Water in star-forming regions (WISH): Physics and chemistry from clouds to disks as probed by Herschel spectroscopy
AAstronomy & Astrophysics manuscript no. wish_final_Jan2021 © ESO 2021February 5, 2021
Water in star-forming regions: Physics and chemistry from cloudsto disks as probed by
Herschel spectroscopy
E.F. van Dishoeck , , L.E. Kristensen , J.C. Mottram , A.O. Benz , E.A. Bergin , P. Caselli , F. Herpin , M.R.Hogerheijde , , D. Johnstone , , R. Liseau , B. Nisini , M. Tafalla , F.F.S. van der Tak , , F. Wyrowski , A.Baudry , M. Benedettini , P. Bjerkeli , G.A. Blake , J. Braine , S. Bruderer , , S. Cabrit , J. Cernicharo , Y.Choi , , A. Coutens , Th. de Graauw , , C. Dominik , D. Fedele , M. Fich , A. Fuente , K. Furuya , J.R.Goicoechea , D. Harsono , F.P. Helmich , , G.J. Herczeg , T. Jacq , A. Karska , M. Kaufman , E. Keto , T.Lamberts , B. Larsson , S. Leurini , , D.C. Lis , G. Melnick , D. Neufeld , L. Pagani , M. Persson , R.Shipman , V. Taquet , T.A. van Kempen , C. Walsh , S.F. Wampfler , U. Yıldız , and the WISH team (A ffi liations can be found after the references) Received 1 August 2020; Accepted 14 December 2020
ABSTRACT
Context.
Water is a key molecule in the physics and chemistry of star and planet formation, but it is di ffi cult to observe from Earth. The Herschel
Space Observatory provided unprecedented sensitivity as well as spatial and spectral resolution to study water. The Water In Star-forming regionswith
Herschel (WISH) key program was designed to observe water in a wide range of environments and provide a legacy data set to address itsphysics and chemistry.
Aims.
The aim of WISH is to determine which physical components are traced by the gas-phase water lines observed with
Herschel and to quantifythe excitation conditions and water abundances in each of these components. This then provides insight into how and where the bulk of the wateris formed in space and how it is transported from clouds to disks, and ultimately comets and planets.
Methods.
Data and results from WISH are summarized together with those from related open time programs. WISH targeted ∼
80 sources alongthe two axes of luminosity and evolutionary stage: from low- to high-mass protostars (luminosities from < > L (cid:12) ) and from pre-stellarcores to protoplanetary disks. Lines of H O and its isotopologs, HDO, OH, CO, and [O I], were observed with the HIFI and PACS instruments,complemented by other chemically-related molecules that are probes of ultraviolet, X-ray, or grain chemistry. The analysis consists of couplingthe physical structure of the sources with simple chemical networks and using non-LTE radiative transfer calculations to directly compare modelsand observations.
Results.
Most of the far-infrared water emission observed with
Herschel in star-forming regions originates from warm outflowing and shockedgas at a high density and temperature ( > cm − , 300–1000 K, v ∼
25 km s − ), heated by kinetic energy dissipation. This gas is not probed bysingle-dish low- J CO lines, but only by CO lines with J up >
14. The emission is compact, with at least two di ff erent types of velocity componentsseen. Water is a significant, but not dominant, coolant of warm gas in the earliest protostellar stages. The warm gas water abundance is universallylow: orders of magnitude below the H O / H abundance of 4 × − expected if all volatile oxygen is locked in water. In cold pre-stellar cores andouter protostellar envelopes, the water abundance structure is uniquely probed on scales much smaller than the beam through velocity-resolvedline profiles. The inferred gaseous water abundance decreases with depth into the cloud with an enhanced layer at the edge due to photodesorptionof water ice. All of these conclusions hold irrespective of protostellar luminosity. For low-mass protostars, a constant gaseous HDO / H O ratio of ∼ abundance stays constant as a function of position in low-masspre- and protostellar cores. Water abundances in the inner hot cores are high, but with variations from 5 × − to a few × − for low- andhigh-mass sources. Water vapor emission from both young and mature disks is weak. Conclusions.
The main chemical pathways of water at each of the star-formation stages have been identified and quantified. Low warm waterabundances can be explained with shock models that include UV radiation to dissociate water and modify the shock structure. UV fields upto 10 − times the general interstellar radiation field are inferred in the outflow cavity walls on scales of the Herschel beam from varioushydrides. Both high temperature chemistry and ice sputtering contribute to the gaseous water abundance at low velocities, with only gas-phase (re-)formation producing water at high velocities. Combined analyses of water gas and ice show that up to 50% of the oxygen budget may be missing.In cold clouds, an elegant solution is that this apparently missing oxygen is locked up in larger µ m-sized grains that do not contribute to infraredice absorption. The fact that even warm outflows and hot cores do not show H O at full oxygen abundance points to an unidentified refractorycomponent, which is also found in di ff use clouds. The weak water vapor emission from disks indicates that water ice is locked up in larger pebblesearly on in the embedded Class I stage and that these pebbles have settled and drifted inward by the Class II stage. Water is transported from cloudsto disks mostly as ice, with no evidence for strong accretion shocks. Even at abundances that are somewhat lower than expected, many oceansof water are likely present in planet-forming regions. Based on the lessons for galactic protostars, the low- J H O line emission ( E up <
300 K)observed in extragalactic sources is inferred to be predominantly collisionally excited and to originate mostly from compact regions of current starformation activity. Recommendations for future mid- to far-infrared missions are made.
Key words.
Astrochemistry - methods:observational - stars:formation - protoplanetary disks - ISM:abundances - ISM:jets and outflowsArticle number, page 1 of 58 a r X i v : . [ a s t r o - ph . GA ] F e b & A proofs: manuscript no. wish_final_Jan2021
1. Introduction
Of the more than 200 detected interstellar molecules (McGuire2018), water is special because it combines two of the mostabundant elements in the Universe and plays a key role in thephysics and chemistry of star- and planet-forming regions. Onplanets, water is widely acknowledged as essential for poten-tial habitability and the emergence of life (Chyba & Hand 2005;Kaltenegger 2017). This makes the question of how much wateris present in forming planetary systems, and whether this amountdepends on the location and birth environment, highly relevant.Water ice also plays a role in promoting the coagulation of smalldust grains to pebbles, rocks and ultimately planetesimals, thebuilding blocks of planets, by enhancing the mass of solids dueto freeze out. Such icy planetesimals (asteroids, comets), in turn,may have delivered much of the water and organic moleculestrapped in ices to oceans on planets such as Earth that haveformed inside the water iceline (see Morbidelli et al. 2012, 2018;van Dishoeck et al. 2014; Hartmann et al. 2017; Altwegg et al.2019, for reviews).In the early phases of star formation, water vapor is an ex-ceptional tool for studying warm interstellar gas and the physi-cal processes taking place during star formation. This diagnos-tic capability stems from water’s large abundance variations be-tween warm gas, where it is copiously produced (e.g., Draineet al. 1983; Kaufman & Neufeld 1996), and cold gas, where itis mostly frozen out (e.g., Bergin et al. 2002). Thus, water vaporemission can be used as a “beacon” that signals where energy isdeposited into molecular clouds. This happens especially in thedeeply embedded stages when jets and winds from the proto-stars interact with the surroundings, and when the (proto)stellarluminosity heats envelopes and disks. Water vapor lines are alsoparticularly sensitive to small motions inside clouds, such asthose that are due to gravitational collapse or expansion. Wa-ter a is therefore highly complementary to other molecules suchas CO in probing the protostellar environment. Finally, water ac-tively contributes to the energy balance of warm gas as a gascoolant, with its importance likely varying with protostellar evo-lution (Nisini et al. 2002).The quest for understanding the water “trail” from clouds toplanet-forming disks is complicated by the fact that water in theEarth’s atmosphere prevents the direct observation of rotationallines of water gas with ground-based telescopes, except for somehigh excitation (masing) transitions and a few isotopolog lines.Indeed, the first detection of interstellar water was made from theground through the 22 GHz water maser line in Orion (Cheunget al. 1969). Space-based observatories are needed to probe thefull spectrum of water vapor. As a light hydride, water vaporhas its principal rotational transitions at far-infrared rather thanmillimeter wavelengths, for which development of appropriatedetector technology was required.A related complication is the fact that star formation takesplace over many size scales (Shu et al. 1993; Beuther et al.2007): pc-sized clouds can fragment and contract to ∼ < . ∼
150 pc a Unless specified, the term “water” in this article can imply both watervapor and water ice. The term “ice” is used to indicate all volatiles insolid form, which includes – but is not limited to – water ice. distance, these sizes range from many arcmin to < (cid:48) cannot resolve protoplanetary disks at far-infraredwavelengths, in contrast with what the Atacama Large Millime-ter / submillimeter Array (ALMA) now does routinely at longerwavelengths. However, a single-dish telescope can image proto-stellar envelopes and their outflows. Motions range from < . − to >
100 km s − so that very high spectral resolution isneeded to resolve line profiles which can typically only be pro-vided by heterodyne instruments. This high spectral resolutionallows to infer some spatial information for unresolved sourceswith systematic motions, such as for gas in rotating disks.Pioneering infrared and submillimeter space missions haveprovided considerable insight into the water cycle in space. The Infrared Space Observatory (ISO) (Kessler et al. 1996) coveredthe full mid- and far-infrared wavelength range observing bothwater gas and ice, but at modest spectral and spatial resolu-tion and sensitivity. The
Submillimeter Wave Astronomy Satel-lite (SWAS) (Melnick et al. 2000) observed water with hetero-dyne spectral resolution in a 3 . (cid:48) × . (cid:48) beam, but only throughthe ground-state line of ortho-H O and H
O near 557 and 547GHz, as did the Swedish-led
Odin satellite (Nordh et al. 2003).
Odin had a factor of two higher angular (2 (cid:48) circular beam) andspectral (0.5 MHz) resolution than SWAS, with spectra fromthese two missions agreeing very well when degraded to thesame resolution (Larsson & Liseau 2017). Finally, the
SpitzerSpace Telescope (Werner et al. 2004) covered the mid-infraredwavelength range at low spatial and spectral resolution but withmuch higher sensitivity than ISO, thus probing both water iceand highly excited water gas. These missions, combined with the
Kuiper Airborne Observatory and ground-based observations ofwater ice, have demonstrated that water is mostly frozen out asice in cold clouds and that water vapor becomes abundant at hightemperatures such as associated with outflows. Detailed resultsfrom these missions are summarized elsewhere and will not berepeated here (see e.g., Melnick 2009; Hjalmarson et al. 2003;van Dishoeck 2004; van Dishoeck et al. 2013; Pontoppidan et al.2010b).The
Herschel Space Observatory (Pilbratt et al. 2010) cov-ered the 55–672 µ m range and improved on previous spacemissions in all relevant observational parameters. With a 3.5mtelescope, its di ff raction-limited beam ranged from 44 (cid:48)(cid:48) at thelongest to 9 (cid:48)(cid:48) at the shortest wavelengths, increasing the spatialresolution with respect to S WAS at the longest wavelengths bya factor of 5 and making it a much better match to the sizesof protostellar sources. Its Heterodyne Instrument for the Far-Infrared (HIFI) (de Graauw et al. 2010) provided spectra at veryhigh spectral resolution ( R = λ/ ∆ λ up to 10 ) at a single po-sition over the 480–1250 GHz (600–240 µ m) and 1410–1910GHz (210–157 µ m) ranges, thus covering many water lines andfully resolving their profiles. The Photodetector Array Cameraand Spectrometer (PACS, Poglitsch et al. 2010) obtained spectrawith moderate resolving power ( R = (1 − × ) in the 55-190 µ m range at each pixel of a 5 × ffi cientinstantaneous imaging of water lines and full far-infrared spec-tral scans. The sensitivity of both HIFI and PACS was a factor of > Herschel eminently suited to studyboth hot and cold water in space. The Spectral and Photomet-ric Imaging Receiver (SPIRE) (Gri ffi n et al. 2010) also coveredsome of the lower-frequency water and CO lines at 194–318 and294-672 µ m but at much lower spectral resolution than HIFI andwas generally not used as part of WISH. Article number, page 2 of 58.F. van Dishoeck et al.: WISH: physics and chemistry from cloud to disks probed by
Herschel spectroscopy
Para − K C E up ( K ) HIFI (GHz)PACS ( µ m)PACS + HIFIGround (GHz) Ortho − K C
557 1153 1097108 175 114791361670 1791661175 621 22 J K A , K C Fig. 1.
Water rotational energy levels within the ground vibrational level. Transitions targeted with HIFI are indicated in red with frequencies inGHz, those with PACS in blue with wavelengths in µ m. Dashed lines indicate water transitions detected in PACS and HIFI spectral scans. The179 µ m line observed with both PACS and HIFI is indicated in green. Purple lines indicate water (isotopolog) lines that can be observed from theground with frequencies in GHz; some of them (e.g., 621 GHz) have also been observed with HIFI. The 22 GHz 6 − maser line observed atcm wavelengths is indicated. The goal of the Water In Star-forming regions with
Herschel (WISH) guaranteed time Key Program (KP) was to use watervapor as a physical and chemical diagnostic and follow the wa-ter abundance throughout the di ff erent phases of star- and planetformation. A sample of about 80 young stellar objects (YSOs)was targeted covering two axes: mass and time (or equivalently,evolutionary stage). The WISH sources have luminosities from < > L (cid:12) , thus including low (LM), intermediate (IM)and high mass (HM) protostars. They also cover a large rangeof evolutionary stages, from pre-stellar cores prior to collapse,through the embedded protostellar stages, to the revealed phasewhen the envelope has dissipated and – for the case of low-masssources – the pre-main sequence stars become optically visiblebut are still surrounded by gas-rich disks. The full list of sourcesis included as Table 1 in van Dishoeck et al. (2011), which alsodescribes the WISH program and highlights initial results. Thegas-poor debris disk stage was not covered in WISH but was tar-geted spectroscopically by other programs such as GASPS (Dentet al. 2013).The WISH program was organized along the mass and timeaxes (see Table A.1 in the Appendix) with slightly di ff erent ob-serving strategies for each of the subprograms. For all embeddedprotostars, a comprehensive set of water observations has beencarried out with HIFI and PACS (Fig. 1) with a mix of low- andhigh-excitation lines of H O, H
O, H
O, and the chemicallyrelated species O and OH, targeted. In addition, a number ofhigh- J CO and its isotopologs lines are included in the spectralsettings, as are some hydrides and C + that are diagnostic of thepresence of X-rays or UV that are emitted either by the sourceitself or at the star-disk interface (see Table 2 in van Dishoecket al. 2011). Thus, the WISH data also cover the major coolantsof the gas. For the cold pre-stellar cores and protoplanetary disks,at most a few water lines were observed but with very long inte-grations. Later in the Herschel mission, when additional guaran-teed time became available, some HDO lines were added, sincethe HDO / H O ratio is a particularly powerful tracer of the waterchemistry and its history. The 1 − line of HDO at 893 GHzis a particularly powerful probe of cold HDO gas. Since the WISH sample was limited in the number of sourcesin each luminosity or evolutionary bin, several follow-up opentime programs have been pursued to validate the conclusions.For example, the William Herschel Line Legacy survey (WILL)doubled the number of low-mass protostars in a much more un-biased way (Mottram et al. 2017). A program targeting all pro-tostars in the Cygnus star-forming cloud added greatly to theintermediate- and high-mass samples (San José-García 2015).Very deep integrations on a number of protoplanetary disks al-lowed more general conclusions for those sources to be drawn(Du et al. 2017). Similar extensions hold for many of the WISHsubprograms, the most important of which are summarized inTable A.1. The collection of WISH and these programs is indi-cated here as WISH + .Complementary guaranteed and open time Herschel pro-grams have provided information on water in star-forming re-gions as well, most notably the HIFI spectral surveys as part ofthe CHEmical Survey of Star-forming regions (CHESS; Cecca-relli et al. 2010) and Herschel observations of EXtraOrdinarySources (HEXOS; Bergin et al. 2010b) KPs. The Dust, Ice andGas in Time (DIGIT; Green et al. 2013) and Herschel Orion Pro-tostar Survey (HOPS; Manoj et al. 2013) KPs complementedWISH by carrying out full PACS spectral scans for a largersample of low-mass embedded YSOs. The PRobing InterStel-lar Molecules with Absorption line Studies (PRISMAS) targetedthe chemistry of water and other hydrides in the di ff use interstel-lar gas (Gerin et al. 2016). At the other end of the water trail, theWater and related chemistry in the Solar System (HSSO; Har-togh et al. 2009) probed water and its deuteration in a varietyof solar system objects, including comets. Some of their resultswill be put in the context of the WISH + program.The main aims of the WISH program are to determine:(i) How and where water is formed in space.(ii) Which physical components of a star-forming cloud watertraces, and what the water abundances are in each of thesecomponents. Also, what fraction of the total oxygen reser-voir is accounted for by water, and what the importance is ofwater vapor as a coolant of warm gas. Article number, page 3 of 58 & A proofs: manuscript no. wish_final_Jan2021 (iii) What the water trail is from clouds to disks, and ultimatelyto comets and planets.Overall, WISH + has published nearly 90 papersto date which are summarized on the WISH website . This paper synthesizesthese WISH + results, together with related results from theabove mentioned Herschel programs. It focuses on questions (i)and (iii) and the latter part of question (ii), including previouslyunpublished data, new models and new analyses. The physicalcomponents traced by water vapor have been addressed in anumber of synthesis papers, most notably San José-Garcíaet al. (2016), Herpin et al. (2016), van der Tak et al. (2019),Kristensen et al. (2017b) and Mottram et al. (2017) for WISH + ,Karska et al. (2018) for WISH + WILL + DIGIT, and Manoj et al.(2016) for HOPS + DIGIT, with the latter paper limited to CO.In terms of question (iii), this paper describes the evolutionof water from cores to disks. The bigger picture of the deliveryof water to planets and Earth itself is summarized in other re-views (e.g., van Dishoeck et al. 2014; Morbidelli et al. 2018).Also, this paper is centered around
Herschel ’s contributions tothese questions. Other facilities provide important complemen-tary information, most notably ground-based infrared telescopeson ices (e.g., Boogert et al. 2015), and millimeter interferome-ters such as ALMA and the Northern Extended Millimeter Array(NOEMA) through high resolution images of quiescent warmwater in protostellar sources (e.g., van der Tak et al. 2006; Pers-son et al. 2014; Harsono et al. 2020), but they will be mentionedonly sparsely.This paper is organized as follows. Section 2 summarizesthe main observing strategy, sample selection and data reduc-tion. Section 3 and Appendix D provide an overview of the mainwater formation and destruction processes, especially those thatare relevant for star-forming regions. Also, the WISH modelingapproach is outlined. Section 4 summarizes the main results of
Herschel
WISH + and related programs on water, especially itsdistribution, line profiles and excitation, and links them to obser-vations of the CO rotational ladder. The inferred characteristicshold universally from low- to high-mass protostars. Section 5describes the surprisingly low water abundance inferred for thedi ff erent types of warm outflowing gas and shocks associatedwith star-forming regions and the implications for shock models.Comparison is made with chemically related molecules such asOH observed with Herschel and shock tracers such as SiO andgrain surface products such as CH OH.Sections 6–9 follow water from pre-stellar cores through pro-tostars to disks for low- and high-mass sources. Section 6 fo-cuses on the cold pre-stellar cores, whereas Section 7 describessimilar results for the outer cold protostellar envelopes using thehigh-resolution HIFI line profiles to infer the water abundanceprofile as a function of position in the envelope. New analyses ofHDO and NH (another grain surface chemistry product) are pre-sented. Section 8 highlights the puzzling low water abundancesin the inner hot core of several protostellar envelopes where allices should have sublimated. Section 9 focuses on recent newresults on water in young and mature disks around pre-mainsequence stars, setting the stage for planet formation. Section10 and Appendix G bring all the information on oxygen reser-voirs in star-forming regions together and compare them withdi ff use clouds and comets to address the puzzling question ofthe missing oxygen. Section 11 concludes with a discussion andoverview of what we have learned from Herschel on the abovewater questions, from low to high-mass protostars, and fromclouds to disks. Section 12 summarizes the main points in bul- let form and looks forward to future observations of warm waterusing ALMA, the
James Webb Space Telescope (JWST) and Ex-tremely Large Telescopes (ELTs) on the ground and in space,as well as surveys of water ice using the SPHEREx mission. Italso makes recommendations for future far-infrared space mis-sion concepts such as the
Space IR telescope for Cosmology andAstrophysics (SPICA) (Roelfsema et al. 2018) and the
OriginsSpace Telescope (Meixner et al. 2019). For deep observations ofcold water vapor, the WISH + Herschel data will remain uniquefor a long time.
