales of ∼
100 km s − resemble the high-velocity line wings observedin ro-vibrational transitions of CO in some of the same sources(Mitchell et al. 1990; Herczeg et al. 2011), suggested to origi- http: // herschel.esac.esa.int / Docs / PACS / html / pacs om.html The highest-excited OH doublet at 71 µ m is not considered furtherin the analysis, because the 70-73 µ m region observed with PACS isa ff ected by spectral leakage and thus is badly flux-calibrated.4. Karska et al. 2013: Far-infrared molecular lines from Low- to High-Mass Star Forming Regions observed with Herschel Fig. 3.
Normalized spectral regions of all our sources at the central position at 64-68, 148-157, and 169-182 µ m. Objects are shownin the evolutionary sequence, with the most evolved ones on top. Lines of CO are shown in red, H O in blue, OH in light blue, CHin orange, and NH and C in green. Spectra are shifted vertically to improve the clarity of the figure.nate in the wind impacting the outflow cavities. Note howeverthat not all absorption needs to be associated with the source:it can also be due to foreground absorption (e.g. outflow lobeor the ISM). For example, ISO Fabry Perot and new Herschel / HIFI observations of H O and OH in Orion also show line wingsin absorption / emission extending to velocities up to ∼
100 kms − (Cernicharo et al. 2006; Goicoechea et al. 2006, Choi et al.in prep.), but the ISO-LWS Fabry-Perot observations of Oriondid not reveal P-Cygni profiles in the [O i ] 63 µ m line.A comparison of selected line-rich parts of the spectra for allour sources is presented in Figure 3. The spectra are normalizedby the continuum emission to better visualize the line absorptiondepths. However, the unresolved profiles of PACS underestimatethe true absorption depths and cannot be used to estimate theoptical depths.The 64-68 µ m segment covers the highly-excited H O linesat 66.4, 67.1, and 67.3 µ m ( E up ≈
410 K); high- J CO lines at65.7 ( J = J = Π / J = / − / ( E up ≈
510 K) doublet at 65 µ m. The low-lying H Olines are detected for all sources. The high- J CO lines are notdetected in this spectral region. The OH doublet is detected for6 out of 10 sources (see also Figure C.3 in the Appendix). The main lines seen in the 148-157 µ m region are: CH Π / J = / – Π / J = / transition at 149 µ m (in absorption),CO 17-16, and H O 3 − line at 156.2 µ m ( E up ≈ ∼ − µ m seen towards thehot core G327-0.6 are most likely C ro-vibrational transitions(Cernicharo et al. 2000, Paper II).The most commonly detected lines in the 170-182 µ m spec-tral region include the CO 15-14 line and the H O lines at 174.6,179.5 and 180.5 µ m ( E up ≈ Olines change from object to object: the H O 2 − line at179.5 µ m is in absorption for all sources except W3 IRS5; theH O 2 − at 180.5 µ m is in emission for the three mostevolved sources (top 3 spectra on Figure 3). The ammonia lineat ∼ µ m, NH (3,2)a-(2,2)s, is detected towards 4 sources(G327-0.6, W51N-e1, G34.26, and NGC6334-I). An absorptionline at ∼ µ m corresponds to H O 2 − and / or H O + − − + lines (Goicoechea & Cernicharo 2001). Extended dis-cussion of ions and molecules other than CO, H O, and OH willappear in Paper II.Line profiles of H O observed with HIFI show a vari-ety of emission and absorption components that are not re-solved by PACS (Chavarr´ıa et al. 2010; Kristensen et al. 2012;
5. Karska et al. 2013: Far-infrared molecular lines from Low- to High-Mass Star Forming Regions observed with Herschel
Fig. 4.
Relative contributions of [O i ] (dark blue), CO (yellow),H O (orange), and OH (red) cooling to the total far-IR gas cool-ing at central position are shown from left to right horizontallyfor each source. The objects follow the evolutionary sequencewith the most evolved sources on top.van der Tak et al. 2013). The only H O lines observed in com-mon by the two instruments within the WISH program are theground-state transitions: 2 − at 179.5 µ m (1670 GHz) and2 − at 180.5 µ m (1661 GHz), both dominated by absorptionsand therefore not optimal to estimate to what extent a complexwater profile is diluted at the PACS spectral resolution. However,HIFI observations of the lines between excited rotational states,which dominate our detected PACS lines, are generally in emis-sion at the longer wavelengths probed by HIFI. In the case of CO 10–9, the line profiles of YSOs observed with HIFI con-sist of a broad outflow and a narrow quiescent component withthe relative fraction of the integrated intensity of the narrow tobroad components being typically 30-70% for low-mass sources(Yıldız et al. 2013). For the single case of a high-mass YSO, W3IRS5, this fraction is about 50% (San Jos´e-Garc´ıa et al. 2013).To summarize, PACS spectra of high-mass sources from oursample show detections of many molecular lines up to high ex-citation energies. CO, H O, OH, and CH lines are seen towardsall objects, whereas weaker lines of other molecules are detectedtowards less than half of the sources. CO lines are always seenin emission, CH in absorption, whereas other species show dif-ferent profiles depending on the transition and the object. TableC.1 shows the CO line fluxes for all lines in the PACS range.