2. Observations
Water is an asymmetric rotor with energy levels characterized bythe quantum numbers J K A , K C , grouped into ortho ( K A + K C = odd)and para ( K A + K C = even) ladders because of the nuclear spinstatistics of the two hydrogen atoms. In contrast with CO, itsenergy level structure is highly irregular, resulting in transitionsscattered across the far-infrared and submillimeter wavelengthrange. Figure 1 summarizes the H O lines in the vibrationalground state that were targeted by WISH using HIFI and PACS.Also, water lines serendipitously detected in full spectral scansare indicated. The full list of lines, including isotopologs and re-lated species, is summarized in Table 2 of van Dishoeck et al.(2011). To this list, the HDO 1 − O5 − line at 970.3150 GHz ( E up =
599 K) b for high-masssources.The targeted water vapor lines were chosen to cover bothlow- and high-energy levels as well as a range of opacities inthree groups. First, lines connected with the ground state levelsof o- and p-H O at 557, 1113, and 1670 GHz c are prime diag-nostics of the cold gas. They usually show strong self-absorptionor can even be purely in absorption against the strong continuumprovided by the source itself (Fig. 2).The second group, the ‘medium- J ’ lines originating fromlevels with E up around 100–250 K, probe the warm gas and arenot (or less) a ff ected by absorption. The 2 − line at 988GHz, 3 − at 1097 GHz, and 3 − at 1153 GHz (for LMand IM sources), are particularly useful; the latter two lines orig-inate from the same upper level but have Einstein A coe ffi cientsthat di ff er by a factor of 6.3. Because of high optical depths ofH O, optically thin(ner) isotopic lines are crucial for the inter-pretation. Thus, deep integrations on a number of ground-stateand medium- J H O and H
O lines are included in WISH forembedded sources. They proved to be particularly useful for thehigh-mass protostars. For the low-mass subsample, H
O spe-cific settings were only observed toward Class 0 sources whichhave brighter lines than the Class I sources. No water lines hadbeen seen prior to
Herschel for Class I sources, so longer inte-grations were put on the 1 −
557 GHz transition to ensuredetections.The third group are highly excited lines originating fromlevels >
300 K, which are only populated in high temperaturegas and strong shocks (e.g., higher-lying backbone lines). Mostof these lines are covered by PACS in its individual line scanmode and are thus not velocity resolved. They have E up up to b Energies of levels are indicated in short as E up in K units rather than E up / k B with k B the Boltzmann constant. c Frequencies are abbreviated to values commonly used to designatethese lines in the WISH papers. They may di ff er by 1 GHz from theproperly rounded values.Article number, page 4 of 58.F. van Dishoeck et al.: WISH: physics and chemistry from cloud to disks probed by Herschel spectroscopy − −
50 0 50 100 υ - υ source (km s − ) T M B ( K ) BHR71
EHV EHVBroadOffset Narrow
NGC2071 / 5
OtherBroadNarrowOffset (abs)
W3-IRS5 / 130
Other Broad Foreground absNarrow
Fig. 2.
Water spectra of the 1 −
557 GHz line taken with
Herschel -HIFI (38 (cid:48)(cid:48) beam) toward a low-, intermediate- and high-mass YSOs (topto bottom). The spectra highlight the complexity of water profiles withmultiple kinematic absorption and emission components (see text andTable 1 for the interpretation of these components). Other refers to linesfrom molecules other than H O. ff ects thequality of individual line spectra in full spectral scans as illus-trated in Figure 3.CO lines are included in the HIFI settings up to the J = E up =
304 K), both of the main and minor CO iso-topologs (Yıldız et al. 2013b; San José-García et al. 2013). Asdescribed in § 4, these lines are not yet high enough in energyto probe the same physical components as the mid- J water lines.The important CO J = E up =
752 K) does so and wastargeted in a separate program, COPS-HIFI (Kristensen et al.2017b). PACS covers CO lines starting at J = J = E up = × ∼ (cid:48) × (cid:48) region forall sources. Additional, sometimes larger, maps were made withHIFI and PACS of selected water and CO lines together with thecontinuum for a few selected low-mass (e.g., Nisini et al. 2010;Bjerkeli et al. 2012) and high-mass protostars (e.g., Jacq et al.2016; Leurini et al. 2017, Kwon et al., unpublished).Simple hydrides such as CH + , OH + , and H O + , also havetheir lowest transitions at far-infrared wavelengths that can beobserved with Herschel . Since these molecules are either part OH . . CO 22-21
138 139020 H O -
162 16402040 F l uxd e n s it y ( J y ) CO 16-15
174 175
Wavelength ( µ m) H O 3 -2
179 180050 H O 2 -1 Fig. 3.
Examples of PACS spectra of NGC 1333 IRAS4B with the samespectrum shown twice in each panel: (top) taken in the full spectral scanmode at Nyquist sampling; (bottom) taken in individual line scan mode,which has a factor 3 higher spectral sampling. of the water network, have formation pathways that involve en-ergy barriers or are dissociation products, they provide comple-mentary information on the chemistry, gas temperatures and theirradiation by far-UV or X-rays and were therefore included inWISH. Moreover, HCO + , an ion whose main destroyer is water,is serendipitously covered through its J = O 1 − setting.Even though WISH had a generous guaranteed time alloca-tion of 425 hr, a trade-o ff between number of sources, lines andintegration times per line had to be made. Integration times rangefrom as little as 10–15 minutes per line for embedded protostarswith bright water emission, to 10–20 hr per setting to detect thevery weak water lines from cold pre-stellar cores and protoplan-etary disks. The WISH KP source sample of ∼
80 sources is summarized inTable 1 of van Dishoeck et al. (2011), with updates on sourcedistances and luminosities listed in Table 1 of Kristensen et al.(2012) for low-mass and Table 1 of van der Tak et al. (2019)for high-mass protostars. Spectral energy distributions (SEDs)of most sources, including updated far-infrared fluxes from
Her-schel , are presented in Kristensen et al. (2012); Karska et al.(2013); van der Tak et al. (2013); Green et al. (2013); Mottramet al. (2017).WISH selected its source samples among well knownarchetypal sources, based on a wide array of complementarydata. The program distinguishes low (LM, < L (cid:12) ), interme-diate (IM, 10 − L (cid:12) ) and high-mass (HM, > L (cid:12) ) pro-tostars. Also, for a subset of low-mass protostars, positions inthe outflow, well separated from the source, were targeted in aseparate subprogram. In terms of evolutionary stages, pre-stellarcores, embedded protostars (both Class 0 and Class I, separatedby T bol < and >
70 K, Evans et al. 2009), and pre-main sequencestars with disks (Class II), are included. For high mass protostars,the evolutionary stages range from Infrared Dark Clouds (IRDC)to High-Mass Protostellar Objects (HMPO, both IR “quiet” andIR bright) to UltraCompact H II regions (UC HII), although thelatter two stages have some overlap (Beuther et al. 2007; Motteet al. 2018).The WISH sample favors luminous sources with particu-larly prominent and extended outflows in the case of low-mass
Article number, page 5 of 58 & A proofs: manuscript no. wish_final_Jan2021 sources. The latter was a requirement to be able to target separateon-source and o ff -source positions in the outflow subprogram.As demonstrated by other Herschel programs such as WILL andDIGIT, such large extended outflows are not representative of thegeneral population of low-mass protostars (Mottram et al. 2017;Karska et al. 2018). All WISH conclusions have, however, beenverified by these additional samples.The number of sources per WISH (sub)category ranges froma few for the cold line-poor sources (pre-stellar cores, disks)to 10–20 for warm line-rich objects. The latter samples arelarge enough to allow individual source peculiarities to be dis-tinguished from general trends. The radiation diagnostics sub-program observed hydrides other than H O in a subsample oflow- and high-mass protostars. As noted above, the number ofsources per subcategory has in many cases been more than dou-bled thanks to the additional programs listed in Table A.1. De-tails of the observations and data reduction are summarized inAppendix B.
3. Water chemistry
For the interpretation of the
Herschel data, it is important tocompare observed water abundances with the maximum waterabundance that can be expected based on the available elementsin interstellar gas. The average abundance of elemental oxygenwith respect to total hydrogen nuclei in the interstellar mediumis measured to be [O] = . × − (Przybilla et al. 2008), with anuncertainty of about 20%. This value is close to the current bestestimate of the solar oxygen abundance of 4 . × − (Asplundet al. 2009; Grevesse et al. 2010) and nearly identical to the so-lar system value (Lodders 2010). The notation [X] indicates theabundance of element X in all its forms. Some fraction of thisoxygen, (0 . − . × − (16–24%), is locked up in refrac-tory silicate material in the di ff use interstellar medium (Whittet2010). The abundance of volatile oxygen (that is, the oxygen nottied up in some refractory form) is measured to be 3 . × − in di ff use clouds (Meyer et al. 1998), so this is the maximumamount of oxygen that can cycle between water vapor and ice indense clouds.Counting up all the forms of detected oxygen in di ff useclouds, the sum is less than the overall elemental oxygen abun-dance. This missing fraction increases with density (Jenkins2009). Thus, a fraction of oxygen is postulated to be in someyet unknown refractory form, called Unidentified Depleted Oxy-gen (UDO), which may increase from ∼
20% in di ff use clouds upto 50% in dense star-forming regions (Whittet 2010; Draine &Hensley 2020). For comparison, the abundance of elemental car-bon is 3 × − , with about 50–70% of the carbon thought to belocked up in solid carbonaceous material (Henning & Salama1998). If CO (gas + ice) contains the bulk of the volatile car-bon, its fractional abundance should thus be about 1 × − withrespect to total hydrogen, thus accounting for up to 30% of thevolatile oxygen. Indeed, direct observations of warm CO and H gas provide a maximum value of CO / [H] = . × − (Lacy et al.1994), that is, a little less than half the volatile oxygen. Subtract-ing the amount of oxygen in CO from 3 . × − leaves ∼ × − for other forms of volatile oxygen in dense clouds.Often water abundances are cited with respect to H ratherthan total hydrogen. Thus, the maximum expected abundance ifmost volatile oxygen is driven into water is H O / H = × − ,with H O / CO ≈ (1 . − O abundance takesinto account the fraction of oxygen locked up in CO and sili-
Fig. 4.
Simplified gas-phase and solid-state reaction network leadingto the formation and destruction of H O. s-X denotes species X on theice surface. Three routes to water can be distinguished: (i) ion-moleculegas-phase chemistry which dominates H O formation at low temper-atures and low densities (green); (ii) high-temperature neutral-neutralchemistry which is e ff ective above ∼
250 K when energy barriers canbe overcome (red); and (iii) solid-state chemistry (blue). The full solid-state network is presented in Lamberts et al. (2013). Chemically relatedmolecules discussed in this paper such as OH, O and HCO + are shownas well. cates (at 24%), as well as a minimal 20% fraction of UDO suchas found in di ff use clouds. A more detailed discussion of theoxygen budget can be found in § 10 and Appendix G. Detailed networks of the water chemistry in the gas and on thegrains are presented in van Dishoeck et al. (2013) and Lam-berts et al. (2013), and the three main routes to water are de-scribed in Appendix D: (i) cold ion-molecule chemistry; (ii)high-temperature gas-phase chemistry; and (iii) ice chemistry.In addition, the link between the water network and that formingCO , another potentially significant oxygen carrier, is describedthere. The main reactions are illustrated in Fig. 4For coupling chemistry with detailed physical or hydrody-namical models, often smaller chemical networks are preferredto make such calculations practical and computationally feasiblefor large grids of models. Moreover, these simple models oftenallow the key physical and chemical processes to be identified.Two independent simple water chemistry models have been de-veloped within WISH, that by Keto et al. (2014) for cold pre-stellar cores and the SWaN (Simple Water Network) model forprotostellar envelopes by Schmalzl et al. (2014). They includeonly a few of the reactions shown in Fig. 4: (a) cycling of waterfrom gas to ice through freeze-out and UV photodesorption; (b)photodissociation of gaseous water to OH and atomic O, and (c)grain-surface formation of water ice from atomic O.The Keto et al. (2014) model explicitly includes gaseous OHas an intermediate channel, as well as the high temperature O + H and OH + H reactions, although the latter reactions neverbecome significant in cold cores. The SWaN model ignores allOH reactions, but explicitly includes thermal desorption of waterin the hot core region above 100 K.Both simplified models have been extensively tested againstthe full water chemistry models of Hollenbach et al. (2009) (see Article number, page 6 of 58.F. van Dishoeck et al.: WISH: physics and chemistry from cloud to disks probed by
Herschel spectroscopy
Fig. 3 and 4 of Keto et al. 2014) and those of Visser et al. (2011),Walsh et al. (2012) and Albertsson et al. (2014) (see Fig. B.2 ofSchmalzl et al. 2014). While di ff erences of factors of a few inabsolute abundances can readily occur, the overall profiles aresimilar and robustly seen in all models as long as the same phys-ical and chemical parameters are adopted.Multilayer ice models such as developed by Taquet et al.(2014) and Furuya et al. (2016) predict similar abundance pro-files of water as the two-phase simple SWaN models do, but dif-ferences in absolute water abundances can be introduced becausemultilayer interstellar ices are inhomogenous.To compare models with observations, not just the waterchemistry but also the water excitation and radiative transferneed to be treated correctly. The various methods and modelingapproaches adopted by WISH are described in § 4.4 and Ap-pendix F. Deuterated water, HDO and D O, is formed through the sameprocesses as shown in Fig. 4 but there are a number of chemicalreactions that can enhance the HDO / H O and D O / H O ratiosby orders of magnitude compared with the overall [D] / [H] ratioof 2 . × − found in the local interstellar medium (Prodanovi´cet al. 2010). Details can be found in Taquet et al. (2014) and Fu-ruya et al. (2016), and in Ceccarelli et al. (2014) for deuteriumfractionation in general (see also § 7.3). The processes are il-lustrated in Fig. 5. The most e ff ective water fractionation occurson grains, due to the fact that the relative number of D atomslanding from the gas on the grain is enhanced compared withthat of H atoms. In other words, atomic D / H in cold gas is muchhigher than the overall [D] / [H] ratio (Tielens 1983; Roberts et al.2003). This high D / H ratio landing on grains naturally leads toenhanced formation of OD, HDO and D O ice according to thegrain-surface formation routes.The high atomic D / H ratio in the gas, in turn, arises from thewell-known fractionation reactions initiated by the H + + HD → H D + + H reaction, which is exoergic by ∼
230 K and is thusvery e ff ective at low temperatures ≤
25 K (e.g., Watson 1976;Aikawa & Herbst 1999; Stark et al. 1999; Sipilä et al. 2015). TheH D + abundance is further enhanced when the ortho-H abun-dance drops (preventing the back reaction) and when the mainH + and H D + destroyer, CO, freezes out on the grains (Paganiet al. 1992, 2009; Roberts et al. 2003; Sipilä et al. 2010). The lat-ter processes become more important as the cloud evolves froma lower density to a higher density phase (Dartois et al. 2003;Pagani et al. 2013; Brünken et al. 2014; Furuya et al. 2015). Dis-sociative recombination of H D + and other ions such as HD + ,D + and DCO + with electrons then produces enhanced atomicD which gets incorporated in the ices (Fig. 5). As a result, theouter ice layers, which are produced when the cloud is denserand colder, have higher HDO / H O ratios — by orders of mag-nitude — than the inner layers and the bulk of the ice. The en-hanced H D + also leads to enhanced H DO + and thus HDO incold gas, which may play a role at the lower density edge of thecloud.Another characteristic of this layered ice chemistry is thatthe D O / HDO ratio is much higher than the HDO / H O ratio (Fu-ruya et al. 2016). Moreover, the deuteration of organic moleculesformed through hydrogenation of CO ice in the later cloud stagesis generally much higher than that of water (Cazaux et al. 2011;Taquet et al. 2014; Furuya et al. 2016).The desorption processes – photodesorption at low ice tem-peratures (Öberg et al. 2009; Arasa et al. 2015; Cruz-Diaz et al. s-D O H O, OH, H O CR, H IceLow-T H D + H DCO+HD D D s-Ds-HDO DesorptionFreeze-out D, OD
COCOe eeH H HDHD H HDOD O H O H D + , HD EvolutionGrainH O-dominated iceLow HDOHigh HDO, D OCO-dominated ice
Fig. 5.
Simplified reaction network illustrating the important reactionsin the deuteration of water and other molecules. The left-hand side il-lustrates the cold gas-phase chemistry, leading to high fractionation ofgaseous H D + and atomic D, and ultimately gaseous HDO and D O.The right-hand side illustrates how this enhanced D ends up on the iceand leads to enhanced solid HDO and D O. The bottom cartoon showsthe di ff erent ice layers on a grain: the H O-dominated layer formedearly in the cloud evolution with low HDO / H O < − , and the CO-dominated layer formed at higher densities which is rich in deuteratedwater, with D O / HDO >> HDO / H O (Furuya et al. 2015). ff ect on the deuterium fractionation. Inother words, the gaseous HDO / H O and D O / H O ratios shouldreflect the ice ratios following desorption if no subsequent gas-phase reactions are involved. It is important to note that pho-todesorption is only e ff ective in the outer few layers of the ice,whereas thermal desorption removes the bulk ice mantle. Thisselective formation and removal of ice layers turns out to be im-portant in the interpretation of HIFI observations of HDO / H Oratios in cold versus warm gas.In warm gas, the exchange reaction D + OH → H + ODis likely barrierless and can be e ff ective in enhancing OD, espe-cially since the reverse reaction seems to have a barrier of around800 K (Thi et al. 2010). Photodissociation of HDO can also en-hance OD compared with OH by a factor of 2–3, which couldbe a route to further fractionation. Finally, in high temperaturegas in disks the exchange reaction H O + HD ↔ HDO + H isoften included. Similarly, there are reactions inside the ices thatcan both enhance and reduce the water fractionation (Lambertset al. 2015, 2016) but are not considered here (see discussion inFuruya et al. 2016).
4. Herschel water spectra and maps
In this section, we briefly summarize the main features charac-terizing water emission observed with
Herschel , largely basedon WISH + programs. Figure 2 illustrates the complexity of in-dividual water profiles for protostars observed with HIFI overmore than ±
50 km s − whereas Fig. 6 shows the spatial distri-bution imaged with PACS. Figure 7 presents examples of fullspectral scans with PACS with many lines detected. The mainconclusion from the combined analysis of the water maps, the Article number, page 7 of 58 & A proofs: manuscript no. wish_final_Jan2021
L1448 0.05 pc
CB2B3 R2R3R4
L1157
CR1R2 B1B2
NGC1333-IRAS4 C C R2 B2
VLA1623
C GSS30 HH313AHH313BR2B1 C RB
HH211
Fig. 6.
Gallery of water maps in the 2 − µ m line made with Herschel -PACS of the low-mass outflow sources L1448-MM ( d =
235 pc),L1157 (325 pc), NGC 1333 IRAS4 (235 pc), VLA1623 (140 pc) and HH211 (250 pc). The white bars indicate a 0.05 pc scale. The central sourcesare indicated by C, with coordinates and details of data given by Nisini et al. (2010, 2013); Bjerkeli et al. (2012); Santangelo et al. (2014a,b);Dionatos et al. (2018). broad water line profiles, and its excitation, is that most of theobserved gaseous water is universally associated with warm out-flowing and shocked gas of several hundred K. All water linesobserved by
Herschel within WISH show thermal emission, sononmasing, in contrast with the 22 GHz maser often associatedwith star-forming regions. At least two di ff erent types of kine-matic components are involved, with water being a significant(but not necessarily dominant) coolant. In contrast, water emis-sion is not associated with the slower, colder entrained outflowgas traced by low- J CO line wings.Cold quiescent water vapor is also detected but is primar-ily seen in absorption at the protostellar position; only very fewcloud positions show weak narrow (FWHM (cid:46) few km s − ) wa-ter emission lines. These are further discussed in § 7. Recall that Herschel does not observe the bulk of water in cold clouds di-rectly, since this is locked up in ice.
Figure 6 presents maps in the H O 179 µ m line imaged over sev-eral arcmin scales at ∼ (cid:48)(cid:48) resolution with Herschel -PACS fora number of low-mass sources (Nisini et al. 2010, 2013; San-tangelo et al. 2014b; Dionatos et al. 2018). Water vapor emis-sion is clearly associated with the powerful large scale outflowsfrom these sources. Close inspection shows, however, that thewater emission is systematically shifted from that of low- J COemission commonly used to trace outflows (e.g., Tafalla et al.2013): water only spatially coincides with maps of the higher- J CO lines with J up >
14 and with the H mid-IR lines S(1)–S(4) (e.g., Nisini et al. 2013; Bjerkeli et al. 2012; Santangeloet al. 2014a,b; Neufeld et al. 2014) (Fig. 8). Thus, water and high- J CO probe a fundamentally di ff erent gas component thanthat commonly studied with low- J CO lines. This also limits theuse of low- J CO lines to determine, for example, abundances inshocked gas (§ 5).Bright compact water emission is seen at the central proto-star position, in contrast with the thermal emission from othermolecules that are associated with outflows such as SiO andCH OH. This suggests an additional production route of waterin the warm inner protostellar envelope beyond thermal desorp-tion of ices. Along the outflow, the water emission is clumpy,with unresolved individual peaks (at the 13” resolution of the179 µ m line) corresponding to shock spots. In addition, weakermore extended water emission is observed.The [O I] 63 µ m line is commonly associated with jets pow-ering the outflows. Indeed, Herschel -PACS maps show strong [OI] emission along the outflow direction (Fig. 8). In some cases( ∼
10% of sources) even velocity resolved PACS spectra havebeen seen, indicating speeds of more than 90 km s − with respectto that of the source, characteristic of jets (e.g., van Kempen et al.2010; Nisini et al. 2015; Dionatos & Güdel 2017; Karska et al.2018). Comparison of mass flux rates from [O I] and CO sug-gests that this atomic gas is not the dominant driver in the earlieststages: the jets in the Class 0 phase are mostly molecular. How-ever, toward the Class I stage, the jets become primarily atomic,and ultimately ionized in the Class II stage.The HH 211 outflow shown in Fig. 8 is an example of asource where the atomic jet has just enough momentum to powerthe molecular jets and the large scale outflow (Dionatos et al.2018). Maps in the chemically related OH molecule exist foronly a few sources and usually show more compact emissionthan that of water. In the case of HH 211, OH peaks primarily on Article number, page 8 of 58.F. van Dishoeck et al.: WISH: physics and chemistry from cloud to disks probed by
Herschel spectroscopy
50 75 100 125 150 175 200
Wavelength ( µ m) F l uxd e n s it y ( J y ) IRAS16293-2422NGC1333-I4ANGC1333-I4B
COH OOH[OI]
Fig. 7.
Full PACS spectral scans of the low-mass protostars NGC 1333 IRAS4B (Herczeg et al. 2012) and IRAS 1333 IRAS4A, both taken as partof WISH, compared with that of IRAS16293 -2422, taken as part of CHESS. The IRAS4B spectrum, which is extracted at the o ff set and localizedshock position just south of the protostar, is much richer in lines than the other two sources, likely because of less extinction at far-infraredwavelengths in the blue outflow lobe. source and at the bow-shock position. It is important to note thatHH 211 is the only protostellar source which was mapped com-pletely with Herschel -PACS full spectral scans. All other out-flows were only mapped in selected lines.The sources included in Figure 6 represent those with well-known extended outflows; the WISH low-mass protostar sam-ple contains a relatively large fraction of them. However, thisis not representative of the low-mass population. As noted in§ 2.2, the bulk of the low-mass Class 0 and I sources in the fullWISH + WILL + DIGIT sample show little or no extended emis-sion beyond the central PACS 13 (cid:48)(cid:48) spaxel: out of 90 sources only18 show extended water emission in the 47 (cid:48)(cid:48)
PACS footprint(Karska et al. 2013, 2018). This means that the currently activemechanism exciting water in the bulk of the sources is limited toa ∼ < < − line showemission confined to (cid:46) (cid:48)(cid:48) (Jacq et al. 2016; van der Tak et al.2019). Herschel -PACS full spectral scans from the WISH, WILL andDIGIT programs (see Fig. 7 for examples) have directly mea-sured the total far-infrared line cooling, L FIRL of the warm gas(Karska et al. 2014a; Green et al. 2016; Karska et al. 2018). Forthose sources for which only selected individual line scans havebeen observed, corrections for the missing lines have been made.For the low-mass source sample consisting of 90 targets, L FIRL does not change significantly with evolution from the Class 0 tothe Class I stage: median values are 4 . × − L (cid:12) for Class 0and 3 . × − L (cid:12) for Class I sources, respectively (Karska et al.2018).The relative contribution of individual coolants does changewith evolution and between low- and high-mass sources. Fig-ure 9 (left) summarizes the far-infrared line emission of CO,H O, OH and [O I] as a fraction of L FIRL . The LM Class 0 andI numbers are taken from Karska et al. (2018), and this sam-ple is large enough that statistical uncertainties are small. TheClass II values are from Karska et al. (2013) based on the samplefrom Podio et al. (2012). This sample is biased toward Class IIsources with strong optical jets which may be more appropriatefor comparison with the Class 0 and I outflows than disk-onlysources. The Class II sample is too small to conclude whetherdi ff erences between Class I and II are significant. The HM dataare taken from Karska et al. (2014a) and for Orion Peak 1 fromGoicoechea et al. (2015).In the earliest stages of LM sources, CO and H O are thedominant coolants, with [O I] becoming relatively more impor-tant in the Class I and II stages. Interestingly, the fractional cool-ing of CO decreases by a factor of 2 whereas that of H O staysroughly constant from Class 0 to I. This conclusion di ff ers fromthat of Nisini et al. (2002), who found a significant decrease inwater cooling based on earlier ISO data. The absolute [O I] linecooling is similar from Class 0 to Class II, but its fraction in-creases as the jet changes from being mostly molecular to beingprimarily atomic (Nisini et al. 2015).High-mass sources are found to have a smaller fraction ofH O cooling compared with low-mass sources: their far-infraredline emission is dominated by CO. The main reason for thisdi ff erence is that several H O and some OH lines are foundin absorption rather than emission at the high-mass protostellarsource positions, suppressing their contributions to the cooling.
Article number, page 9 of 58 & A proofs: manuscript no. wish_final_Jan2021 −
15 OH 119 µ m[O I ] 63 µ m H O vs CO 2 − Fig. 8.
Map of the water 179 µ m emission in the blue outflow lobe of the HH 211 low-mass Class 0 source (L bol ≈ . L (cid:12) , d =
250 pc) comparedwith that of other lines observed with
Herschel -PACS. The color maps present, from left to right: a high- J CO line, the OH 119 µ m doublet, the[O I] 63 µ m line and the H O 179 µ m line overlaid with a low- J CO ( J = O map is overlaid in light gray contours on the otherpanels. Note the similarity between the H O and CO high- J maps, but not with CO low- J . Data from Tafalla et al. (2013); Dionatos et al. (2018). However, while globally the energy released in the water linesis low, the central emission is higher with the absorption trans-ferring energy to the cooler layers. To assess the e ff ect of thison-source absorption on the cooling budget, the Peak 1 outflowposition in Orion has been added to the figure. This comparisondemonstrates that the conclusion that CO dominates the coolingalso holds o ff source where there is no absorption in the waterlines.Consistent with Nisini et al. (2002), the total far-infrared linecooling is found to strongly correlate with L bol , see Fig. 14 inKarska et al. (2014a) and Fig. 17 in Karska et al. (2018). Theratio of line cooling over dust emission, L FIRL / L bol , decreasesfrom ∼ − for low- and intermediate-mass sources to ∼ − for high-mass sources. Nevertheless, this implies that total far-infrared line luminosity can be used as a direct measure of themechanical luminosity deposited by the jet and wind into theprotostellar surroundings. These data also provide estimates ofthe mass loss rates (Maret et al. 2009; Mottram et al. 2017;Karska et al. 2018). The water lines imaged by PACS reveal the location of warmwater vapor emission, but are spectrally unresolved. In contrast,HIFI spectra reveal the full kinematic structure of the gas butare available mostly at a single position. Figure 2 illustrates thatthe observed water line profiles are universally broad toward alltypes of protostars, from low to high mass, and show dynami-cal components not previously seen from the ground in low- J CO or other molecules (e.g., Kristensen et al. 2010, 2011, 2012,2013; Johnstone et al. 2010; Herpin et al. 2012; van der Tak et al.2013; Mottram et al. 2014; San José-García et al. 2016; Conrad& Fich 2020). They are complex, with emission out to ±
100 kms − , and often have narrow absorption superposed on the ground- CO H O OH O . . . . . . . . L X / L F I R L CO H O OH O
Fig. 9.
Fractions of gas cooling contributed by far-infrared lines of CO,H O, OH and [O I] to the total far-infrared line cooling, as function ofevolutionary stage for low-mass sources (left) and for high mass sources(right). The Class II sample is small and di ff erences with Class I are notstatistically significant. [O I] is seen to become relatively more impor-tant with evolution, whereas H O and OH contribute little for high-masssources. state lines at the source position. Nevertheless, comparing di ff er-ent water transitions, a maximum of only four di ff erent Gaussiancomponents can be distinguished, each of them associated witha di ff erent physical component (Fig. 2, Table 1, see also Fig. C.2in Appendix):1. A broad component (typical FWHM >
15 km s − ) centered atthe source velocity and heated by kinetic energy dissipation. Article number, page 10 of 58.F. van Dishoeck et al.: WISH: physics and chemistry from cloud to disks probed by
Herschel spectroscopy
Table 1.
Physical components seen in water and high- J CO emission lines profiles a Source type Velocity FWHM a CO Excitation Possible origin b Characteristics (km s − ) (K)Pre-stellar narrow < <
10 quiescent, infallLow mass Class 0 broad, 24 300 warm outflow (cavity shocks, disk wind, turbulent mixing layers)medium / o ff set 18 700 hot spot shockmedium / EHV 700 hot EHV bulletsnarrow 3.5 <
100 quiescent, infall, expansionLow mass Class I broad 15 300 warm outflow (cavity shocks, disk wind, turbulent mixing layers)narrow 2.6 <
100 quiescent, infall, expansionIntermediate mass broad 32 300 warm outflow (cavity shocks, disk wind, turbulent mixing layers)narrow 4.6 <
100 quiescent, infall, expansionHigh mass broad 24 300 warm outflow (cavity shocks, disk wind, turbulent mixing layers)narrow 5.6 <
100 quiescent, infall, expansionDisks narrow < a Median width (FWHM) of water lines at the source position, taken from San José-García et al. (2016) and Table 5.4 of SanJosé-García (2015). Values refer to the median of the WISH sample; those including other samples, such as the WILL sample, aresimilar within the uncertainties of a few km s − . b Based on Table 1 of Kristensen et al. (2017b).One scenario is that this water emission is produced in non-dissociative C − type shocks along the outflow cavity, called“cavity shocks” (Mottram et al. 2014). Alternative explana-tions for the observed line profiles and gas heating includeMHD “disk winds” heated by ion-neutral drift (Yvart et al.2016), and “turbulent mixing layers” forming between theprotostellar wide-angle wind and infalling envelope (Lianget al. 2020).2. A medium-broad o ff set component (typical FWHM = − ) slightly blue-shifted from the source velocity. Basedon its velocity o ff set and distinct chemistry, this compo-nent is likely tracing dissociative J − type shocks, called “spotshocks”, at the base of the outflow cavity.3. A medium-broad component with similar FWHM as 2but significantly o ff set from the source velocity, previouslycalled Extremely High Velocity (EHV) gas or “bullets”(Bachiller et al. 1991). Although separately listed here, types2 and 3 form a continuous sequence of velocity o ff set fea-tures in the profiles and are therefore all considered J -typespot shocks (Mottram et al. 2014).4. A narrow component (FWHM < − ) at the source ve-locity, seen primarily in ground-state water lines in absorp-tion and arising in quiescent cloud or envelope material. Theabsorption can also be part of a P-Cygni type profile, con-sisting of strong emission lines with blue-shifted absorptionpointing to an expanding shell of gas. Also, inverse P-Cygniprofiles are seen with red-shifted absorption indicating infall.Components 1–3 involve heating by dissipation of kineticenergy through one or more of the proposed mechanisms. In con-trast, component 4 involves radiative heating, either by the bulkluminosity of the protostar heating the dust in the surroundingenvelope which then transfers its heat to the gas through colli-sions (a ‘passively’ heated envelope), or by UV irradiation of theyoung star-disk system impinging on the gas and dust in outflowcavity walls or disk surfaces and heating it through the photo-electric e ff ect.Roughly 70-80% of the integrated H O emission comes fromthe broad part of the profile (component 1), whereas the medium- broad o ff set parts contribute 20-30% (components 2 +
3) for low-mass sources. For intermediate and high-mass sources, the qui-escent component 4 contributes an increasingly higher fraction,up to 40% for the highest luminosity cases (San José-Garcíaet al. 2016).The first two components are also seen in CO J = E up =
752 K, n crit ≈ × cm − ) profiles observed with HIFIfor low-mass protostars (Kristensen et al. 2017b). Extremelyhigh velocity features are occasionally present as well in the16–15 lines. The fact that they are not yet seen in the CO 10–9 profiles emphasizes the need for high spectral resolution ob-servations at 1.5 THz and above. The broad and medium-broadvelocity components can be associated physically with the twocomponents universally seen in CO rotational diagrams of thesame sources, as the T rot ∼
300 K (broad) and ∼
700 K (medium-broad) component, respectively (Table 1; see also below).At the outflow positions, well o ff set from the LM and IMsource positions by > − on either the blue or red side ofthe flow (Vasta et al. 2012; Tafalla et al. 2013). The wings at theoutflow positions generally follow those seen at the source posi-tion in shape and extent although the level of agreement dependson whether a low or high excitation line is chosen for compari-son (see below and examples in Bjerkeli et al. 2012; Santangeloet al. 2014b). Similar to the source position, the outflow posi-tions show evidence for two physical components, at intermedi-ate and high velocity o ff sets, which have di ff erent distributionsand source sizes (Santangelo et al. 2014b). In addition, EHV gasis seen in a few cases, most notably at various positions in theL1448-MM outflow (Nisini et al. 2013) (see Fig. 2 for BHR71).The EHV component is actually more prominent in water linesthan in the low- J CO lines in which it was originally detected(Bachiller et al. 1990, 1991).