4. Analysis
Emission lines observed in the PACS wavelength range are usedto calculate the contribution of di ff erent species to the total linecooling from high-mass protostars. Our goal is to compare thecooling of warm gas by molecules and atoms with the coolingby dust and connect them with the evolutionary stages of theobjects. Relative contributions to the cooling between di ff erentmolecules are also determined, which can be an indicator of thephysical processes in the environments of young protostars (e.g.Nisini et al. 2002; Karska et al. 2013).We define the total far-IR line cooling ( L FIRL ) as the sumof all emission line luminosities from the fine-structure [O i ]lines (at 63 and 145 µ m) and the detected molecules, follow- ing Nisini et al. (2002) and Karska et al. (2013). [C ii ], the mostimportant line coolant of di ff use interstellar gas, is also expectedto be a significant cooling agent in high-mass star forming re-gions and extragalactic sources. It is not explicitly included inour analysis, however, because the calculated fluxes are stronglya ff ected by o ff -source emission and often saturated (see alsoSection 2). The emitted [C ii ] luminosity is shown below for onlytwo sources, where reliable fluxes were obtained. Cooling inother ionic lines such as [O iii ], [N ii ], and [N iii ] is also excluded,due to the o ff -position contamination and the fact that thoselines trace a di ff erent physical component than the moleculesand [O i ]. Since the only molecules with emission lines are CO,H O, and OH, the equation for the total far-IR line cooling canbe written as: L FIRL = L OI + L CO + L H O + L OH .Table 2 summarizes our measurements. The amount of cool-ing by dust is described by the bolometric luminosity and equals ∼ -10 L ⊙ for our sources (van der Tak et al. 2013). The to-tal far-IR line cooling ranges from ∼ ∼
40 L ⊙ , several or-ders of magnitude less than the dust cooling. Relative contri-butions of oxygen atoms and molecules to the gas cooling areillustrated in Figure 4. Atomic cooling is the largest for the moreevolved sources in our sample (see Table 1), in particular forNGC7538 IRS1 and AFGL2591, where it is the dominant linecooling channel. For those two sources, additional cooling by[C ii ] is determined and amounts to ∼ L FIRL (Table 2).Typically, atomic cooling accounts for ∼
20 % of the total linefar-IR line cooling. Molecular line cooling is dominated by CO,which is responsible for ∼
15 up to 85% of L FIRL , with a mediancontribution of 74%. H O and OH median contributions to thefar-IR cooling are less than 1%, because many of their transitionsare detected in absorption. Assuming that the absorptions arisein the same gas, as found for the case of Orion (Cernicharo et al.2006), they therefore do not contribute to the cooling, but heat-ing of the gas. Still, the contribution of H O to the total FIRcooling increases slightly for more evolved sources, from ∼ Detections of multiple rotational transitions of CO, H O andOH allow us to determine the rotational temperatures of theemitting or absorbing gas using Boltzmann diagrams (e.g.Goldsmith & Langer 1999). For H O and OH, the densities arelikely not high enough to approach a Boltzmann distribution andtherefore the diagrams presented below are less meaningful.Emission line fluxes are used to calculate the number ofemitting molecules, N u , for each molecular transition usingEquation (1), assuming that the lines are optically thin. Here, F λ denotes the flux of the line at wavelength λ , d is the distanceto the source, A is the Einstein coe ffi cient, c is the speed of lightand h is Planck’s constant: N u = π d F λ λ hcA (1)The base 10 logarithm of N u over degeneracy of the upper level g u is shown as a function of the upper level energy, E u , inBoltzmann diagrams (Figures 5 and 6). The rotational temper-ature is calculated in a standard way, from the slope b of thelinear fit to the data in the natural logarithm units, T rot = − / b .Because the size of the emitting region is not resolved by ourinstrument, the calculation of column densities requires addi-tional assumptions and therefore only the total numbers of emit-ting molecules is determined. The formula for the total number
6. Karska et al. 2013: Far-infrared molecular lines from Low- to High-Mass Star Forming Regions observed with Herschel
Table 2.
Far-IR line cooling by molecules and atoms in units of L ⊙ . Source L bol L FIRL L mol L OI L CO L H2O L OH L CII (L ⊙ ) (L ⊙ ) (L ⊙ ) (L ⊙ )G327-0.6 7.3 10 + Notes.
Columns show: (1) bolometric luminosity, L bol , (2) total FIR line cooling, L FIRL ( L mol + L OI ), (3) molecular cooling, L mol , and (4) coolingby oxygen atoms, L OI . Cooling by individual molecules is shown in column (5) CO, (6) H O, and (7) OH. Absence of emission lines that wouldcontribute to the cooling are shown with ”–” (absorption lines are detected for H O and OH). Errors are written in brackets and include 20%calibration error on individual line fluxes. of emitting molecules, N tot , is derived from the expression fortotal column density, N tot = Q · exp( a ), where Q is the partitionfunction for the temperature and a is the y-intercept. Correctingfor a viewing angle, Ω = d /π R , and multiplying by the gasemitting area of radius R , yields: N tot = Q · exp( a ) · d (2)For details, see e.g. Karska et al. (2013) and Green et al. (2013).For absorption lines, column densities, N l , are calcu-lated from line equivalent widths, W λ , using Equation 3 (e.g.Wright et al. 2000). N l = π cW λ g l λ Ag u (3)This relation assumes that the lines are optically thin, the cover-ing factor is unity and the excitation temperature T ex ≪ hc / k λ for all lines.Some of the emission / absorption lines are likely P-Cygniprofiles (Cernicharo et al. 2006) that are not resolved with PACSand which we assume to be pure emission / absorbing lines.The natural logarithm of N l over the degeneracy of the lowerlevel, g l , is shown on the Boltzmann diagrams as a function ofthe lower level energy, E l (Figures 6, 7 and 9). Rotational tem-peratures and total column densities are calculated in the sameway as for the emission lines (see above).In the following sections, the excitation of CO, H O, andOH are discussed separately. Table 3 summarizes the valuesof rotational temperatures, T rot , and total numbers of emittingmolecules or column densities, for those species. Figure 5 shows CO rotational diagrams for all our sources.CO detections, up to J = E u = ∼
370 K (W3IRS5 andG34.26 + T rot , CO ∼ ±
60 K. Thehighest temperatures are seen for objects where high- J CO tran-sitions with E u > The value in the brackets (23) shows the average error of rotationaltemperatures for di ff erent sources, whereas ±
60 is the standard devia-tion of rotational temperatures.