Figure 10 (left) compares the complexity and number of com-ponents seen in velocity-resolved water line profiles with evolu-
Article number, page 11 of 58 & A proofs: manuscript no. wish_final_Jan2021 −
50 0 50 υ - υ source (km s − ) T M B ( d i s t a n ce no r m a li ze d ) Prestellar x 50Class 0Class I x 10Disk x 300 −
50 0 50 υ - υ source (km s − ) IRDC / 35HMPO / 85UCHII / 225
Fig. 10.
Water 1 −
557 GHz spectra at various evolutionary stages,from pre-stellar cores to protostars and disks. Left: Low-mass YSOs,showing from top to bottom the starless core Oph H-MM1, the Class 0source Ser SMM4, the Class I source Elias 29 and the disk TW Hya.Right: high-mass YSOs, showing from top to bottom the starless coreG11-NH , the HMPO DR21(OH), and the UC HII region NGC 7538IRS1. tionary stage across low-mass sources, whereas Fig. 10 (right)does so for high-mass objects.Water emission from quiescent starless or pre-stellar coresprior to star formation is very weak or not detected, both for low-mass (Caselli et al. 2010, 2012) and higher mass cores (Shipmanet al. 2014). If any feature is seen, it is usually in absorptionagainst the weak central continuum (Fig. 10).As soon as the protostar is formed, complex water lineprofiles appear, especially in the earliest most deeply embed-ded stages. The complexity decreases with time: very fewmedium / o ff set components attributed to J − type spot shocks areseen in low-mass Class I sources. The FWHM and Full Width atZero Intensity (FWZI) of the profiles also decrease from Class0 to Class I, suggesting that the outflows become less powerfulas the source evolves (Table 1). These results hold irrespectiveof source sample used to determine the mean or median widthwithin each Class. Figure 10 (right) shows that within the high-mass stages from HMPO to UC HII, no significant di ff erencesin line widths are seen (San José-García et al. 2016; San José-García 2015).For low-mass Class II sources without an envelope or molec-ular outflow, water emission is extremely weak and remains un-detected for the bulk of the protoplanetary disks observed ei-ther with HIFI (Bergin et al. 2010a; Du et al. 2017) or PACS(Fedele et al. 2013; Dent et al. 2013; Rivière-Marichalar et al.2015; Alonso-Martínez et al. 2017).Overall, comparing the di ff erent evolutionary stages inFig. 10, it is clear that water emission ‘turns’ on only when thereis star formation activity in the cloud and it diminishes quicklyonce sources exit the embedded phase.Water vapor emission is generally not detected from molec-ular clouds away from protostars. The only exception are PhotonDominated Regions (PDRs), that is, clouds exposed to enhancedUV radiation where water has been observed to have a singlenarrow emission line component. Examples are the Orion BarPDR (Choi et al. 2014; Putaud et al. 2019), the Orion MolecularRidge (Melnick et al. 2020) and the extended ρ Oph cloud (seeFig. 1 in Bjerkeli et al. 2012).
Nearby intermediate mass sources show very similar features tothose of low-mass Class 0 sources, with all four componentspresent (Johnstone et al. 2010; van Kempen et al. 2016) (seeNGC 2071 in Fig. 2). More distant intermediate mass sourcesand high-mass sources generally show only two components:the broad profile (component 1) and the quiescent envelope ab-sorption and / or emission (component 4) (San José-García et al.2016) (Fig. 10). Either these sources are more evolved, such asthe low-mass Class I sources, or the J − type spot shocks are toobeam diluted due to their larger distance to be detected.The FWHM and FWZI of the water profiles change littlefrom low- to high-mass protostars, suggesting a similar under-lying launching mechanism of the jet. This point is illustratedby the average profiles for three water transitions for each of thetypes of sources presented in San José-García et al. (2016) (seeFig. C.2 in Appendix). Also, the maximum velocity v max reachedby gas seen in CO is similar to that of water and independent ofthe CO line used, even though the line profiles change greatlyfrom low- to high- J and their FWHM are generally smaller thanthat of water (Kristensen et al. 2017b; San José-García et al.2016). At the outflow positions, v max is also similar among trac-ers such as H O, CO and SiO (Santangelo et al. 2012; Bjerkeliet al. 2012).The narrow emission component originating in the quiescentenvelope becomes more prominent for intermediate and high-mass sources due to their higher envelope mass, up to 40% ofthe emission seen in the low-lying H
O lines. This componentis better probed through the excited lines of the less abundant iso-topologs H
O and even H
O: indeed, their line profiles showa narrow(er) emission component with FWHM < − thatlacks the broad outflow wings (Johnstone et al. 2010; Chavarríaet al. 2010; Marseille et al. 2010; Choi et al. 2015; Herpin et al.2016).Interestingly, for high-mass sources, the H O 1 − lineprofiles at 1101 GHz seen in absorption are remarkably similarto the di ff erence between the H O 2 −
988 GHz and 2 −
752 GHz profiles, whose upper energy levels lie at 101 and 137K respectively. This suggests that the narrow H
O absorptionoriginates in the envelope just inside the 100 K radius wherewater sublimates from the grains; it does not arise in the coldenvelope gas (see also § 7 and 8) (Jacq et al. 2016; van der Taket al. 2019).These H
O features are not seen for low-mass sources: inspite of low noise levels, very few low mass sources show detec-tions in the WISH and WILL samples, even in the ground-statelines (Mottram et al. 2014, 2017). Also, when detected (Fig. 11),the H
O lines are in emission and broad, in contrast with nar-row excited CO and C O lines (Kristensen et al. 2010; Yıldızet al. 2013b). The spatially extended broad H
O outflow com-ponent may actually block any narrower emission line arisingcloser to the protostar if it is optically thick. One exception maybe the molecule-rich low-mass protostar IRAS16293-2422, partof the CHESS survey (Coutens et al. 2012), where both H
Oand H
O lines have been detected with widths that are clearlynarrower than those of H O (Fig. 11).Is quiescent warm H
O emission completely absent in low-mass sources? A dedicated deep (up to 5 hr integration per line)open time program has revealed narrow features superposed onthe broad outflow profile in the excited water 3 − linesnear 1097 GHz for all three water isotopologs, but only in a fewsources (Visser et al. 2013) (see § 8 and Fig. C.3 in Appendix). Article number, page 12 of 58.F. van Dishoeck et al.: WISH: physics and chemistry from cloud to disks probed by
Herschel spectroscopy −
50 0 50 υ - υ source (km s − )01234 T M B ( K ) IRAS162931 − H O × O −
50 0 50 υ - υ source (km s − )0 . . . . T M B ( K ) NGC1333 − IRAS4A1 − H O × O −
50 0 50 υ - υ source (km s − )0 . . . . T M B ( K ) NGC1333 − IRAS4B1 − H O × O Fig. 11.
Water 1 −
557 GHz and H
O 548 GHz lines for threelow-mass sources. The H
O line for IRAS16293-2422 is seen to be nar-rower compared with that of the NGC 1333 sources. Data from Coutenset al. (2012) and Mottram et al. (2014).
Taken together, the broad and medium / o ff set components ofthe water line profiles can be used to study its chemistry in warmoutflowing and shocked gas, from low- to high-mass sources.For intermediate and high-mass sources, the narrow H O emis-sion profiles probe the water abundance in the quiescent enve-lope. For low-mass sources, only the narrow part of the H
O(inverse) P-Cygni profiles can be used to determine the quies-cent water abundance, since their H
O profiles are still broadand optically thick, hiding any narrow emission component. Aswill be shown in § 7, these complex (inverse) P-Cygni profilesare actually a remarkably powerful tool to determine the waterabundance profile throughout the cold envelope and cloud.
The integrated intensities and luminosities of the various watervapor lines obtained with HIFI and PACS are well correlatedwith each other at the central protostellar positions, both for low-and high-mass sources (e.g., Karska et al. 2014a; Mottram et al.2014; San José-García et al. 2016). Rotational diagrams usingmostly PACS lines (Fig. 12, middle) indicate median water ex-citation temperatures of ∼
140 K for low-mass sources (Herczeget al. 2012; Goicoechea et al. 2012; Karska et al. 2018) increas-ing up to ∼
250 K for high-mass sources (Karska et al. 2014a),with some spread around the mean. The DIGIT survey finds anaverage value of ∼
190 K (Green et al. 2013). Because water issubthermally excited, these temperatures are lower limits to thekinetic temperature(s) of the emitting gas. Also, because of highbut varying optical depths of the lines, these water excitationtemperatures have a limited meaning.Do these conclusions also hold at positions o ff source? Forlow-mass sources, water vapor maps have been made only in twoof the ground state transitions (557 GHz and 1670 GHz / µ m).The flux ratio of these two lines is again very constant acrossthe entire outflow (Tafalla et al. 2013). Moreover, as Figure 13shows, the protostellar positions do not stand out in this correla-tion, suggesting similar conditions on and o ff source.This good correlation between water lines holds not onlyfor integrated fluxes but even as a function of velocity. HIFIvelocity-resolved observations of the di ff erent water transitionswith E up ranging from 50 to 250 K reveal surprisingly similarline profiles at the low-mass protostellar positions, irrespectiveof excitation and beam size (Mottram et al. 2014). The same isoften true for the intermediate and high-mass protostars for thoselines whose profiles are not a ff ected by absorption (Herpin et al.2016; San José-García et al. 2016) (Fig. C.2 in Appendix). E up / k B (K) l og ( N up / g up ) T rot = 343 K T rot = 902 K CO E up / k B (K) l og ( N up / g up ) T rot = 175 K H O E up / k B (K) l og ( N up / g up ) T rot = 92 K T rot = 853 K OH Fig. 12.
Comparison of the rotational diagram of CO (top), H O (mid-dle) and OH (bottom) for the low-mass YSO NGC 1333 IRAS4B ob-tained from the full PACS spectral scan. The red points indicate thoselines targeted in WISH and WILL for sources for which no full spectralscans have been obtained. The scatter in the water diagram is dominatedby the di ff erence in optical depth of the lines. Di ff erent energy rangesare covered for each of the species. Surveys of the CO rotational ladder for many protostars withPACS have shown that its excitation for J up >
14 can be univer-sally fitted by two rotational temperatures: T rot ∼
300 K (warm)and ∼
700 K (hot) (Fig. 12 top, Fig. 14), respectively, with about80% of the CO flux originating in the warm component and 20%in the hot component (e.g., Goicoechea et al. 2012; Karska et al.2013; Green et al. 2013, 2016; Manoj et al. 2013; Lee et al. 2013;Dionatos & Güdel 2017; Karska et al. 2018; Yang et al. 2018).This constancy holds both across evolutionary stages as well asacross mass or luminosity (Karska et al. 2014a; Matuszak et al.
Article number, page 13 of 58 & A proofs: manuscript no. wish_final_Jan2021
Fig. 13.
Correlation of the water 1 −
557 GHz (HIFI) and 2 − / µ m (PACS) line intensities convolved to a commonangular resolution toward various outflow and source positions. Theoutflow data points (blue) are from Tafalla et al. (2013), whereas thesource positions (red) use data from Kristensen et al. (2012). This fig-ure demonstrates similar line ratios and thus similar conditions at the onand o ff source positions. Also, the two sources with strong EHV com-ponents, L1157 and L1448-MM (green), do not stand out in this figure. T >
500 K.A key step forward in the interpretation of these two CO ro-tational temperature components was made by Kristensen et al.(2017b) through velocity resolved CO J = T rot ∼
300 K component, and the medium-o ff set spot shocks to the ∼
700 K component.Taken together,
Herschel has shown that three ranges of COlines can be identified (Goicoechea et al. 2012; Yıldız et al.2013b; Karska et al. 2018): (i) low- J ( J up <
14) probing warmquiescent and entrained outflow gas; (ii) mid- J ( J up = J ( J up > E up / k B (K)404244464850525456 l og ( N up / g up ) + o ff s e t ( ) K ( ) K ( ) K ( ) K ( ) K ( ) K ( ) K ( ) K ( ) K W3-IRS5Orion BN/KLSer SMM1N1333-I4AN1333-I4B L bol ( L (cid:12) )10 Fig. 14.
CO excitation diagrams from low to high-mass protostars basedon (velocity-unresolved) PACS data. Two excitation temperature com-ponents of T rot ∼
300 and ∼
700 K are seen across the mass range.The Orion rotation diagram refers to Peak 1. Data from Herczeg et al.(2012); Karska et al. (2014a); Goicoechea et al. (2012, 2015).
To infer the physical conditions of thegas from these observed rotational temperatures, non-LTE ex-citation calculations need to be performed using programs suchas RADEX (van der Tak et al. 2007). Such analyses have beenprimarily carried out for LM and IM sources. For CO, it hasbeen demonstrated that both components require high densities( > cm − ), with kinetic temperatures T k of ∼ T k ≈ T rot << T k , canprovide good solutions to the entire CO excitation ladder as well(Neufeld 2012), but this solution is not consistent with the mul-tiple velocity components seen for high- J CO and H O line pro-files (Kristensen et al. 2017b) and will therefore not be consid-ered further here.Do these same conditions also reproduce the water excita-tion? Analyses of the HIFI lines using RADEX for a singleslab model demonstrate that water is indeed collisionally andsub-thermally excited. Best fits are obtained for densities in therange of 10 − cm − , both for the broad component and themedium-broad or o ff set spot shocks. Kinetic temperature is notwell constrained due to the limited range of energy levels cov-ered with HIFI but is typically a few hundred K. Moreover, to fitthe absolute fluxes, the size of the emitting area has to be smallcompared with the beam size, only of order ten to a few hundredau (Mottram et al. 2014).PACS observed much higher excitation lines of water. Anearly analysis for the line-rich low-mass protostar NGC 1333IRAS4B (Fig. 7, 14) resulted in a single component fit with T k ≈ n ≈ × cm − and an emitting radius of ∼ < (cid:48)(cid:48) ) centered in the blue outflow lobe located slightly o ff source (Herczeg et al. 2012). Thermal excitation at a much lowertemperature of 170 K and very high densities n > cm − onsource, as suggested based on mid-infrared Spitzer observations(Watson et al. 2007), was excluded in this case. Similarly, model-ing of the PACS spectral scan of the Serpens SMM1 intermediatemass source gave T k ≈
800 K, n ≈ × cm − with an emittingregion of radius ∼
500 au (Goicoechea et al. 2012). Comparableconditions are found for L1448-MM (Lee et al. 2013).In summary, all detailed analyses for low- and intermediate-mass sources agree that the on-source CO and H O emissionoriginates from high density n > cm − , warm ( T k (cid:38)
300 Kup to 1000 K) gas with a small emitting area. There is little ev-idence for a very hot component of several thousand K basedon far-infrared data. For the more distant high-mass sources, theanalysis is complicated by the fact that shock and envelope emis-sion are more intermixed. However, if the RADEX analysis islimited to the water wing emission, the conditions are inferred tobe similar, albeit with somewhat larger emitting areas up to 5000au (see §3.6 in San José-García et al. 2016).
Outflow positions.