Temperatures of ∼
300 K are attributed to the ‘warm’component in low-mass YSOs (e.g. Goicoechea et al. 2012;Karska et al. 2013; Green et al. 2013), where they are calculatedusing transitions from J up =
14 ( E u =
580 K) to 24 ( E u = E u ∼ J up ≥
25 transitions are attributed to the ‘hot’ component. Sucha turning point is not seen on the diagrams of our high-masssources with detections extending beyond the J = −
23 tran-sition.In NGC7538 IRS1, a possible break is seen around E u ∼ T rot1 ∼ ±
10 K and T rot2 ∼ ±
35 K. The lat-ter temperature is consistent within errors with the ‘warm’ com-ponent seen towards low-mass sources. The colder tempera-ture resembles the ∼ −
100 K ‘cool’ component seen in J up ≤
14 transitions (e.g. Goicoechea et al. 2012; Karska et al.2013; van der Wiel et al. 2013), detected at wavelengths longerthan the PACS range.The absence of the hot component towards all our sourcesis not significant according to the calculated upper limits. Inaddition to limited S / N and line-to-continuum ratio, there areother e ff ects that may prevent the hot component from beingdetected. These include the fact that the continuum becomesmore optically thick at the shorter wavelengths (see also below)and / or a smaller filling factor of the hot component in the PACSbeam compared with low-mass sources. More generally, boththe ‘warm’ and ‘hot’ components could still be part of a singlephysical structure such as proposed in Neufeld (2012).The average logarithm of the number of emitting (warm)CO molecules, log N , is similar for all objects, and equals52.4(0.1) ± N in the range from 51.6 to 53.1are derived. DR21(OH), one of the lowest bolometric luminos-ity sources in our sample, exhibits one of the lowest N (CO) con-tents, whereas the highest CO content is found for W51N-e1, themost luminous source. O Figure 6 shows rotational diagrams of H O calculated for fivesources with H O lines detected both in emission and in absorp-tion. Because the emitting region is not resolved by our observa-tions, the diagrams calculated using the emission and absorption
7. Karska et al. 2013: Far-infrared molecular lines from Low- to High-Mass Star Forming Regions observed with Herschel
Fig. 5.
Rotational diagrams of CO for all objects in our sample. The base 10 logarithm of the number of emitting molecules from alevel u , N u , divided by the degeneracy of the level, g u , is shown as a function of energy of the upper level in kelvins, E up . Detectionsare shown as filled circles, whereas three sigma upper limits are shown as empty circles. Blue lines show linear fits to the data andthe corresponding rotational temperatures. The vertical red line in the G34.26 + O diagrams for thethree sources where all H O lines are seen in absorption. The rotational diagrams show a substantial scatter, largerthan the errors of individual data points, caused by large opac-
8. Karska et al. 2013: Far-infrared molecular lines from Low- to High-Mass Star Forming Regions observed with Herschel
Fig. 6.
Rotational diagrams of H O calculated using emission lines (left column) and absorption lines (right column), respectively.For emission line diagrams, the logarithm with base 10 of total number of molecules in a level u , N , divided by the degeneracy ofthe level, g u , is shown as a function of energy of the upper level in Kelvins, E low . For absorptions lines, the natural logarithm of thecolumn density in a level l , N l , divided by the degeneracy of the level, g l , is shown in y-axis. A one component linear fit is shownwith the corresponding value of rotational temperature and error of the fit in brackets. A solid line is used for the cases where at least10 lines are detected. In the W3 IRS5 emission panel, para-H O lines are shown in grey and ortho-H O lines in black, respectively.ities, subthermal excitation (see Herczeg et al. 2012 and below) and possible radiation excitation by far-IR dust emission pump-
9. Karska et al. 2013: Far-infrared molecular lines from Low- to High-Mass Star Forming Regions observed with Herschel
Table 3.
CO, H O, and OH rotational excitation
Source Warm CO H O (em.) H O (abs.) R a OH Π / b OH c T rot (K) log N T rot (K) log N T rot (K) ln N low (arc sec) T rot (K) T rot (K)G327-0.6 295(60) 51.9(0.3) . . . . . . + Notes.
Rotational temperatures of H O and OH within the Π / ladder are calculated using at least 10 and 3 lines / doublets, respectively, andare shown in boldface. Non-detections are marked with dots. For OH temperatures determined using only 2 transitions, the associated error is notgiven and marked with ”-”. ( a ) Size of the H O emitting region assuming that all H O lines trace the same physical component (see Section 4.2.2.). ( b ) Rotational temperatureof OH calculated using the OH Π / ladder transitions only, see Figure 9 and Section 4.2.3. ( c ) Rotational temperature of OH calculated using alllines detected in absorption.
Fig. 7.
Rotational diagrams of H O calculated using absorptionlines. A single component fit is used to calculate the temperatureshown in the panels.ing. The determination of rotational temperatures is thereforesubject to significant errors when only a limited number of linesis detected. In Table 3 H O temperatures calculated using at least10 lines are indicated in boldface. Ro-vibrational spectra of H Ofrom ISO-SWS towards massive protostars (four of them in com-mon with our sample) show rotational temperatures of H O ashigh as 500 K, in agreement within the errors with our measure-ments (Boonman & van Dishoeck 2003).The largest number of H O lines is detected in the two mostevolved sources – G5.89-0.39 and W3 IRS5. Single componentfits to 12 and 37 water emission lines, respectively, give similarrotational temperatures of ∼
250 K (see Figure 6), with no sys-tematic di ff erences between o -H O and p -H O lines. Rotationaltemperatures determined for the remaining sources, with at least5 detections of H O in emission, are higher, T rot ∼ −
450 K,but are less accurate.
Fig. 8.
Top : Rotational diagram of H O calculated assuming akinetic temperature T = density n = cm − ,H O column density 2 · cm − , and line width ∆ V = − . Lines observed in absorption in NGC6334-I are shown inblue. Fits are done to all lines in the PACS range included inthe LAMDA database (Sch¨oier et al. 2005) (in black) and to theNGC6334-I lines (in blue). Darker shades denote optically thinlines. Bottom : The same as above, but assuming a H O columndensity 10 cm − , such that all lines are optically thin.Rotational temperatures calculated from a single componentfit to the absorption line diagrams are ∼
200 K for all sources.
10. Karska et al. 2013: Far-infrared molecular lines from Low- to High-Mass Star Forming Regions observed with Herschel
Fig. 9.
Rotational diagrams of OH calculated using absorption lines, similar to Figure 7 for H O. The single component fit to allOH transitions detected in absorption is shown with a solid line with the corresponding temperature. A separate single componentfit is done for the OH Π / ladder transitions and drawn in a dashed line. The respective rotational temperatures are tabulated inTable 3.They are in good agreement with the values obtained from theemission diagrams for G5.89-0.39 and W3 IRS5, suggesting thatall H O lines originate in the same physical component (see e.g.Cernicharo et al. 2006). In such a scenario, the column densi- ties should also agree, and the comparison of the total numberof molecules calculated from emission lines and columns deter-mined from the absorption lines yields the size of the emittingregion, R . The radius of the H O emitting area under that con-
11. Karska et al. 2013: Far-infrared molecular lines from Low- to High-Mass Star Forming Regions observed with Herschel dition equals ∼ ff ects of optical depth and subthermal exci-tation on the derived temperatures from the absorption lines,equivalent widths of lines are calculated using the radiativetransfer code Radex (van der Tak et al. 2007) and translated tocolumn densities using Equation (2). The adopted physical con-ditions of T kin = n = cm − are typical ofwarm, shocked region where water is excited (Goicoechea et al.2012). The models were calculated using all H O lines in thePACS range included in the LAMDA database (Sch¨oier et al.2005). The latest available H O collisional coe ffi cients are used(Daniel et al. 2011, and references therein).Figure 8 shows the ‘theoretical’ rotational diagrams calcu-lated for low and high column densities. A single component fitto all lines in the high column density model gives a tempera-ture of ∼
120 K, consistent with subthermal excitation, and atotal column density of 1.2 · cm − (ln N low ∼ ∼
180 K and a slightly higher column density (1.9 · cm − ). These observed lines are typically highly optically thickwith τ ∼ few tens, up to 100.In the low column density model all lines are optically thin.A fit to all lines gives a temperature of ∼
120 K, similar to thehigh column density model. The level populations are clearlysubthermal ( T rot ≪ T kin ), resulting in the scatter in the diagram.These examples illustrate the di ffi culty in using the inferred ro-tational temperatures to characterize a complex environment ofhigh-mass star forming regions.The continuum opacity at PACS wavelengths is typically ofthe order of a few in the observed sources, as indicated by thesource structures derived by van der Tak et al. (2013), and be-comes higher at shorter wavelengths. This implies that the ab-sorbing H O is on the frontside of the source. Also, any emissionat short wavelengths must originate outside the region where thedust is optically thick.
Figure 9 presents rotational diagrams of OH calculated using ab-sorption lines. A single component is fitted to all detected lines.A separate fit is done for the lines originating in the Π / lad-der, which are mostly collisionally excited. This fit excludes theintra-ladder 79 µ m doublet, connecting with the ground state,that readily gets optically thick (Wampfler et al. 2013). A pos-sible line-of-sight contribution by unrelated foreground is ex-pected in the ground-state 119 µ m line, which is included inthe fit. The resulting rotational temperatures for each source areshown in Table 3, separately for those two fits.The average rotational temperature for the Π / ladder isvery similar for all objects and equals 100(12) ± T rot , OH ∼ ±
5. Discussion: from low to high mass
Several physical components have been proposed as a sourceof far-IR CO emission in isolated low-mass young stel-lar objects: (i) the inner parts of the quiescent envelope,passively heated by a central source (Ceccarelli et al. 1996;Doty & Neufeld 1997); (ii) gas in cavity walls heated by UV
Table 4.
Input parameters for the C O envelope emissionmodel a Object X X in b T ev FWHM O / O c (K) (km s − )G327-0.6 3.8 10 − – – 5.0 387W51N-e1 3.0 10 − – – 4.7 417DR21(OH) 1.3 10 − – – 2.4 531NGC6334-I 0.5 10 − −
35 4.2 437G5.89-0.39 0.1 10 − −
40 4.5 460NGC7538-I1 2.1 10 − −
35 2.1 614
Notes. ( a ) Parameters for the remaining sources will be presented inSan Jos´e-Garc´ıa et al. (in prep.). ( b ) A jump abundance profile is usedto model NGC6334-I, G5.89-0.39 and NGC7538 IRS1. ( c ) The ratio de-pends on the source’s distance from the Galaxy center (Wilson & Rood1994). photons (van Kempen et al. 2009, 2010; Visser et al. 2012); (iii)currently shocked gas along the outflow walls produced by theprotostellar wind-envelope interaction (van Kempen et al. 2010;Visser et al. 2012; Karska et al. 2013).The quiescent envelope of high-mass protostars is warmerand denser than for low-mass YSOs and therefore its contribu-tion to the far-IR CO emission is expected to be larger. In thissection, we determine this contribution for a subsample of oursources using the density and temperature structure of each en-velope obtained by van der Tak et al. (2013).In this work, the continuum emission for all our objects ismodeled using a modified 3D Whitney-Robitaille continuum ra-diative transfer code (Robitaille 2011; Whitney et al. 2013). Forsimplicity, the van der Tak et al. models do not contain any cavityor disk, and assume a spherically symmetric power law densitystructure of the envelope, n ∝ r − p , where p is a free parame-ter. The size and mass of the envelope, and the power law ex-ponent, p , are calculated by best-fit comparison to the spectralenergy distributions and radial emission profiles at 450 and 850 µ m (Shirley et al. 2000). The models solve for the dust tempera-ture as function of radius and assume that the gas temperature isequal to the dust temperature. For more detailed discussion, seevan der Tak et al. (2013).The envelope temperature and density structure fromvan der Tak et al. (2013) is used as input to the 1D radiativetransfer code RATRAN (Hogerheijde & van der Tak 2000) inorder to reproduce simultaneously the strenghts of opticallythin C O lines from J = Oconstant abundance X and the line width, FWHM. For threesources, a ‘jump’ abundance profile structure is needed, de-scribed by the evaporation temperature T ev and inner abun-dance X in . The parameters derived from the fits are sum-marized in Table 4; the C O observations are taken fromSan Jos´e-Garc´ıa et al. (2013), while the modeling details and re-sults for all our objects will be presented in I. San Jos´e-Garc´ıa(in prep.).The parameters from Table 4 are used as inputs for theRATRAN models of CO. The integrated CO line emissionobtained from RATRAN is convolved with the telescope beamand compared with observed line fluxes.Figure 10 compares the envelope model for NGC7538 IRS1with the CO J up =
14 to 22 observations from Herschel / PACS,CO 3-2 (in 14” beam, San Jos´e-Garc´ıa et al. 2013) and CO 7-6 (in 8” beam, Boonman et al. 2003) from the James ClerkMaxwell Telescope. By design, the model fits the line profile
12. Karska et al. 2013: Far-infrared molecular lines from Low- to High-Mass Star Forming Regions observed with Herschel
Table 5.