In contrast with the protostellar positions,the velocity-resolved profiles of water lines at outflow spots inlow-mass sources can be markedly di ff erent from each other.Low excitation lines often show excess emission at high velocitycompared to the higher excitation lines. This can be interpretedin terms of density variations, with the gas at low velocity beingdenser than that at high velocity (Santangelo et al. 2012; Vastaet al. 2012). Indeed, when the water line intensities are consis- Article number, page 14 of 58.F. van Dishoeck et al.: WISH: physics and chemistry from cloud to disks probed by
Herschel spectroscopy tently fitted together with spectrally resolved high − J CO andmid-IR H observations, the physical conditions of the two wa-ter components are as follows (Santangelo et al. 2014b): a warmcomponent at T k ∼ − n ≈ − cm − , and size10 − (cid:48)(cid:48) ; and a hot component at T k > n ≈ − cm − , and size 1 − (cid:48)(cid:48) , for sources at ∼
200 pc (see also Busquetet al. 2014).These numbers are consistent with the analysis of the wa-ter excitation based on the 557 / nT ≈ × cm − K (Tafalla et al. 2013). Such pressures are about 4 or-ders of magnitude higher than the surrounding quiescent gas. For T k < n > cm − , consis-tent with the warm component. The association with H mid-infrared emission gives T k >
300 K as a lower limit.Thus, while the physical parameters at the outflow positionsare similar to those inferred at the protostellar positions (warmand dense), the emitting areas are much larger o ff source, withsizes up to a few thousand au for the warm component. Also,temperatures can be higher in the o ff set hot component. Thiscompact hot component is suggested to be associated with the jetimpacting the surrounding material in a bow shock, whereas theextended warm and dense component originates from the weakershocks further downstream. Since shocks at o ff -source positionscan expand in more directions than close to the protostellar base,the larger emitting area and somewhat lower densities found atthe outflow positions are not in contradiction. The absorption and narrow emission features probe the quiescentenvelope material. Their line profiles are asymmetric in a frac-tion of sources (see, for example, Fig. 2 and 10), and indicateeither infall motions (inverse P-Cygni) or expansion (regular P-Cygni) on envelope scales (few thousand au). The incidence ofinfall profiles is higher in low-mass Class 0 than Class I sources(Kristensen et al. 2012; Mottram et al. 2017). Still, the majorityof the Class 0 sources surprisingly do not show infall motions inany molecular feature (Mottram et al. 2017). The incidence ofexpansion is similar for Class 0 and I. Taken together, the WISH + WILL samples show infall in 7 + =
13 sources and expansionin 6 + = + =
78 sources. The smallsample of intermediate mass sources shows primarily P-Cygniprofiles indicative of expansion.The majority of the high-mass sources, about 2 /
3, show in-fall whereas the remainder show expansion (Herpin et al. 2016;van der Tak et al. 2019). Infall is seen not just in the outer coldenvelope gas but also in warm dense gas close to the protostarthrough H
O absorption (Jacq et al. 2016; van der Tak et al.2019). High mass sources often have many more narrow absorp-tion lines o ff set from the source velocity due to low density dif-fuse clouds along the line of sight, some of which may be closeby and infalling onto the high-mass core (Marseille et al. 2010;van der Wiel et al. 2010; Gerin et al. 2016).Quantitative modeling of the (outflow subtracted) water lineprofiles provides mass infall rates from cloud to envelope of or-der 10 − − − M (cid:12) yr − for low-mass protostars (Mottram et al.2013), to 10 − − − M (cid:12) yr − for high-mass protostars (van derTak et al. 2019). The latter values are comparable to those de-rived from inverse P-Cygni profiles observed in NH , either withHIFI or with the Stratospheric Observatory for Infrared Astron-omy (SOFIA) (Wyrowski et al. 2016; Hajigholi et al. 2016). Table 1 summarizes the profile characteristics and their proposedorigin. The key point is that the bulk of the broad water andhigh- J CO ( J up >
14) emission go together and arise in twophysical components associated with outflows and shocks, seenas T rot =
300 and 700 K in the CO data and as the broad andmedium-broad o ff set components. The implied physical condi-tions for both components are high densities ( n > − cm − )and with warm ( T k ≈
400 K) and hot ( T k ≈ ∼
100 au forlow-mass sources up to 5000 au for distant high-mass sources.The narrow emission and absorption lines probe the morequiescent envelope gas. Overall, the infall and expansion signa-tures highlight water’s ability to pick up small motions alongthe line of sight on envelope scales. This will prove to be ex-tremely useful to constrain the physical structure and abundanceprofiles of water and related molecules in the cold outer partsof envelopes (§ 7). Water is in this aspect more e ff ective thanwell-known envelope tracers such as low- J C O or HCO + . The ortho / para ratio of water, or equivalently its “spin” temper-ature, is in principle also a diagnostic of its physical and chem-ical history. At low temperatures ( <
10 K) the ortho / para ratiotends to 0 whereas at T >
50 K, the equilibrium ortho / para ra-tio would be 3. Thus, low ortho / para ratios have been invokedto trace a low formation temperature of H O on cold grain sur-faces (Mumma & Charnley 2011). However, the ortho / para ratiois controlled by a complex combination of processes (e.g., Tie-lens 2013; van Dishoeck et al. 2013), including possible changesin the ortho / para ratio upon ice desorption (Hama et al. 2018),and will not be summarized here. Overall, water spin tempera-tures may tell astronomers less about the water formation loca-tion than previously thought.Within the WISH project, the data are consistent with an or-tho / para of 3 for warm water (e.g., Herczeg et al. 2012; Mottramet al. 2014; Herpin et al. 2016). There have been some claims ofortho / para ratios lower than 3 in other sources (e.g., Hogerheijdeet al. 2011; Choi et al. 2014; Dionatos et al. 2018) but di ff erencesin optical depth of the ortho and para lines can also mimic suchlow ratios if not properly accounted for (Salinas et al. 2016).In fact, a more detailed analysis of many water lines observedwith HIFI in the Orion Bar shows values close to 3 (Putaud et al.2019).In summary, there is no convincing evidence for water ortho-to-para ratios significantly below the LTE value of 3 to be presentin the ISM.
5. Shocks
Assuming that the bulk of the water emission observed by
Her-schel originates from high temperature shocked gas, WISH + allows tests of chemical models of di ff erent shock types. Keyquestions to address are (i) whether the basic prediction that allvolatile oxygen not locked up in CO is driven into water at hightemperatures holds, in which case values of H O / H ≈ × − and H O / CO ≈ (1 . −
2) are expected (§ 3.1), (ii) what the rela-tive contributions of ice sputtering and gas-phase (re)formationare to the total gaseous water abundance.
Article number, page 15 of 58 & A proofs: manuscript no. wish_final_Jan2021
Table 2.
Summary of H O / H abundance determinations using Herschel data in warm outflows and shocks
Source Type Instrument H O / H H from Reference(10 − ) On source
Several LM HIFI 0.02 (LV + HV) CO 16-15 a Kristensen et al. (2017b)AFGL 2591 HM HIFI 0.001 CO 16-15 Kristensen, unpubl.NGC1333 IRAS4B LM PACS 1.0 high- J CO Herczeg et al. (2012)Serpens SMM1 LM / IM PACS 0.4 (hot) high- J CO Goicoechea et al. (2012)NGC 7129 IRS IM HIFI / PACS 0.2–0.3 high- J CO Johnstone et al. (2010)Fich et al. (2010)NGC 6334I HM HIFI 0.4 mid- J CO Emprechtinger et al. (2010)DR21(OH) HM HIFI 0.32 b high- J CO van der Tak et al. (2010)
Outflow
L1157 B1 LM HIFI 0.008(LV) mid- J CO Lefloch et al. (2010)0.8(HV)L1157 B2 / R LM-OF HIFI / PACS 0.01 warm H Vasta et al. (2012)L1448 B2 LM-OF HIFI / PACS 0.03 (warm) warm H Santangelo et al. (2012, 2013)0.03-0.1 (hot) hot H L1448 LM HIFI / PACS 0.005-0.01 warm H Nisini et al. (2013)VLA1623 LM-OF HIFI < c warm H Bjerkeli et al. (2012)Several B / R LM / IM-OF HIFI / PACS 0.003 warm H Tafalla et al. (2013)HH 54 OF HIFI < Santangelo et al. (2014a)L1157 B1 LM-OF HIFI / PACS 0.007-0.02 (warm) warm H Busquet et al. (2014)HIFI / PACS 1–3 (hot) hot H NGC 1333 I4A R2 LM-OF HIFI / PACS 0.007-0.01 (warm) CO 16–15 a Santangelo et al. (2014b)0.3–0.7 (hot)NGC 2071 B / R IM-OF PACS 0.3-0.8 warm H + high − J CO Neufeld et al. (2014)Orion-KL HM-OF PACS ≤ J CO a Goicoechea et al. (2015) > J CONote: uncertainties claimed by the authors range are typically a factor of 2 to a few. HV = high-velocity; LV = low; B = Broad com-ponent velocity. In most cases, the abundance refers to the outflow at the central source position. B and R indicated blue- andred-shifted outflow spots (OF) o ff set from the source. a assuming CO / H ≈ − ; if CO / H = × − advocated in § 3.1 is used, the water abundances are increased by a factor of two. b from para-H O assuming ortho / para = c ortho-H O abundance.
Table 2 summarizes measurements of the water abundance inwarm outflowing and shocked gas made by
Herschel . A list ofpre-
Herschel results, mostly using
S WAS or ISO, can be foundin Table 4 of van Dishoeck et al. (2013).
Herschel has advancedthe field in several ways. First, thanks to HIFI, the di ff erent ve-locity components (see § 4.3) can now be quantified separately.Although the precise origin of the broad component is still underdiscussion (see Table 1), this does not matter for the observa-tional derivation of abundances, and this component will there-fore be denoted in this section as shock. A second key advan-tage is the simultaneous measurement of (spectrally-resolved)high − J CO originating in the same gas. Moreover,
Spitzer ob-servations of the mid-IR H lines are available for several o ff -source shock positions providing a direct measure of the warmH (S(1)-S(5), few hundred K) or hot H (S(5)-S(9), > J CO or H data removes oneof the largest uncertainty in the X (H O) = N (H O) / N (H ) abun-dance determinations, namely the H column (the denominator). Early measurements often used the line wings of low − J CO linesto estimate the H column, which Herschel has now shown notto be appropriate (§ 4). Hence, any literature values using low- J CO as reference are not included in Table 2.The most reliable determination of the water abundance as afunction of velocity comes from the detailed study of 24 low-mass protostars by Kristensen et al. (2017b) using HIFI CO16–15 spectra as the reference. A remarkably constant H O / COabundance ratio of 0.02 is found, independent of velocity. Thisconstant abundance ratio even holds for the EHV “bullet” com-ponent such as seen in L1448-MM. Assuming CO / H = − ,both the warm broad and hot spot shock components are found tohave surprisingly low water abundances, with H O / H = × − .Only very hot gas may have a higher water abundance (Franklinet al. 2008): using the H O and CO PACS lines in line-rich low-mass sources, H O abundances up to ∼ − have been foundfor the hottest gas (Goicoechea et al. 2012).For high-mass protostars, few observations of CO 16–15 areavailable as reference. For the case of AFGL 2591, the observedline intensity ratio of H O 557 GHz / CO 16–15 is ∼ O / CO ∼ − when taking beam dilution into account and followingthe same analysis as for the low-mass sources (Kristensen, un- Article number, page 16 of 58.F. van Dishoeck et al.: WISH: physics and chemistry from cloud to disks probed by
Herschel spectroscopy −
20 0 20 υ - υ source (km s − ) T M B ( K ) I15398 x 3x 2Elias29 x 4x 5Ser-SMM1 / 8/ 6W3-IRS5 /450/600OH 1834 GHzH O 1 -1 Fig. 15.
Comparison of HIFI spectra of the H O 1 −
557 GHzline (blue) and the OH 1834 GHz triplet (red) toward a number of lowand high-mass regions. The splitting of the three OH lines in the tripletand their relative strength are indicated at the bottom of the figure. Datafrom Kristensen & Wampfler (unpublished). published). An early analysis of the DR 21 shock using CO10–9 as reference found an abundance ratio H O / CO ∼ × − in the shock. Thus, the H O abundance in high-mass outflowsmay be even lower than that inferred for low-mass protostars,although not as low as found by Choi et al. (2015). Indeed,van der Tak et al. (2019) conclude based on H
O absorptionin line wings toward high-mass protostars that water abundancesin outflows are an order of magnitude higher than in the coldenvelopes, where values of 10 − − − have been found (Boon-man et al. 2003a, and § 7 and 8). Taken together, the H O / H abundances for high mass outflows are found to be 10 − − − . At the outflow positions, where direct measurements of the warmH columns are available, the largest uncertainty stems from theuncertainty in the physical parameters used to determine the H Ocolumn densities, since the same data can be fitted with di ff erentwater column densities and filling factors. Nevertheless, if theH O, high- J CO and H data are fitted simultaneously, the abun-dance of the warm H O component is better constrained. The de-rived abundance is universally found to be low, ranging between10 − − − (Santangelo et al. 2012, 2013; Bjerkeli et al. 2012;Nisini et al. 2013; Tafalla et al. 2013; Busquet et al. 2014; San-tangelo et al. 2014a). There are suggestions that the abundance inthe hot component ( > > × − ,but this is more di ffi cult to constrain with the available data (San-tangelo et al. 2013, 2014b). Abundances close to 10 − have beenfound only for the very hot gas seen in the highest- J H O andCO lines but the lack of H O lines originating from very highenergy levels ( > Herschel data make it more dif- ficult to constrain this component (Lefloch et al. 2010; Busquetet al. 2014; Neufeld et al. 2014; Goicoechea et al. 2015).Melnick et al. (2008) combined mid-infrared observationsof pure rotational H O and H lines to infer water abundancesH O / H in the range (0 . − × − for shock positions in theintermediate mass source NGC 2071 o ff set from the source posi-tion. These mid-infrared data probe higher excitation H O levelsthan the
Herschel data and are thus better suited to test the waterpredictions for > O / H ≈ (1 . − . × − for the IRAS 17233–3606 out-flow by modeling consistently both Herschel H O and H µ minfrared data, finding a maximum fractional water abundance inthe shock layer of 10 − .Independent evidence that H O does not contain all of theoxygen comes from mid-infrared absorption line observationswith ISO and SOFIA of hot water and CO along the lines ofsight to massive protostars, many of them the same sources asincluded in the WISH sample (van Dishoeck & Helmich 1996;Boonman & van Dishoeck 2003). In particular, high spectral res-olution absorption line observations at 6 µ m using SOFIA-EXESgive N (H O) / N (CO) ≈ . ∼ Herschel are two–three orders of magnitude lower than predicted by stan-dard high temperature (shock) models that drive all volatile oxy-gen into water. This conclusion holds in both velocity compo-nents on-source, and at o ff -source positions as well. Water abun-dances as high as ∼ − are perhaps only attained in the veryhot ( > Herschel . Eventhen, there is no firm observational evidence yet for any shockassociated with outflows in which H O / H is as high as 4 × − ,the value expected if all volatile oxygen is in H O. See § 10 andAppendix G for further discussion. O vs OH
A key intermediate species in the O → H O reaction is OH,so given the fact that the H O abundance is found to be low,does this imply that the OH abundance is high? WISH targeteda number of OH transitions with PACS, particularly the low-lying transitions at 119, 84 and 163 µ m (Fig. 3), and coveredlow-, intermediate-, and high-mass sources. The PACS instru-ment did not resolve the hyperfine components of the OH transi-tions (the splitting is typically < − ), but it did resolve the Λ doubling. Strong indications for an enhanced OH abundancecome from the observed OH / H O flux ratio in a large number oflow-mass protostars, which is found to be much higher than ex-pected based on models of shocks propagating into dark molec-ular clouds (Karska et al. 2014b, 2018). Strong OH emission isalso detected with PACS for high-mass protostars (Karska et al.2014a).Apart from these velocity-unresolved PACS observations,the 1834 and 1837 GHz OH triplets, which make up the 163 µ m lines, were observed with HIFI and detected toward W3IRS5 and Ser SMM1 (Wampfler et al. 2011; Kristensen et al.2013). The 1837 GHz triplet is completely resolved toward W3IRS5 and consists of narrow ( ∼ − ) hyperfine componentssuperposed on top of a broad ( ∼
20 km s − ) outflow compo-nent. Toward Ser SMM1, the 1834 GHz transition consists of asingle broad component, which appears blue-shifted ( ∼
10 kms − ) (Fig. 15). The hyperfine splitting is actually smaller for thistriplet than for the 1837 triplet. The 1837 GHz line was also ob- Article number, page 17 of 58 & A proofs: manuscript no. wish_final_Jan2021 served toward a number of other low-mass sources within WISHbut not detected. Comparing their limits with the PACS 163 µ mdetections implies that the lines for these sources must also bebroad (FWHM at least 10-15 km s − ) (Wampfler et al. 2010).As part of WISH + , the COPS-HIFI program observed the1834 GHz triplet together with CO 16–15 for 24 sources (Kris-tensen et al. 2017b, Kristensen & Wampfler unpublished) (Fig-ure 15). Where detected, the OH line is broad and follows H O.The OH abundance with respect to H O in the broad outflowcomponent has been constrained to > / H O abundance ratios > O and thus increasing theabundance of OH. However, in no case has OH been found tolock up a significant fraction of the oxygen, perhaps implyingthat gas temperatures remain below the activation threshold ofthe critical reactions (see also § 5.4.2, Fig. 16).
The cold outer protostellar envelopes are punctured by outflowcavities on scales of a few thousand au. Besides shocks stirringup the gas, UV radiation can escape through the cavities and im-pinge on the walls thereby heating them up (Spaans et al. 1995).Thus, outflow cavity walls are a source of bright far-infrared lineradiation and simple hydrides such as CH + , OH + and H O + turnout to be a good probe of them. These hydrides are readily de-tected in the WISH data for many sources, often in absorption,with their velocity o ff sets pointing to a spot shock or a more qui-escent cavity wall origin. The results of the WISH hydride sub-program are summarized in Benz et al. (2016) with the analysisbuilding on detailed modeling of the e ff ects of UV and X-rayson protostellar envelopes by Stäuber et al. (2005) and Brudereret al. (2009, 2010).In brief, the observed hydride fluxes and flux ratios (e.g.,CH + / OH + ) can be reproduced by 2D models of UV illuminatedoutflow cavity walls on scales of the Herschel beam. The im-plied UV fluxes are up to a few × times the interstellar ra-diation field, at distances of ∼ ∼ (cid:12) , is required. Alternatively, some ofthe UV can be produced by the high-velocity shocks themselves.For high-mass regions, the FUV flux required to produce the ob-served molecular ratios is smaller than the unattenuated flux ex-pected from the central object(s) at that radius, implying someextinction in the outflow cavity or, alternatively, bloating of theprotostar (Hosokawa et al. 2010). Another important conclusionis that there is no molecular evidence for X-ray induced chem-istry in low-mass objects on the observed scales of a few ×
200 400 600 800 1000 T (K)-1012345 l og ( G ) - - - - - - - − − − − l og ( X ( H O )) Fig. 16.