Far-IR CO emission: observations and envelope models
Object L CO (obs) L CO (env)(L ⊙ ) (L ⊙ ) (%)High-mass YSOs a G327-0.6 1.9 1.6 84W51N-e1 25.8 14.6 56DR21(OH) 1.2 0.7 58NGC6334-I 3.4 2.4 71G5.89-0.39 3.9 1.8 46NGC7538-I1 2.1 1.6 77Low-mass YSOs b NGC1333 I2A 4.1 10 − − − − − − Notes. ( a ) Observed CO luminosities are calculated using detected tran-sitions only, from J = ( b ) Results from Visser et al. (2012). The ob-served CO emission is taken to be the total CO emission originatingfrom all modeled physical components. of C O 9-8 (San Jos´e-Garc´ıa et al. 2013, from Herschel / HIFI)shown in the bottom of Figure 10. The pure envelope modelslightly underproduces the CO 10-9 line, because it does notinclude any broad entrained outflow component ( T ex ∼
70 K,Yıldız et al. 2013). Adding such an outflow to the model (seeMottram et al. 2013) provides an excellent fit to the total lineprofile. For the case of CO, the pure envelope model repro-duces the 3-2 and 7-6 lines within a factor of two, with the dis-crepancy again being due to the missing outflow in the model.This envelope model reproduces the CO integrated intensitiesfor transitions up to J up =
18, but the larger J CO fluxes areunderestimated by a large factor.Comparison of the observed and modeled integrated COline emission for the remaining sources from Table 4 is shownin Figure 11. As found for the case of NGC7538 IRS1, the con-tribution of the quiescent envelope emission can be as high as70-100% of that of the J = J transitions. Only 3-22% of CO J = ∼ − J CO lines. The broadline profiles of high- J ( J ≥
10) CO lines (San Jos´e-Garc´ıa et al.2013) argue in favor of a shock contribution to the far-IR emis-sion in CO. There may also be a contribution from UV-heatingof the outflow cavities by the photons from the protostellaraccretion shocks or produced by high velocity shocks insidethe cavities as found for low-mass YSOs (Visser et al. 2012)but this component is best distinguished by high- J CO lines(van Kempen et al. 2009). Physical models similar to those de-veloped for low-mass sources by Visser et al. (2012), which in-clude the di ff erent physical components, are needed to comparethe relative contribution of the envelope emission, shocks, andUV-heating in the high-mass sources, but this is out of the scopeof this paper. Fig. 10.
Top:
Comparison of integrated line fluxes of CO ob-served by Herschel / PACS and from the ground (black dots witherrorbars) and the predictions of the quiescent envelope pas-sively heated by the luminosity of the source (red crosses) forNGC7538-IRS1.
Bottom:
The same model compared with theJCMT CO 3-2 and Herschel / HIFI CO 10-9 and C O 9-8observed line profiles. Additional model including an outflowcomponent is shown in blue dashed line.
The basic excitation analysis using Boltzmann diagrams inSection 4.2 shows remarkably similar rotational temperatures ofeach molecule for all our high-mass sources, irrespective of theirluminosity or evolutionary stage. The average values of thosetemperatures are: 300 K for CO, 220 K for H O, and 80 K forOH.Figure 12 presents our results in the context of low-and intermediate-mass YSOs studies by (Fich et al. 2010;Herczeg et al. 2012; Goicoechea et al. 2012; Manoj et al. 2013;Wampfler et al. 2013; Karska et al. 2013; Green et al. 2013;Lee et al. 2013). Rotational temperatures from the Water In Starforming regions with Herschel (WISH), the Dust, Ice and Gas inTime (DIGIT), and the Herschel Orion Protostar Survey (HOPS)programs are shown separately. OH rotational temperatures ofNGC1333 I4B, Serpens SMM1, and L1448 are taken from theliterature, whereas temperatures for the two additional low-mass YSOs and four intermediate-mass YSOs are calculatedin Appendix B based on the line fluxes from Wampfler et al.(2013).Rotational temperatures of CO are remarkably similar formost sources in the luminosity range from 10 − to 10 L ⊙ andequal to ∼ −
350 K. For the high-mass sources, this refersto the shocked component, not the quiescent envelope compo-nent discussed in §
13. Karska et al. 2013: Far-infrared molecular lines from Low- to High-Mass Star Forming Regions observed with Herschel
Fig. 11.
Comparison of integrated line fluxes of CO observedby PACS only (black dots with errorbars) and the predictions ofthe quiescent envelope passively heated by the luminosity of thesource (red crosses).ditions of the gas have been proposed: (i) CO is subthermallyexcited in hot ( T kin ≥ K), low-density ( n (H ) ≤ cm − )gas (Neufeld 2012, Manoj et al. 2013); or (ii) CO is close toLTE in warm ( T kin ∼ T rot ) and dense ( n (H ) > n crit ∼ cm − )gas (Karska et al. 2013). The low-density scenario (i) in the caseof even more massive protostars studied in this work is ratherunlikely. Even though no CO lines are detected in our PACSspectra, three of our high-mass protostars were observed in thefundamental v = − CO (Mitchell et al. 1990). The Boltzmann distribution of high- J populations of CO, indicated by a single component on ro-tational diagrams, implies densities above 10 cm − for W33Aand NGC7538-I1 and > cm − for W3 IRS5.Rotational temperatures of H O increase for the more mas-sive and more luminous YSOs from about 120 K to 220 K(Figure 12). The similarity between the temperatures obtainedfrom the absorption and emission lines argues that they arisein the same physical component in high-mass YSOs (see alsoCernicharo et al. 2006). Due to the high critical density, the wa-ter lines are most likely subthermally excited in both low- andhigh-mass YSOs (see above discussion in Section 4.2.2), but inthe denser environment of high-mass protostars, the gas is closerto LTE and therefore the rotational temperatures are higher. Highoptical depths of H O lines drive the rotational temperatures tohigher values, both for the low- and high-mass YSOs. Lines arein emission, when the angular size of the emitting region ( ∆Ω L ) Fig. 12.
Rotational temperatures of CO, H O, and OH for low-to high mass star forming regions. WISH, DIGIT and HOPSteam’s results obtained with PACS are shown in blue, red, andnavy blue, respectively. Dotted lines show the median valuesof the rotational temperature from each database; for the caseof WISH, the median is calculated separately for objects with L bol < L ⊙ and L bol > L ⊙ , except for OH whereintermediate-mass YSOs covering 10 > L bol >
50 L ⊙ are alsoshown separately. The CO and H O excitation of intermediate-mass sources has not yet been surveyed with Herschel.multiplied by the blackbody at excitation temperature is largerthan the continuum flux at the same wavelength, ∆Ω L × B ν ( T ex ) > F cont ,ν (4)and are in absorption in the opposite case.Rotational temperatures of OH show a broad range of valuesfor low-mass YSOs (from about 50 to 150 K), whereas they areremarkably constant for intermediate mass YSOs ( ∼
35 K; seeAppendix D) and high-mass YSOs ( ∼
80 K). The low tempera-tures found towards the intermediate mass YSOs may be a resultof the di ff erent lines detected towards those sources rather thana di ff erent excitation mechanism. Figure 13 shows relations between selected line luminosities ofCO, H O, and [O i ] transitions and the physical parameters ofthe young stellar objects. Our sample of objects is extended to
14. Karska et al. 2013: Far-infrared molecular lines from Low- to High-Mass Star Forming Regions observed with Herschel the low-mass deeply embedded objects studied with PACS inHerczeg et al. (2012); Goicoechea et al. (2012); Wampfler et al.(2013); Karska et al. (2013) and intermediate-mass objects fromFich et al. (2010); Wampfler et al. (2013). This allows us tostudy a broad range of luminosities, from ∼ L ⊙ , andenvelope masses, from 0.1 to 10 M ⊙ .The typical distance to low-mass sources is 200 pc, whereasto high-mass sources – 3 kpc. For this comparison, the full PACSarray maps of low-mass regions are taken ( ∼ ′′ ), which corre-sponds to spatial scales of 10 AU; only a factor of 3 smaller thanthe central spaxel ( ∼ ′′ ) observation of high-mass objects. Onthe other hand, the physical sizes of the low-mass sources aresmaller than those of high-mass sources by a factor that is com-parable to the di ff erence in average distance of low- and high-mass sources, so one could argue that one should compare justthe central spaxels for both cases. Using only the central spaxelfor the low-mass YSOs does not a ff ect the results, however (seeFigure 9 in Karska et al. 2013).The choice of CO, H O, OH, and [O i ] transitions is based onthe number of detections of those lines in both samples and theiremission profiles. The strengths of the correlations are quanti-fied using the Pearson coe ffi cient, r . For the number of sourcesstudied here, the 3 σ correlation corresponds to r ≈ . σ correlation to r ≈ . σ correlations between the se-lected line luminosities and bolometric luminosities as well asenvelope masses. The more luminous the source, the larger is itsluminosity in CO, H O, and [O i ] lines. Similarly, the more mas-sive is the envelope surrounding the growing protostar, the largeris the observed line luminosity in those species. The strength ofthe correlations over such broad luminosity ranges and envelopemasses suggests that the physical processes responsible for theline emission are similar.In the case of low-mass young stellar objects, Karska et al.(2013) linked the CO and H O emission seen with PACS withthe non-dissociative shocks along the outflow walls, most likelyirradiated by the UV photons. The [O i ] emission, on the otherhand, was attributed mainly to the dissociative shocks at thepoint of direct impact of the wind on the dense envelope. In thehigh-mass sources the envelope densities and the strength of ra-diation are higher, but all in all the origin of the emission can besimilar. Figure 14 compares the total far-IR cooling in lines, its molec-ular and atomic contributions, and the cooling by dust for theYSOs in the luminosity range from ∼ L ⊙ .The far-IR line cooling, L FIRL , correlates strongly (5 σ ) withthe bolometric luminosity, L bol , in agreement with studies onlow-mass YSOs (Nisini et al. 2002; Karska et al. 2013). Underthe assumption that L FIRL is proportional to the shock energy, thestrong correlation between L bol and L FIRL has been interpretedby Nisini et al. (2002) as a result of the jet power / velocity beingcorrelated with the escape velocity from the protostellar surfaceor an initial increase of the accretion and ejection rate.The ratio of molecular and atomic line cooling, L mol / L atom , issimilar for YSOs of di ff erent luminosities, although a large scat-ter is present. Cooling in molecules is about 4 times higher thancooling in oxygen atoms. If cooling by [C ii ] was included in theatomic cooling, the L mol / L atom ratio would decrease for the high-mass sources. In low-mass YSOs, the [C ii ] emission accountsfor less than 1% of the total cooling in lines (Goicoechea et al.2012, for Ser SMM1). In the high-mass sources, this contribu- Fig. 13.
Correlations of line emission with bolometric luminos-ity (left column) and envelope mass (right column) from top tobottom: CO 14-13, H O 3 -2 , [O i ] at 145 µ m and OH 163 µ mline luminosities. Low- and intermediate-mass young stellar ob-jects emission is measured over 5 × O line). High-mass YSOs fluxes are measuredin the central position and shown in black. Pearson coe ffi cient r is given for each correlation.tion is expected to be higher due to the carbon ionizing and COdissociating FUV radiation. In Section 4.1 we estimate the [C ii ]cooling in two high-mass sources to 10-25% of L FIRL .The ratio of cooling by gas (molecular and atomic lines)and dust ( L bol ) decreases from 1.3 · − to 6.2 · − from low tohigh-mass YSOs. It reflects the fact that H O, and to a smallerextent OH, contributes less to the line luminosity in the high-mass sources, because many of its lines are detected in absorp-tion. The detection of H O and OH lines in absorption provesthat IR pumping is at least partly responsible for the excita-tion of these molecules and the resulting emission lines (e.g.Goicoechea et al. 2006; Wampfler et al. 2013). In the case wherecollisions play a marginal role, even the detected emission linesof those species do not necessarily cool the gas.The above numbers do not include cooling from moleculesoutside the PACS wavelength range. This contribution can be
15. Karska et al. 2013: Far-infrared molecular lines from Low- to High-Mass Star Forming Regions observed with Herschel
Fig. 14.
From top to bottom: (1) Total far-IR line cooling; (2)Ratio of molecular to atomic cooling; (3) Gas to dust coolingratio, L FIRL / L bol , where L FIRL = L mol + L atom ; as a function ofbolometric luminosity. Low-mass YSOs are shown in red (Class0) and blue (Class I), whereas high-mass ones are in black.significant for CO, where for low-mass sources the low- J linesare found to increase the total CO line luminosity by about 30%(Karska et al. 2013). Thus, the CO contribution to the total gascooling is likely to be even larger than suggested by Table 2.