Water abundance with respect to H as function of temperatureand UV radiation field using a simplified high-temperature chemistrymodel. The assumed density is 10 cm − ; for higher or lower densitiesthe abundance curves shift up or down, respectively, but the trend is thesame. Based on Kristensen et al. (2017b). Early shock models under dark cloud conditions find completeconversion of all volatile oxygen into water for any shock thathas temperatures greater than ∼ + H and OH + H reactions drive rapid wa-ter formation (Draine et al. 1983; Kaufman & Neufeld 1996;Flower & Pineau des Forêts 2010) (Fig. 4). The detection ofhigh water abundances in the Orion shock by ISO was cited asa confirmation of this prediction (Harwit et al. 1998). For a non-dissociative shock, the temperature just behind the shock is typ-ically T s = b M − . [ v s /
10 km s − ] . with the magnetic fieldstrength b M = B / ( n H / cm − ) / µ G usually taken to be standard b M = v s = −
15 km s − to obtain the temperatures of a few hundredK required to drive the reactions.If instead most oxygen is locked up in water ice in the pre-shocked gas, higher shock velocities are needed to sputter wa-ter ice, typically (cid:38)
15 km s − . Such shocks have peak tem-peratures of T s ≈
650 K (Draine et al. 1983; Jiménez-Serraet al. 2008; Gusdorf et al. 2011). Standard theoretical modelsfor non-dissociative shocks predict the complete vaporization ofice mantles resulting in H O / CO ∼ (1 . −
2) (§ 3.1). Even if theice comes o ff as atomic O or OH, it should be quickly convertedto H O at these temperatures and densities, if shielded from dis-sociating radiation.In dissociative J -type shocks much higher temperatures > K can be reached at shock velocities (cid:38)
25 km s − (McKee& Hollenbach 1980), which are su ffi ciently high to collisionallydissociate H and H O. These molecules, together with CO, sub-sequently reform downstream in the cooling shock gas resultingin similar abundances.In conclusion, the low water abundances and low H O / OHintensity ratios found by WISH + are clearly in conflict with thestandard shock model predictions. Article number, page 18 of 58.F. van Dishoeck et al.: WISH: physics and chemistry from cloud to disks probed by
Herschel spectroscopy
The above findings have led to the development of a new class ofUV irradiated shock models (Melnick & Kaufman 2015; Godardet al. 2019). In such models, the H O abundance is lowered byphotodissociation into OH and O. The UV radiation can eithercome from the shock itself, if fast enough (Neufeld & Dalgarno1989), or it can be external, for example from the disk-star ac-cretion boundary layer (Spaans et al. 1995).Lowering the water abundance in the shocked gas by pho-todissociation is the simplest e ff ect of UV radiation. If the pre-shocked gas is also UV irradiated and some of its H dissociated,the entire shock structure changes: the UV photons increase theatomic and ionization fraction in the pre-shocked gas, resultingin a shock layer that is more compressed (smaller) and hotter fora given shock velocity and density. The UV radiation can alsounlock more oxygen from water ice in the pre-shocked gas andthus increase the amount of atomic oxygen that can be convertedinto water.Results from a limited set of irradiated shock models by M.Kaufman (unpublished) for UV fields of G = − cm − indeed show much better agreement with Her-schel data than the older models when G > G = − close to the protostar on scales of ∼ J CO emission lines(Spaans et al. 1995; Yıldız et al. 2012; Visser et al. 2012; Leeet al. 2013; Yıldız et al. 2015). Moreover, as noted above (§ 5.3),the presence of several hydrides such as CH + and OH + for bothlow- and high-mass protostars obtained as part of WISH pointsto UV-irradiated outflow walls with similar G values (Brudereret al. 2009, 2010; Benz et al. 2010, 2016).Figure 16 illustrates the results of a simple chemical modelbalancing the high temperature formation of water through thereactions of O and OH with H with photodissociation of H O,where the UV radiation field is given by its enhancement G withrespect to the standard interstellar radiation field (see Kristensenet al. 2017b, for details). For G values of 10 − , the expectedwater abundance is indeed in the 10 − − − range at temper-atures of several hundred K. Temperatures > − if UV radiation is present. T chemistry vs sputtering: SiO, CH OH, NH To what extent is the water seen in shocks produced by high tem-perature gas-phase chemistry, as assumed in Fig. 16, versus icesputtering? Shocks with velocities of 15 km s − needed for waterice sputtering are observed for most embedded sources, exceptperhaps for the low-mass Class I objects. To better distinguishthe two mechanisms and isolate the contribution due to sput-tering, it is useful to compare the H O line profiles with thoseof other abundant ice mantle constituents, most notably CH OHand NH .Figure 17 compares H O and CH OH line profiles for a low,intermediate and high mass protostar, observed with
Herschel -HIFI or from the ground in a similar beam (Suutarinen et al.2014; van Kempen et al. 2014; Herpin et al. 2016). The H Oprofiles clearly extend to much higher velocities than those ofCH OH, with the latter profiles dropping to zero beyond ± − from source velocity, independent of the source lumi-nosity. A quantitative analysis shows one order of magnitude − −
25 0 25 50 υ - υ source (km s − ) . . . . . . . T M B ( K ) N1333-I4ANGC2071 x 0.2x 3N6334-I(N) / 3CH OHH O 2 -1 Fig. 17.
Comparison of the H O 2 −
988 GHz line with that of anice species, CH OH, for a low, intermediate and high-mass source. In allcases the CH OH line is seen to be much narrower than that of H O. ForNGC 1333 IRAS4A, the 7 − line ( E up =
78 K) at 338.4 GHz observedwith the JCMT (15 (cid:48)(cid:48) beam) is shown; for NGC 2071 the 6 − line( E up =
86 K) at 766.7 GHz observed with HIFI is included (van Kempenet al. 2014) and for NGC 6334-I(N) the 5 − line ( E up =
56 K) at 538.5GHz observed with HIFI is presented. decrease in the CH OH / H O column density ratio as the veloc-ity increases to 15 km s − (Suutarinen et al. 2014). A similardecrease in the CH OH / CO ratio is found, demonstrating thatCH OH is destroyed in higher velocity shocks, and that not justH O increased at larger velocities. A similar lack of CH OH(and H CO) emission compared with H O at higher velocitieshas been seen at shock positions o ff source (e.g., Codella et al.2010; Busquet et al. 2014; Vasta et al. 2012).These data suggest that sputtering of ices does indeed oc-cur but that the process either shuts o ff at velocities above 15km s − , or that any CH OH that is sputtered at higher velocitiesis destroyed. Destruction of molecules can occur either duringsputtering or by reactions with atomic H in the shock. The factthat H O is seen up to much higher velocities implies that eitherH O is not destroyed (see below) or that gas-phase formationof H O must be active above ±
10 km s − to re-form any wa-ter destroyed during sputtering or by other processes such as UVphotodissociation (Fig. 16). The data also suggest that sputteringtakes place at somewhat lower velocities than the expected 15km s − , or that the sputtered molecules have been slowed downsince being released into the gas.For NH , the same phenomenon is observed (Fig. 18). NH is also found to emit only at low velocities up to ±
15 km s − atshock positions o ff source (e.g., Gómez-Ruiz et al. 2016). SinceNH has similar excitation requirements as H O, this must be achemical e ff ect. Similar to the case for CH OH, destruction byatomic H in the shock is proposed (Viti et al. 2011). However, thebarrier for the H + NH reaction is about 5000 K, much higher Article number, page 19 of 58 & A proofs: manuscript no. wish_final_Jan2021 −
50 0 50 υ - υ source (km s − )02468 T M B ( K ) H O 1 − NH − + − / OH 7 − CO 10 − / − × Fig. 18.
Comparison of the line profile of water with those of otherchemically or physically related molecules for the low-mass protostarNGC 1333 IRAS4A. SiO and CO 10–9 trace shock and outflow; NH and CH OH are both ice mantle products, similar to H O. The fact thattheir profiles are narrower suggests that H O is not just sputtered fromices but also produced by high temperature chemistry in the gas. Theanticorrelation of the HCO + abundance with that of water is illustratedby its lack of broad line profiles. than that for H + CH OH of about 2200–3000 K. The barrierfor H + H O destruction is even higher, ∼ K, allowing H Oto survive following sputtering. If only temperature-sensitive de-struction would play a role, the relative abundances of these threemolecules could provide a thermometer for the higher velocitygas. More likely, however, high-temperature gas-phase forma-tion of H O in the high velocity gas controls the di ff erence be-tween these three species.Shocks are also expected to sputter silicates from grain man-tles and cores for shock velocities above 20–25 km s − , with SiOlong known to be one of the best tracers of shocks (e.g., Martín-Pintado et al. 1992; Caselli et al. 1997; Schilke et al. 1997;Jiménez-Serra et al. 2008). How do the SiO and H O profilescompare? Figure 18 presents the case of NGC 1333 IRAS4Ashowing similarly broad profiles for the two molecules, as alsofound in some other sources (e.g., Leurini et al. 2014). In othersources such as L1448-MM, SiO is only found in the EHV bul-lets but not in the broad component (Nisini et al. 2013). In yetother cases, SiO is found at intermediate velocities (e.g., Vastaet al. 2012; Busquet et al. 2014). One explanation for L1448-MM is that the jet gas is rich in atomic Si originating from itsdust-free launch position in the inner disk. This then leads toSiO formation in the jet by gas-phase processes, a mechanism proposed originally by Glassgold et al. (1991) and revisited the-oretically by Tabone et al. (2020), not because of sputtering ofambient dust grains. Whether SiO is also seen in the broad com-ponent then depends on the details of the wind-cloud interaction.An alternative option is that these di ff erences reflect the timeevolution of the SiO profiles going from high to low velocitiesas the shocked gas slows down, as modeled by Jiménez-Serraet al. (2009). Spatial evidence for SiO evolution from high tolower velocities with distance from the source (or equivalently,time) is found in ALMA images of outflows (Tychoniec et al.2019). Clearly, H O line profiles are less sensitive to such timeevolution e ff ects. O vs HCO + As Figure 4 shows, HCO + is e ff ectively destroyed by reactionswith H O so that an anticorrelation between these two speciesis expected, both in position and velocity (e.g., Phillips et al.1992; Bergin et al. 1998). This anticorrelation has recently beendemonstrated observationally through mm interferometric im-ages of H
O (or CH OH as a proxy) and H CO + (Jørgensenet al. 2013; van ’t Ho ff et al. 2018a). These lines trace quiescentwarm envelope gas, however, not the shocked gas.The broad shocked H O gas line profiles observed with
Herschel -HIFI can be compared with HIFI HCO + J = + J = Herschel beam (Carney et al. 2016). Figure 18shows that HCO + generally avoids the high velocities and islimited to ± − around source velocity, although in somesources a weak underlying component out to ±
10 km s − is seen(e.g., Kristensen et al. 2010; San José-García 2015; Carney et al.2016; Benz et al. 2016; Mottram et al. 2017). HCO + thus seemsto trace primarily parts of the envelope that are away from loca-tions in the outflow where water is enhanced, consistent with thechemical expectations.In summary, the combination of H O data with those of otherspecies confirm the chemical schemes outlined in Fig. 4, includ-ing the anticorrelation with HCO + . Both high temperature chem-istry and ice sputtering contribute to the water abundance, withice sputtering limited to the low velocities. UV radiation reducesthe water abundance in outflows to orders of magnitude belowthe expected abundance of H O / H = × − .
6. Cold dense pre-stellar cores
Herschel -HIFI was unique in its ability to obtain velocity-resolved water line profiles in cold dense clouds, and will remainso in the coming decades. Such data allowed a much deeper anal-ysis than just deriving beam-integrated column densities, as isnormally done from spectra taken at a single position. In fact,the detailed line profiles combined with water’s ability to probemotions at a fraction of a km s − make it possible to reconstructthe actual water abundance profile as a function of position inthe core, in spite of not spatially resolving or mapping the cloud.This allows stringent testing of the primary chemical processescontrolling the gaseous water abundance. At the same time, con-straints on physical parameters that are di ffi cult to determine oth-erwise are obtained, most notably the external and internal UVradiation fields that control photodesorption and photodissocia-tion (Caselli et al. 2012; Schmalzl et al. 2014).A similar analysis can be carried out for other species, inparticular NH observed with Herschel-
HIFI. Since NH is also Article number, page 20 of 58.F. van Dishoeck et al.: WISH: physics and chemistry from cloud to disks probed by
Herschel spectroscopy abundant in ices, comparison with H O can be insightful intohow, where and why the nitrogen chemistry in dark clouds isso di ff erent from that of other species (Hily-Blant et al. 2017;Caselli et al. 2017).This section describes the results of the water abundance pro-file analysis for dense cores prior to star formation combinedwith our new analysis of the NH profile; the next section doesso for the cold outer parts of protostellar envelopes which havean internal heat source. In both cases, simplified chemistry net-works as described in § 3.2 are coupled with a physical struc-ture of the source. These chemistry models can be run either insteady-state or in a time-dependent mode at each position in thecloud. For pre-stellar cores at densities > cm − the chem-istry is faster than the free-fall time of 0.03 Myr, so the chem-istry reaches steady-state before the physical conditions change(Keto et al. 2014). In the time-dependent case, abundances at t = Within the WISH program, only one low-mass pre-stellar core,L1544, was observed deep enough with HIFI in the ground-stateo-H O line for detection and detailed analysis (13.6 hr integra-tion) (Caselli et al. 2012). The line is very weak but shows aninverse P-Cygni profile. Assuming an ortho / para ratio of 3, theinferred H O column for L1544 integrated over the entire lineof sight is > × cm − from the absorption part of the fea-ture. This corresponds to a fractional gaseous water abundanceH O / H > . × − , with the lower limit stemming from themoderate optical depth of the line. The fact that water is seen inemission in the blue-shifted part of the line profile implies highcentral densities, about 10 cm − . The same deep spectrum alsoincludes the o-NH − transition, showing clear emissionlines and resolving the hyperfine structure for the first time inspace (Caselli et al. 2017).A second starless core observed in WISH, B68, has a slightlyless stringent upper limit of N (o-H O) < . × cm − , ora fractional abundance H O / H < × − . Both of these re-sults re-inforce the original conclusions from S WAS and
Odin (Bergin et al. 2002; Klotz et al. 2008) that the gaseous waterabundance is very low in cold clouds and that most of the watermust be locked up as water ice.Two additional very deep integrations (15 hr each) were ob-tained in an OT2 program (PI: P. Caselli; Caselli et al., unpub-lished) on the low-mass cores L183 and Oph H-MM1. Bothshow weak absorption features, with that for Oph H-MM1 in-cluded in Fig. 10. As for B68, their lower central densities com-pared with L1544 do not favor water emission but their centralcontinuum emission is strong enough to absorb against. Upperlimits on the gaseous H O abundance integrated along the lineof sight are again low, < − . Four high mass cores that were thought to have no associated starformation prior to
Herschel launch were targeted within WISH.The low-lying H O lines are detected in all four sources (Ship-man et al. 2014). Three of the four sources show outflow wings,however, demonstrating that they are not truly pre-stellar. Never-theless, after subtraction of the outflow component, their inverseP-Cygni profiles greatly resemble those of L1544 and of the low- au Fig. 19.
Best-fit H O (top) and NH (bottom) abundance profiles (withrespect to H ) as functions of position in the NGC 1333 IRAS4A pro-tostellar envelope (full lines) and the L1544 pre-stellar core (dashedlines) based on analysis of HIFI spectra. Note the similarity of the outerwater abundance profiles for the pre-stellar core and outer protostel-lar envelope, and the very di ff erent chemical behavior of H O versusNH . Based on results from Caselli et al. (2012); Mottram et al. (2013);Caselli et al. (2017) and new analyses (see Fig. 22). The inner hot coreH O abundance for IRAS4A is taken from Persson et al. (2016). mass protostars discussed below, but over a wider range of ve-locities. The fourth source, G11.11-0.12-NH , shows no signs ofstar formation. It has narrow (FWHM ≈ − ) water absorp-tion lines (Fig. 10), even in H O, that allow an accurate watercolumn density determination. Assuming ortho / para =
3, the in-ferred abundance is again low, 4 × − . O and NH abundance profiles Herschel -HIFI confirmed the overall picture of low gas-phasewater and high ice abundances in cold clouds, but could do muchmore by inferring the water abundance profiles across the cores.These profiles can then be compared with models such as thoseby Hollenbach et al. (2009) through coupling of the chemistrywith a physical model of the source to infer critical parameters(Caselli et al. 2012).L1544 is an excellent example, since it has gradients in tem-perature, density and velocity through the core that are well con-strained by other observations. The observed inverse P-Cygniprofile allows the blue-shifted emission part to be connected withthe dense central part of the core, and the red-shifted absorptionwith the infalling lower-density outer part. The strong absorptionindeed points to a relatively high water abundance of ∼ − inthe outer layer as expected due to photodesorption of water iceby the external UV radiation field (Fig. 19, Appendix D). Thecentral emission requires a significantly higher water abundancethan the original value of 10 − in the Hollenbach et al. model.Adding cosmic-ray induced UV photodesorption of water ice tothe model raises the central abundance to ∼ − − − , suf- Article number, page 21 of 58 & A proofs: manuscript no. wish_final_Jan2021 ficient to explain the emission. The best fitting water abundanceprofile for L1544 is included in Fig. 19.If grain surface formation followed by photodesorptionwould dominate the production of gas-phase NH as well, itsabundance structure should follow that of H O. However, suchan abundance structure gives a poor fit to the data (Caselli et al.2017). A much better fit is obtained with a constant or slightlyincreasing (rather than decreasing) abundance structure towardthe center. Such an abundance profile was previously inferredfrom radio observations with the Very Large Array (VLA) of p-NH by Crapsi et al. (2007). Thus, gaseous NH and H O havevery di ff erent distributions in dense cores, with a much highercontribution of cold gas-phase chemistry than ice chemistry forNH compared with H O (Le Gal et al. 2014; Sipilä et al. 2019).
7. Protostellar envelopes: Cold outer part
Is this behavior for H O and NH also seen for protostellarsources? The cold outer parts of protostellar envelopes are inmany aspects similar to pre-stellar cores, but protostars have amore strongly centrally concentrated density structure as wellas an internal heating source, facilitating both absorption in theouter layers and emission from the inner part of their envelopes.As discussed in § 4.3 and 4.5, the low-lying H O lines towardprotostars –from low to high mass– indeed show narrow (in-verse) P-Cygni profiles superposed on the outflow componentsin some fraction of sources (Fig. 11), indicating either infallor expansion. When the outflow components are subtracted, theresidual line profiles can be analyzed in a similar way as thosefor the pre-stellar cores.
To analyze the water data, the temperature and density struc-ture of protostellar envelopes has been coupled with the (sim-plified) SWaN water network (§ 3.2) (Schmalzl et al. 2014).An important di ff erence with pre-stellar cores is the tempera-ture structure: as soon as a protostar has turned on in the cen-ter, a temperature gradient is established throughout the enve-lope, with temperature decreasing with radius R as roughly R − . in the region where the far-infrared dust emission is opticallythin. The actual dust temperature structure can be computed withfull continuum radiative transfer calculations given a luminosityand density structure of the core, both fitted to the SED and ex-tent of the submillimeter continuum emission assuming spheri-cal symmetry (e.g., Jørgensen et al. 2002; Kristensen et al. 2012).Gas and dust temperatures are taken to be coupled, which is avalid assumption at these high densities. The chemistry is runin time-dependent mode. The initial abundances for the proto-stellar phase are taken from a model of a cold, constant density( n H = × cm − ) pre-stellar cloud at an age of ∼ T <
100 K, the abundance structure is remark-ably similar to that of L1544, with a water peak abundance of ∼ − due to ice photodesorption at A V ≈ few mag and then a rapid decrease in gaseous water abundance going inward. Thissteep decrease is largely due to the steeply increasing H density,which is the denominator in H O / H , coupled with H O freeze-out on grain mantles. Once the threshold for thermal desorptionis reached, the gaseous water abundance quickly returns to theoverall oxygen abundance in the model of ∼ − .This general behavior has been demonstrated in detail fora number of low-mass protostars, both Class 0 (Mottram et al.2014) and Class 0-I (Schmalzl et al. 2014). The main parametersthat control the shape of the line profiles are the external UV ra-diation G ISRF and the internal cosmic ray field G CR (all in unitsof the Habing 1968 field d ), together with the details of the ve-locity profile (infall or expansion). Best fit models generally findan internal cosmic ray induced UV field that is slightly belowthe normal value ( G CR (cid:46) − ), with 10 − the standard num-ber for a cosmic-ray ionization rate of ∼ − s − (Shen et al.2004). More precisely, the amount of water vapor produced byphotodesorption scales with the product n gr σ gr G CR with n gr thenumber of grains and σ gr its geometrical cross section. Thus,the results scale with the inverse of grain radius a , and some-what lower values of this product can also imply dust growth tomicron-sized particles deep inside the core, as has been foundobservationally (Pagani et al. 2010).The results are insensitive to the timescale of the protostellarphase t proto over the 0.1–1 Myr range. The absolute values are,however, highly sensitive to the duration of the pre-stellar stage.In fact, the observations of both water gas and ice for the sameline of sight can only be reconciled within this modeling frame-work if a rather short pre-stellar phase is assumed of ∼ O 1 −
557 GHz line profiles. For thissource, both a low external and internal radiation field are re-quired, together with a radially expanding velocity field to getthe P-Cygni line profile. The best-fitting H O abundance profilein model D is similar to that found in Fig. 19 for IRAS 4A.More generally, the inferred external UV field is very lowfor all sources analyzed, G ISRF ≈ .
01, suggesting that the outerprotostellar envelopes on scales of ∼ A V = − G ISRF > The alternative approach to analyze water data is to use a simplestep-function profile with an outer (freeze-out at T <
100 K) andinner (sublimation >
100 K) abundance. A slightly more sophis-ticated method is a “drop abundance” profile, with the latter in-cluding an outer ice photodesorption layer (Coutens et al. 2012).As shown in Mottram et al. (2013, their Fig. 4 and 14), the mainfeatures of the water chemistry are captured by the drop profile,but the outer abundances have a limited meaning. For example,for NGC 1333 IRAS4A the best fit drop abundance profile has anouter abundance of 3 × − and a photodesorption layer abun- d To convert to units of Draine (1978) field, divide G by a factor of1.7Article number, page 22 of 58.F. van Dishoeck et al.: WISH: physics and chemistry from cloud to disks probed by Herschel spectroscopy -4 0 4 υ − υ source (km s − )-1000100200 T M B ( m K ) t post =1.0 Myr G ISRF =10 G cr =10 − Model A β , υ r ( k m s − ) βυ r − − − − − − − X H O A V (mag) -4 0 4 t post =1.0 Myr G ISRF =10 G cr =10 − Model B − -4 0 4 t post =1.0 Myr G ISRF =10 G cr =10 − Model C − -4 0 4 t post =1.0 Myr G ISRF =10 − G cr =10 − Model D − Fig. 20.
Illustration how various physical and chemical parameters can be constrained using the SWaN network. H O 1 −
557 GHz observa-tions of the Class I source L1551-IRS5 are compared with SWaN models assuming di ff erent parameters for the FUV fluxes and velocity profile.The top row shows the observations (black) together with the synthetic spectra (red), whereas the bottom row shows the adopted velocity profile(green) and water abundance profile (with respect to H ) (black). A constant Doppler broadening of β = − is assumed, whereas the radialvelocity profile is either taken to be constant at 0 (Models A and B) or to have an expansion profile (Models C and D). The radiation fields G ISRF and G CR strongly decrease from model A to model D. A pre-stellar stage duration of 0.1 Myr is used. Based on Schmalzl et al. (2014). Fig. 21.
Inner hot core (red circles) and outer cold envelope (blue cir-cles) water abundances with respect to H in high mass protostars, de-rived using a step-function analysis. Inferred abundances for a low mass(IRAS16293-2422) and an intermediate mass (NGC 7129) source areshown for comparison. Uncertainties are typically factors of a few, butsee text for the limited meaning of outer abundances. The gray bar at thetop shows the warm water abundance expected if all oxygen is driveninto water. Based on Herpin et al. (2016); Choi (2015); Coutens et al.(2012); Johnstone et al. (2010). dance of 3 × − , whereas Figure 19 shows a steadily decreasingabundance profile between 10 − and 10 − .One intermediate (NGC 7129) and thirteen high-mass proto-stars have been analyzed using the step-function profiles (John-stone et al. 2010; Marseille et al. 2010; Chavarría et al. 2010;Herpin et al. 2016; Choi 2015). The inferred outer abundances X out (with respect to H ) range from 10 − to 10 − and are plottedin Figure 21. There is no obvious trend with luminosity or enve-lope mass, which is not surprising given its sensitivity to externalUV and the discussion above on analysis methods. O vs NH abundance profile HIFI observations of the NH −
572 GHz line have beentaken for a number of low-mass sources as part of an open timeprogram (PI: P. Hily-Blant). As for pre-stellar cores, the gas-phase NH abundance in the outer envelope would be expectedto follow that of H O if both were produced primarily on thegrains and returned to the gas by photodesorption of ice. How-ever, such an abundance profile is again a poor fit to the NH spectrum in the best studied case of NGC 1333 IRAS4A. Thebest-fit NH abundance is a constant abundance of 7 × − overmuch of the envelope, with a slightly lower value in the out-ermost photodissociation layer (Fig. 22). Overall, the inferredNH abundance profile for this low-mass protostar is similar tothat found for the pre-stellar core L1544 (Fig. 19). The similarityof the observed NH spectra for other sources suggests that thisabundance structure is a common feature.The inferred constant NH abundance is inconsistent withcurrent chemical models (Caselli et al. 2017). Yet NH ice is ob-served at an abundance of a few % of that of H O ice in a widevariety of sources (e.g., Bottinelli et al. 2010; Öberg et al. 2011a;Boogert et al. 2015), so the abundance of s − NH is known tobe about 10 − with respect to hydrogen. The bulk of the ob-served gaseous NH with HIFI has an abundance of a few × − ,which is a few % of that of ice. This value is higher than canbe explained by photodesorption of ice. Thus, the observed gas-phase NH is likely formed primarily through gas-phase reac-tions (Le Gal et al. 2014) in a route that preserves a constantabundance with increasing density and that is apparently not in-cluded, or with a too low rate coe ffi cient, in current gas-phasemodels (Caselli et al. 2017; Sipilä et al. 2019). Lines of deuterated water, HDO, have been observed as part ofthe (extended) WISH program, most notably the 1 − lineat 893 GHz connecting with the HDO ground state. Figure 23provides examples of high quality HIFI 1 − spectra near 1 Article number, page 23 of 58 & A proofs: manuscript no. wish_final_Jan2021 − − υ - υ source (km s − )0 . . . . . T M B ( K ) NH − − H O 2 − au au Fig. 22.
Left: Outflow-subtracted HIFI line profiles of H O, HDO and NH (blue) toward the low-mass protostar NGC 1333 IRAS4A overlaidwith the best fit model line profiles (red). These model spectra correspond to the abundance structures presented in the middle and right panels.The inner H O and HDO abundances are set to those found by Persson et al. (2014, 2016).
THz of both H O, H
O and HDO, all obtained in a similar beamof ∼ (cid:48)(cid:48) .Similar to the H O lines, the 893 GHz line is very well-suited to determine the HDO abundance profile, and thus theHDO / H O abundance ratio, in the cold outer envelope. Higher-lying HDO lines such as 2 − at 241 GHz and 3 − at225 GHz, both of which can be observed from the ground, aremore sensitive to the inner warm HDO abundance (see § 8).The HDO abundance profile has been determined using thesame procedure as for H O and NH . The simplest assumption isthat HDO follows H O but scaled by a constant factor. Figure 22shows that this assumption works very well for HDO: a constantHDO / H O abundance ratio of 0.025 provides a very good fit tothe observed HDO and outflow-subtracted H O line profiles forthe case of the low-mass protostar NGC 1333 IRAS4A. A flatabundance profile, such as for NH , does not fit well. Thus, HDOfollows H O.Since the observed gaseous H O results from photodesorp-tion of water ice, this result could imply that the same holds forHDO. Detailed modeling of the H O and HDO photodesorptionprocesses by Arasa et al. (2015) has shown that di ff erences in ef-ficiencies are very small, so that no significant corrections to theobserved numbers are needed because of di ff erences in desorp-tion e ffi ciencies. However, since photodesorption only proceedsfrom the top few ice layers, these data would then only probeHDO / H O in the outermost ice layers, not in the bulk of the ice(Fig. 5). Moreover, cold gas-phase chemistry can also contributeto the observed ratios (§ 3.3 and below).Using a step-function model, Coutens et al. (2012, 2013a)also infer very high HDO / H O ratios of 0.05 for the outermostphotodesorption layers of the NGC 1333 IRAS4A and IRAS16293 -2422 envelopes. Averaged over the entire cold part ofthe envelopes, however, values of 0.002–0.02 are found (Coutenset al. 2013a). For the Serpens SMM1 source, narrow unsaturatedH
O and HDO absorption lines are detected (Fig. 23) imply-ing HDO / H O = + / H O = / H O ratios are much higher than the over-all [D] / [H] ratio of 2 × − in the gas, suggesting significantfractionation of water in cold cores, with the upper range valuesclose to those found of some other highly-deuterated molecules such as H CO or NH (Ceccarelli et al. 2014). It is importantto stress once more, however, that the observed cold HDO / H Ovalues do not reflect the bulk HDO / H O ice ratios, which aretypically (0 . − × − as measured in warm hot cores whereall ices have just sublimated (Persson et al. 2014; Jensen et al.2019).As discussed in § 3.3 and illustrated in Fig. 5, high levels ofcold water deuteration are well reproduced by various gas-grainchemical models (Cazaux et al. 2011; Aikawa et al. 2012; Taquetet al. 2013; Albertsson et al. 2014; Wakelam et al. 2014), espe-cially if the multilayer structure of the ice is taken into account(Taquet et al. 2014; Furuya et al. 2016). The bulk of the H O isformed early in the cloud evolution, whereas most of the deuter-ation takes place in the subsequent high density stage wherethe forward reaction producing H D + receives a boost from COfreeze-out. The outer layers of the ice are therefore expected tobe richer in deuterated water than the inner layers.Figure 24 presents models of H O, HDO and D O by Furuyaet al. (2016) that follow the chemistry along an infalling streamline from the cold outer part to the warm inner envelope (see alsoFig. 3 in that paper). No free atomic oxygen is needed in the highdensity phase to make deuterated water: the OH resulting fromphotodissociated water ice can be used to produce HDO. Themodels of Furuya et al. (2016) and Taquet et al. (2014) find thatcold gas-phase chemistry dominates over water ice photodesorp-tion in setting the cold gaseous HDO / H O value, and that bothresult in a roughly constant HDO / H O ratio with depth into theenvelope (Fig. 24, bottom, full lines at T <
100 K). The modelvalue of HDO / H O > . / H O ratio of 0.025 that has been observed for NGC1333IRAS4A.The main message is that the observed HDO / H O ratiosin cold gas can be understood using multilayer ice models butthat they do not reflect at all the bulk ice ratio. To obtain theHDO / H O ratio for the bulk of the ice, for example for compar-ison with cometary values, observations of hot cores are neededwhere the entire ice mantle is thermally desorbed above 100 K,as mentioned above and further discussed in § 8.Alternatively, the bulk HDO / H O ice ratio could be mea-sured by infrared absorption. The deepest direct observationallimits of infrared ice features give HDO / H O < Article number, page 24 of 58.F. van Dishoeck et al.: WISH: physics and chemistry from cloud to disks probed by
Herschel spectroscopy −
50 0 50 υ - υ source (km s − )012 T M B ( K ) Ser − SMM11 − H OHDOH O −
50 0 50 υ - υ source (km s − )0123456789 T M B ( K ) IRAS162931 − H OHDO × O −
50 0 50 υ - υ source (km s − )0123 T M B ( K ) NGC1333 − IRAS4A1 − H OHDO × O −
50 0 50 υ - υ source (km s − )0123 T M B ( K ) NGC1333 − IRAS4B1 − H OHDO × O Fig. 23.
HIFI observations of H O and H
O 1 −
557 and 548 GHz lines together with the HDO 1 −
893 GHz line toward a number oflow-mass protostars. t final − t (years) − − − − A bund a n ce r a ti o D O/HDO HDO/H O Temperature (K)10 − − − − − − n X / n H H OO atom
Fig. 24. H O, HDO and D O gas and ice abundances as function oftime following an infalling trajectory from cold outer envelope (left)to inner hot core (right) using the full multilayer chemistry by Furuyaet al. (2016). The top panel shows the water ice abundance (dotted line)in the cold outer part of the envelope, sublimating into gaseous water(solid line) when the parcel enters the ∼
100 K region. The bottom panelshows the computed HDO / H O (black) and D O / HDO (red) abundanceratios along the trajectory. Solid lines: gas phase, dotted lines: ice sur-face layers (top ∼
5% of ice); dashed lines: bulk ice (bottom ∼ O / HDO ice >> bulk HDO / H O ice, as also reflected in thewarm gas ratios. The cold gaseous HDO / H O ratio is very high in thismodel due to cold gas-phase chemistry; the observed HDO / H O valuefor IRAS4A is HDO / H O = the tentative HDO ice detections of Aikawa et al. (2012) at alevel of HDO / H O of a few %. Deeper limits may be obtainedwith JWST toward ice-rich sources although values down to10 − or lower, as expected for bulk ice, will remain di ffi cult toprobe due to limitations on feature-to-continuum ratios that canbe achieved with the instruments. abundances Standard gas-phase-only chemistry models of cold dark cloudspredict that most oxygen is transformed into O with time (e.g.,Goldsmith & Langer 1978; Millar et al. 1997). The low abun-dance of gaseous O in dense clouds measured by the S WAS (Goldsmith et al. 2000) and subsequently
Odin (Larsson et al.2007) satellites therefore initially came as a surprise. They werereadily explained by models that include gas-grain chemistrywhich turn most oxygen into water ice (Bergin et al. 2000), aprocess now also demonstrated to proceed in laboratory ice ex-periments at low temperatures (Ioppolo et al. 2008; Miyauchiet al. 2008; Oba et al. 2009).
Herschel -HIFI has been able topush the observational limits on O even deeper, with gaseousO remaining undetected in most sources. O has only firmlybeen identified through multiple transitions in the Orion shock(Goldsmith et al. 2011; Chen et al. 2014) and one region in theOph A cloud (Liseau et al. 2012; Larsson & Liseau 2017), butnot, for example, in the Orion PDR (Melnick et al. 2012). Itsabundance is only O / H = × − in the Oph A cloud.It is clear that the stories of the O and H O chemistries areintimately linked, and that any model that explains water gas andice also needs to be consistent with the O data (Fig. 4). Unfor-tunately, deep upper limits on the O
487 GHz line exist for onlyone low-mass protostellar envelope, the Class 0 source NGC1333 IRAS4A (Yıldız et al. 2013a). Using a drop-abundanceprofile, the inferred 3 σ limit is O / H < × − in the coldouter part, the lowest O limit so far. Even in the inner hot core,the observed O abundance cannot be more than 10 − , so O is not locked up in ice at high abundances and released at hightemperatures in this particular source. Recent ALMA interfer-ometer data also imply a similarly low gaseous O abundance inthe IRAS16293-2422 hot core (Taquet et al. 2018). Deep unpub-lished Herschel O
487 GHz data in a few high-mass protostarssuch as NGC 6334I and AFGL 2591 have also failed to detectany line at the 3 mK km s − (3 σ ) level, pointing to similarly lowO abundances as for IRAS4A.The low observed O abundance of IRAS4A has been mod-eled by Yıldız et al. (2013a) by coupling the temperature anddensity structure of the source with a full gas-grain chemicalmodel. Subsequent radiative transfer models provide line inten-sities that can be compared directly with the HIFI observations.The results indicate that only models with a long and cold pre-stellar stage of (0.7–1) Myr, coupled with a protostellar stage ofabout 10 yr, are consistent with the data. A long timescale atlow temperatures is needed to transform all O and O into waterice, leaving no free oxygen to form O . These same models arealso able to reproduce a very low O abundance in the hot core.The long timescale seems in contradiction with the conclusion Article number, page 25 of 58 & A proofs: manuscript no. wish_final_Jan2021 in § 7.1 of a short pre-stellar timescale to avoid the overproduc-tion of water ice. However, no water ice observations exist forIRAS4A to test this model because it is too faint at mid-infraredwavelengths. This conundrum of reproducing both water gas,water ice and O will be further discussed § 10.A new twist to the O puzzle has been provided by the de-tection of abundant O ice in comet 67P by the ROSINA in-strument on board the Rosetta mission (Bieler et al. 2015; Ru-bin et al. 2019). The inferred ratio is O / H O ≈ .