6. Conclusions
We have characterized the central position
Herschel / PACS spec-tra of 10 high-mass protostars and compared them with the re-sults for low- and intermediate-mass protostars analyzed in asimilar manner. The conclusions are as follows:1. Far-IR gas cooling of high-mass YSOs is dominated by CO(from ∼
15 to 85% of total far-IR line cooling, with mediancontribution of 74%) and to smaller extent by [O i ] (with me-dian value ∼
20 %). H O and OH median contributions tothe far-IR cooling are less than 1%. In contrast for low-massYSOs, the H O, CO, and [O i ] contributions are compara-ble. The e ff ective cooling by H O is reduced because manyfar-IR lines are in absorption. The [O i ] cooling increases formore evolved sources in both mass regimes.2. Rotational diagrams of CO in the PACS range show a single, warm component , corresponding to rotational temperature of ∼
300 K, consistent with low-mass YSOs. Upper limits on high − J CO do not exclude the existence of an additional, hot component in several sources of our sample.3. Emission from the quiescent envelope accounts for ∼ − J CO lines.4. Rotational diagrams of H O are characterized by T rot ∼ ff ects.5. OH rotational diagrams are described by a single rota-tional temperatures of ∼
80 K, consistent with most low-massYSOs, but higher by ∼
45 K than for intermediate-mass ob-jects. Similar to H O, lines are sub-thermally excited.6. Fluxes of the H O 3 − line and the CO 14 −
13 linestrongly correlate with bolometric luminosities and envelopemasses over 6 and 7 orders of magnitude, respectively. Thiscorrelation suggests a common physical mechanism respon-sible for the line excitation, most likely the non-dissociativeshocks based on the studies of low-mass protostars.7. Across the large luminosity range from ∼ L ⊙ , thefar-IR line cooling strongly correlates with the bolometricluminosity, in agreement with studies on low-mass YSOs.The ratio of molecular and atomic line cooling is ∼
4, similarfor all those YSOs.8. Because several H O lines are in absorption, the gas to dustcooling ratio decreases from 1.3 · − to 6.2 · − from lowto high-mass YSOs. Acknowledgements.
Herschel is an ESA space observatory with science in-struments provided by European-led Principal Investigator consortia and withimportant participation from NASA. AK acknowledges support from theChristiane N¨usslein-Volhard-Foundation, the L ′ Or´eal Deutschland and theGerman Commission for UNESCO via the ‘Women in Science’ prize. JRG, LC,and JC thank the Spanish MINECO for funding support from grants AYA2009-07304, AYA2012-32032 and CSD2009-00038 and AYA. JRG is supported by aRamo´on y Cajal research contract. Astrochemistry in Leiden is supported by theNetherlands Research School for Astronomy (NOVA), by a Royal NetherlandsAcademy of Arts and Sciences (KNAW) professor prize, by a Spinoza grant andgrant 614.001.008 from the Netherlands Organisation for Scientific Research(NWO), and by the European Community’s Seventh Framework ProgrammeFP7 / References
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Appendix A: Details of PACS observations
Table A.1 shows the observing log of PACS observations usedin this paper. The observations identifications (OBSID), obser-vation day (OD), date of observation, total integration time, pri-mary wavelength ranges, and pointed coordinates (RA, DEC)are listed. All spectra were obtained in Pointed / Chop-Nod ob-serving mode. Additional remarks are given for several sources. G327-0.6 and W33A observations were mispointed. NGC6334-I, W3IRS5, and NGC7538-IRS1 spectra were partly satu-rated and re-observed (re-obs). Two observations of W51N-e1,G34.26, G5.89, and AFGL2591 were done using di ff erent point-ing. Appendix B: Continuum measurements
Table B.1 shows the continuum fluxes for all our sources mea-sured using the full PACS array. The fluxes were used inthe spectral energy distributions presented by van der Tak et al.(2013).
Appendix C: Tables with fluxes and additionalfigures
Table C.1 shows line fluxes and 3 σ upper limits of CO linestoward all our objects in units of 10 − W cm − . For details, seethe table caption.Figure C.1 show blow-ups of selected spectral regions of W3IRS5 with high- J CO, H O, and OH lines. Figures C.2 and C.3show blow-ups of selected CO and OH transitions towards allsources.
17. Karska et al. 2013: Far-infrared molecular lines from Low- to High-Mass Star Forming Regions observed with Herschel
Table A.1.
Log of PACS observations
Source OBSID OD Date Total time Wavelength ranges RA DEC Remarks(s) ( µ m) ( h m s ) ( o ′ ′′ )G327-0.6 1342216201 659 2011-03-04 6290 102-120 15 53 08.8 -54 37 01.0 mispointed1342216202 659 2011-03-04 4403 55-73 15 53 08.8 -54 37 01.0 mispointedW51N-e1 1342193697 327 2010-04-06 4589 55-73 19 23 43.7 +
14 30 28.8 di ff . point.1342193698 327 2010-04-06 4969 102-120 19 23 43.7 +
14 30 25.3 di ff . point.DR21(OH) 1342209400 551 2010-11-15 4401 55-73 20 39 00.8 +
42 22 48.01342209401 551 2010-11-16 6280 102-120 20 39 00.8 +
42 22 48.0W33A 1342239713 1018 2012-02-25 4403 55-73 18 14 39.1 -17 52 07.0 mispointed1342239714 1018 2012-02-25 3763 102-120 18 14 39.1 -17 52 07.0 mispointed1342239715 1018 2012-02-25 2548 174-210 18 14 39.1 -17 52 07.0 mispointedG34.26 + +
01 14 58.1 di ff . point.1342209734 542 2010-11-07 4969 102-120 18 53 18.7 +
01 15 01.5 di ff . point.NGC6334-I 1342239385 1013 2012-02-21 4403 55-73 17 20 53.3 -35 47 00.0 saturated1342239386 1013 2012-02-21 3763 102-120 17 20 53.3 -35 47 00.0 saturated1342239387 1013 2012-02-21 2548 174-210 17 20 53.3 -35 47 00.0 saturated1342252275 1240 2012-10-05 3771 102-120 17 20 53.3 -35 46 57.2 re-obsNGC7538-I1 1342211544 589 2010-12-24 6290 102-120 23 13 45.3 +
61 28 10.0 saturated1342211545 589 2010-12-24 4403 55-73 23 13 45.3 +
61 28 10.0 saturated1342258102 1329 2013-01-02 3771 102-120 23 13 45.2 +
61 28 10.4 re-obsAFGL2591 1342208914 549 2010-11-14 6280 102-120 20 29 24.7 +
40 11 19.0 di ff . point.1342208938 550 2010-11-15 4403 55-73 20 29 24.9 +
40 11 21.0 di ff . point.W3-IRS5 1342191146 286 2010-02-24 6345 102-120 2 25 40.6 +
62 05 51.0 saturated1342191147 286 2010-02-24 4102 55-73 2 25 40.6 +
62 05 51.0 saturated1342229091 860 2011-09-21 4403 55-73 2 25 40.6 +
62 05 51.0 saturated1342229092 860 2011-09-21 4499 102-120 2 25 40.6 +
62 05 51.0 re-obs1342229093 860 2011-09-21 2249 55-73 2 25 40.6 +
62 05 51.0 re-obsG5.89-0.39 1342217940 691 2011-04-05 4969 102-120 18 00 30.5 -24 04 00.4 di ff . point.1342217941 691 2011-04-05 4589 55-73 18 00 30.5 -24 04 04.4 di ff . point. Table B.1.
Full-array continuum measurements in 10 Jy λ ( µ m) Continuum (10 Jy)G327-0.6 a W51N-e1 DR21(OH) W33A G34.26 NGC6334I NGC7538I1 AFGL2591 W3IRS5 G5.8956.8 2.2 11.0 1.7 1.9 7.9 16.1 8.6 5.5 23.0 17.759.6 2.6 12.3 2.0 2.1 8.2 17.2 8.9 5.7 23.8 18.862.7 3.0 13.3 2.5 2.2 9.2 18.2 9.3 5.8 24.0 19.763.2 3.0 13.6 2.6 2.3 9.4 18.3 9.3 5.8 24.2 19.966.1 3.5 14.3 3.0 2.5 10.6 19.3 9.6 6.0 24.8 20.269.3 3.8 15.2 3.4 2.5 10.9 19.8 9.6 6.0 24.4 20.772.8 4.6 17.8 4.1 3.0 13.2 20.5 9.7 5.9 24.8 15.776.0 4.9 18.6 4.5 2.9 14.0 20.9 9.6 5.6 24.4 15.879.2 5.3 18.9 4.8 3.0 14.4 20.9 9.5 5.6 23.3 15.781.8 5.6 19.2 5.1 3.2 14.8 21.1 9.6 5.7 23.1 15.786.0 6.1 19.8 5.5 3.2 15.6 21.6 9.6 5.5 22.2 15.790.0 > > > > > Notes.
The calibration uncertainty of 20% of the flux should be included for comparisons with other modes of observations or instruments. ( a ) Onespaxel at N-W corner of the PACS map is saturated around 100 µ m region due to the strong continuum; therefore the tabulated values are the lowerlimits to the total continuum flux from the whole map.18 . K a r s k ae t a l . : F a r- i n fr a r e d m o l ec u l a r li n e s fr o m L o w - t o H i gh - M a ss S t a r F o r m i ng R e g i on s ob s e r v e d w it h H e r s c h e l Table C.1.
CO line fluxes in 10 − W cm − Trans. λ lab ( µ m) Central position flux (10 − W cm − )G327-0.6 a W51N-e1 DR21(OH) W33A a G34.26 + < < < < < .
05 5.18(1.40) < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < < Notes.
The uncertainties are 1 σ measured in the continuum on both sides of each line; calibration uncertainty of 20% of the flux should be included for comparisons with other modes of observationsor instruments. 3 σ upper limits calculated using wavelength dependent values of full-width high maximum for a point source observed with PACS are listed for non-detections. CO 23-22 and CO31-30 fluxes are not listed due to severe blending with the H O 4 -3 line at 113.537 µ m and the OH , - , line at 84.4 µ m, which are often in absorption. CO 25-24, CO 26-25 and CO 37-36transitions are located in the regions of overlapping orders, where the flux calibration is unreliable. ( a ) A mispointed observation. The fluxes are calculated from a sum of two spaxel closest to thetrue source position, see Section 2. . Karska et al. 2013: Far-infrared molecular lines from Low- to High-Mass Star Forming Regions observed with Herschel Table D.1.
OH rotational excitation and number of emittingmolecules N u based on emission lines for low- and intermediate-mass sources Source T rot (K) log N Low-mass YSOsNGC1333 I4A 270(700) 52.4(0.9)L1527
Notes.
Rotational diagrams are shown in Figures D.1 and D.2. Objectswith at least 3 detected doublets in Wampfler et al. (2013) are presented.Rotational temperatures of OH calculated with error less than 100 K areshown in boldface.
Appendix D: OH in low and intermediate masssources
Figures D.1 and D.2 show rotational diagrams of OH for low-and intermediate-mass young stellar objects based on the fluxespresented in Wampfler et al. (2013). Only the sources with atleast 3 detected doublets in emission (out of 4 targeted in total)are shown in diagrams. Rotational temperatures and total num-bers of emitting molecules are summarized in Table D.1. In caseof low-mass YSOs, a single component fit is usually not a goodapproximation (with the exception of L1527 and IRAS15398).In the intermediate-mass YSOs, on the other hand, such approx-imation holds and results in a very similar rotational tempera-tures of OH T rot ∼
35 K for all sources except NGC7129 FIRS2.
Fig. D.2.
OH rotational diagrams (from emission lines) forintermediate-mass young stellar objects (fluxes from Wampfleret al. 2013).
20. Karska et al. 2013: Far-infrared molecular lines from Low- to High-Mass Star Forming Regions observed with Herschel
Fig. C.1.
Close-ups of several of the H O, CO and OH lines in W3IRS5 are shown in Figure 1. The rest wavelength of each line isindicated by dashed lines: blue for H O, red for CO and light blue for OH. Identifications of the undetected lines in the presentedspectral regions are shown in brackets.
21. Karska et al. 2013: Far-infrared molecular lines from Low- to High-Mass Star Forming Regions observed with Herschel
Fig. C.2.
Close-ups of several transitions of CO lines in the PACS wavelength range towards all sources. The spectra are continuumsubtracted and shifted vertically for better visualization.
22. Karska et al. 2013: Far-infrared molecular lines from Low- to High-Mass Star Forming Regions observed with Herschel
Fig. C.3.
Normalized spectra of OH doublets for all our sources at central position. Doublets at 71 and 98 µ m are excluded becauseof poor calibration of those spectral regions observed with PACS. OH doublet at 84.4 µ m is a blend with the CO 31-30 line, whereasOH at 65.13 µ m can be a ff ected by H O 6 -5 at 65.17 µ m. Fig. D.1.