03 in the ice,which would correspond to O / H ≈ × − if released in thegas assuming H O / H = − . This is several orders of magnitudehigher than the limit for cold O measured for IRAS4A and closeto the warm O hot core limit (see also Taquet et al. 2018).A detailed parameter study by Taquet et al. (2016) shows thata O / H O ice abundance as high as a few % found for comet 67Pcan only be achieved if the atomic H / O ratio in the gas is low,so that water ice formation is suppressed and O ice productionpromoted. This, in turn, implies a rather specific narrow range ofphysical parameters: a relatively high temperature, ∼ > cm − , and low cosmic ray ionization rate, < − s − , for the pre-stellar cloud out of which our SolarSystem formed. A key result of these models is that they arealso consistent with the low observed HO , H O and O abun-dances in 67P, in contrast with alternative models such as waterice “radiolysis” (that is, processing of ice with ionizing photons)(Mousis et al. 2016). Similarly, other recent models that arguefor a primordial origin of the high O in comet 67P by forma-tion in a cold pre-stellar cloud (e.g., Rawlings et al. 2019) do notattempt to simultaneously reproduce these other species.Such a relatively warm model could also explain the detec-tion of O in the SM1 core of the ρ Oph A cloud, which isknown to have enhanced dust temperatures due to illuminationby nearby massive B-type star(s). A high gaseous O / H O abun-dance is only found at an early chemical age (Taquet et al. 2016)which may not be unrealistic given the measured infall speedsin this core (Larsson & Liseau 2017). For IRAS 4A, these samemodels are consistent with the low observed limits because thepre-stellar cloud is much colder, 10 K.
8. Hot cores: Dry or wet
In the warm inner envelopes close to the protostar, the dust tem-perature becomes higher than 100 K, at which point all waterice sublimates back into the gas. The size of this so-called “hotcore” region scales roughly as R T = ≈ . × ( √ L / L (cid:12) ) cmin a spherically symmetric envelope (Bisschop et al. 2007). For L =
1, 100 and 10 L (cid:12) , the 100 K radius is at 15, 150 and 1500au, respectively. Thus, whether for low-mass sources at 200 pcor high-mass sources at 2 kpc, the hot core region is < (cid:48)(cid:48) on thesky and thus heavily beam diluted in the Herschel beams. Theregion where T >
250 K, the temperature at which all volatileoxygen would be driven into water by gas-phase reactions, iseven smaller.Hot core water abundances have been derived by fitting theHIFI line intensities of the higher-lying water lines, especiallythose of H
O or H
O that are not a ff ected by outflow emis-sion (see Figure C.3 in Appendix). Detailed radiative transfermodels using a step-abundance profile have been performed formost of the WISH high-mass sources (Herpin et al. 2016; Choi2015) and one intermediate mass source (NGC 7129) (John-stone et al. 2010). The resulting inner abundances X in (with re-spect to H ) are included in Figure 21, together with a few re- sults from the literature using a similar approach: the low-masssource IRAS16293-2422 (Coutens et al. 2012) and the high-mass source NGC 6334I (Emprechtinger et al. 2013), both fromthe CHESS program. The narrow H O absorption lines towardhigh-mass sources also indicate a jump in water abundance of atleast an order of magnitude at the ∼
100 K radius (van der Taket al. 2019). For most low-mass sources, a step-function analysishas not been possible due to the lack of narrow HIFI H
O lines(Fig. 11).It is clear from Figure 21 that the inner water abundancesrange from 5 × − to > − , with no trend with luminosityor envelope mass, nor with the ratio L / M env , which is thoughtto be an evolutionary indicator. A few sources have inner abun-dances of 10 − or higher, as expected from ice sublimation andhigh temperature chemistry driving oxygen into water, althoughnone as high as the expected value of 4 × − . The only regionwith a water abundance that may be as high as 6 . × − is asmall compact dense clump near the Orion hot core found usingHEXOS data Neill et al. (2013). Among a subset of WISH high-mass sources, Herpin et al. (2016) have found a possible trendof higher inner water abundances with higher infall or expan-sion velocities. This could suggest that sputtering of ice mantlescontributes to the water production in the inner region.In an attempt to probe the inner warm water abundance forLM protostars with Herschel , very deep (5 hr) HIFI observa-tions of the H
O 3 − E up =
249 K) lines havebeen obtained for 6 sources (Visser et al. 2013). This H
O lineis detected in two low-mass sources, NGC 1333 IRAS2A andSerpens SMM1 (Fig. C.3 in Appendix), and actually shows nar-row profiles with FWHM ≈ − , so the emission is clearlynot associated with the outflow. Deep limits are obtained for theother sources (NGC 1333 IRAS4A, IRAS4B; GSS30, Elias 29);the H O line covered in the same setting is not detected at thesame noise level (8 mK in 0.5 km s − bin).A detailed combined analysis of the IRAS2A H O HIFI1095 GHz spectrum and H
O 3 −
203 GHz NOEMA datareveals that even the HIFI water isotopolog line is likely opti-cally thick when coming from a hot core with a radius of 100au. In contrast, the 203 GHz emission is optically thin. It resultsfrom the fact that the Einstein A coe ffi cient for the 1095 GHzline, 1 . × − s − , is a factor of 3400 larger than that of the 203GHz line, 4 . × − s − . Moreover, the dust continuum emissionfrom the hot core starts to become optically thick at 1 THz on100 au scales, thus shielding some of the water emission. Alto-gether, this means that only a lower limit can be placed on the hotcore water abundance of H O / CO > .
25 or H O / H > × − .Here the C O 10–9 ( E up =
290 K) line obtained in the samespectrum has been used as reference for the warm gas assumingCO / H = × − inferred for this source (Visser et al. 2013;Yıldız et al. 2012).Are most hot cores indeed “dry”, or is this an artifact of theanalysis method? For high-mass sources, the H O 203 GHz line( E up =
203 K) has long been recognized to be an excellent di-agnostic of water in hot cores, with inferred water abundancesof ∼ − − − (e.g., Jacq et al. 1988; van der Tak et al. 2006;Wang et al. 2012). Millimeter interferometers can now also im-age this line in low-mass hot cores with sub-arcsec resolution(Jørgensen & van Dishoeck 2010b; Persson et al. 2012; Jensenet al. 2019). These data show narrow water lines (FWHM ≈ − ) located within a ∼
100 au radius region. Inferred wa-ter abundances are low, only 10 − − × − (with respect toH ) if analyzed within a spherically symmetric model. However,on these scales much of the gas and dust is likely in a flattened Article number, page 26 of 58.F. van Dishoeck et al.: WISH: physics and chemistry from cloud to disks probed by
Herschel spectroscopy disk-like structure, even in the Class 0 phase (Murillo et al. 2013;Tobin et al. 2012, 2016). The temperature in the shielded mid-plane of the disk, where most of the mass is located, is muchlower than 100 K, so the amount of warm H in the denominatorcan be much smaller.Modeling the small scale interferometer structure with anonspherically symmetric model consisting of a disk (compactsource) + envelope indeed confirms that only a small fraction ofthe mass in the inner region is above 100 K (Persson et al. 2016).This raises the fractional water abundances with respect to H byan order of magnitude to 10 − − × − , with the latter valueapplying to IRAS 2A. These water abundances are closer to, butstill below, the expected water abundance of 4 × − .In summary, when corrections for nonspherical symmetryand optical depths are included, hot cores are not as “dry” as pre-viously thought, but they are not “wet” either, with water abun-dances still a factor of a few to two orders of magnitude lowerthan expected (see also Fig. 11 in Persson et al. 2016). Over-all, this comparison highlights the fact that Herschel was notwell suited to determine hot core water abundances, due to thecombination of severe beam dilution, high optical depths in lineand continuum, and di ffi culties in isolating the hot core emis-sion from that in the outflows. Studying the chemistry of hotcores is squarely an interferometer project, with NOEMA andnow also ALMA opening up the opportunity to image the 203GHz line, and other H O lines with small Einstein A coe ffi -cients and low continuum optical depth, for large numbers ofsources. The challenge remains to constrain the amount of warmgas on small scales to determine the water abundance, that is, thedenominator in the H O / H ratio remains the main uncertainty,rather than just the water column. This will require unravellingthe detailed disk-envelope physical structure on <
100 au scales.
There has been no shortage of determinations of the HDO / H Oratio in warm gas, using both ground-based and
Herschel -HIFIobservations. Single-dish ground-based determinations of high-mass protostars, using typically the 225 and 241 GHz HDO linesto measure HDO, typically find values as low as HDO / H O = (2 − × − (e.g., Jacq et al. 1990; Gensheimer et al. 1996; Helmichet al. 1996; van der Tak et al. 2006; Emprechtinger et al. 2013;Coutens et al. 2014), with occasionally values up to (2 − × − (Neill et al. 2013). Interferometer observations of high-mass pro-tostars are still sparse but indicate values of ∼ × − forthe disk-like structure of AFGL 2591 (Wang et al. 2012). Forlow-mass protostars in clustered regions, the interferometer datagive HDO / H O ratios in the range (0 . − × − (Jørgensen& van Dishoeck 2010a; Persson et al. 2014). Interestingly, forisolated low-mass protostars, the HDO / H O ratios are a factorof ∼ / H O ratios in warm gas are lower than those for thecold gas of 0.025 presented in § 7.3 by an order of magnitude ormore. This is consistent with the models of Taquet et al. (2014)and Furuya et al. (2016) discussed in that section and presentedin Fig. 24 (bottom panel, compare full lines at T < and > Class I Class II >100 au Region 3Region 1 Region 2 envelopedisk warm water cold waterhot water
Fig. 25.
Water vapor reservoirs in young Class I systems and matureClass II disks. Region 1 lies inside the water ice line where high temper-ature chemistry drives all oxygen into water near the midplane ( ∼ − abundance). Region 2 lies beyond the snow line and produces a smallamount of water vapor (10 − up to 10 − ) through photodesorption ofwater ice. Region 3 concerns the warm surface layers of the disk out tolarge radii with moderate water vapor abundances limited by UV radi-ation ( ∼ − ). Regions 1 and 3 extend further out into the disk for thewarmer Class I sources. Also, hot water can be emitted in the hot coreand in outflows or shocks along the outflow cavity in Class I sources.The small cartoons at the top illustrate the larger scale disk-envelope-outflow cavity system.
9. Protoplanetary disks
To address the question of the water trail from clouds to planets,the water content of disks at di ff erent evolutionary stages needsto be addressed. This section discusses Herschel ’s contributionto our knowledge of water in forming disks in the embeddedClass 0 / I protostellar stages as well as in mature disks in theoptically visible Class II stage.To put the
Herschel observations in context and set the stagefor the discussion, it is useful to summarize the expected watervapor reservoirs in disks based on chemical models (Fig. 25).Disks, whether young or mature, have large gradients in tem-peratures and densities, both radially and vertically. Their sur-face layers are exposed to intense UV radiation from the youngstar. Thus, the chemistry varies strongly with position in the disk,with the “snow” or ice-lines of di ff erent molecules playing a keyrole (Öberg et al. 2011b). Overall, models of protoplanetary diskchemistry have identified three di ff erent water chemistry regimes(e.g., Woitke et al. 2009; Gorti & Hollenbach 2008; Walsh et al.2012; Albertsson et al. 2014; Furuya et al. 2013; van Dishoecket al. 2014; Walsh et al. 2015). Region 1 lies inside the water iceline of ∼
160 K and has a high water abundance of at least 10 − ,with most oxygen contained in water vapor by ice sublimationand high temperature chemistry. Region 2 lies in the outer diskwell beyond the ice line, and has a low water vapor abundanceranging from 10 − up to ∼ − . This water vapor is producedat intermediate disk heights where UV radiation can penetrateand water ice can be photodesorbed from grains (Dominik et al.2005; Bergin et al. 2010a). Finally, in the warm surface layersof disks (region 3), water vapor is produced by high temperaturechemistry in the “warm finger” which has also abundant OH.There the water abundance is limited by photodissociation to avalue of ∼ − . The preceding sections have shown that most of the water emis-sion from low-mass protostars detected with
Herschel arises inwarm outflows and shocks associated with jets and winds and
Article number, page 27 of 58 & A proofs: manuscript no. wish_final_Jan2021 − . . . T M B ( K ) B335x0.3 L1527 RNO 91 x0.5 −
50 0 50 − . . . T M B ( K ) L1489x0.5 −
50 0 50 υ - υ source (km s − ) TMC1A −
50 0 50
DG Tau
Fig. 26. H O 1 − spectra toward low-mass embedded protostars with and without disk-like structures. RNO91, TMC1A and L1489 are ClassI sources with disks; L1527 is a Class 0 / I source with an embedded disk whereas B335 is a Class 0 source with no clear evidence for a disk downto 20 au. DG Tau is a young Class II source with disk and jet. These are WBS spectra binned to 0.3 km s − while B335 is binned to 0.6 km s − velocity resolution. The horizontal green line indicates the baseline. For TMC1A, the blue line indicates the ALMA CO spectrum of the rotatingdisk to indicate its width (Harsono et al. 2018). their interaction with the surrounding envelope, with the gasheated by kinetic energy dissipation. However, as the protostel-lar system evolves from the deeply embedded Class 0 phase tothe late Class I phase, the envelope mass and outflow force de-crease whereas the disk grows in mass and size (Hueso & Guil-lot 2005). Has
Herschel detected any water emission from theseyoung disks?Prior to the launch of
Herschel , the presence of youngdisks was heavily debated since numerical studies suggested thatstrong magnetic fields inhibit the formation of disks (Galli & Shu1993; Li et al. 2014). Recent millimeter interferometric obser-vations of optically thin molecular lines have however revealedflattened disk structures for several Class 0 and I sources tar-geted by WISH + that are characterized by Keplerian motion(e.g., Lommen et al. 2008; Tobin et al. 2012; Murillo et al. 2013;Harsono et al. 2014; Yen et al. 2017; Takakuwa et al. 2018; van’t Ho ff et al. 2018b; Artur de la Villarmois et al. 2019). Sincethese data clearly show that young disks are present, the disk’scontribution to the H O lines observed with
Herschel needs tobe reassessed.The sizes of the embedded disks are typically 50 to 100 au, so ∼ (cid:48)(cid:48) diameter at 140 pc (Harsono et al. 2014). Young disks in theClass 0 and I phase are expected to be warmer than their ClassII counterparts with similar mass due to their higher accretionrate and higher bolometric luminosity ( L (cid:38) (cid:12) , Harsono et al.2015). Thus, warm young disks may have su ffi cient amounts ofwater vapor that could contribute to the HIFI spectra. A first, purely observational look is to compare the H O 1 −
557 GHz spectra toward a number of Class 0 / I protostars withand without disks (Fig. 26). ALMA observations of the Class0 protostar B335 ( L = . (cid:12) ) show no clear evidence for aKeplerian disk inside of 20 au (Bjerkeli et al. 2019). In contrast,the Class I protostar L1489 with a similar luminosity ( L = . (cid:12) ) is surrounded by a very large rotating structure with a ∼ / I source (1 . − . (cid:12) ) with a ∼
100 au warm young disk seenedge-on (Tobin et al. 2012; van ’t Ho ff et al. 2018b). TMC1A(3.8 L (cid:12) ) is a Class I source with a bonafide disk with 50–100 au radius (Harsono et al. 2014, 2018). Finally, the young Class IIsource DG Tau is included since Podio et al. (2013) suggest thattheir Herschel -HIFI observations of the water ground-state linesat 557 GHz and 1113 GHz are emitted by the disk ( M disk = (cid:12) , L = (cid:12) ). DG Tau is an interesting case in that it also has apowerful optical jet which emits strong UV and X-rays onto thedisk surface (Güdel et al. 2010, 2018).At first glance, all of these profiles look similar. The B335spectrum stands out because of its broad line profile with aFWHM of 40 km s − typical of Class 0 sources while the waterlines toward the Class I protostars, when detected, have a typicalFWHM of 15–20 km s − (see § 4.3). Water is also detected to-ward the Class 0 L1527 source with a FWHM of 20 km s − butits comparatively low value could be due to the outflow lyingclose to the plane of the sky. No water emission is detected fromTMC1A, which does have a large disk as well as a blue-shifteddisk wind (Bjerkeli et al. 2016). Except for TMC1A, the inte-grated water line intensities follow the observed correlation of I (H O) with envelope mass M env and measured outflow forceestablished for a much larger sample (Kristensen et al. 2012;Mottram et al. 2017), but there is clearly no relation with L bol for this subsample. The former relation suggests that the lineprofiles largely reflect the warm gas associated with the outflow,which has weakened for the Class I sources. Association withthe outflow is further strengthened by the fact that spatially re-solved H O emission has been seen for several of these sourcesin the PACS data along the outflow direction, see Figures 4, D.1and D.2 in Karska et al. (2013) for L1527, L1489 and TMC1A.Could the “double peaked” profile reflect disk emission? Fig-ure 26 includes the observed CO 2-1 line profile observed withALMA from the inclined TMC1A disk and its disk wind show-ing a typical FWHM of 8 km s − . This is close to the maximumline width expected for outer disk rotation: even for mature disksaround Herbig stars ( M ∗ > (cid:12) ), CO line profiles do not spanmore than 10 km s − (Thi et al. 2001; Dent et al. 2005). The com-parison between L1489 and DG Tau shows that the DG Tau lineis narrower than most Class I sources, as expected from a disk,although cloud or wind absorption can also (asymmetrically) af-fect the profiles.Are the observed line intensities consistent with those froma disk? The simplest approach is to compare expected line inten-sities from optically thick water emission within the ice line. For Article number, page 28 of 58.F. van Dishoeck et al.: WISH: physics and chemistry from cloud to disks probed by
Herschel spectroscopy z ( a u ) a) log n ( H ) (cm ° ) 456789050100150200 z ( a u ) b) T (K) 204060801001201401601802000 100 200 300 R (au)050100150200 z ( a u ) c) log X ( H O ) ° ° ° ° ° ° ° ° ° Fig. 27.
Cross section of one quadrant of a model of a Herbig disk,tailored to that around HD 100546. (a) H number density, (b) gas tem-perature (taken equal to the dust temperature), and (c) calculated H Oabundance. Further details of the model can be found in the text andreferences. The H O abundance can be divided in three regions as indi-cated (see also Fig. 25). high accretion rates, the midplane temperature increases and thewater ice line shifts outward, from radii of only a few au in ma-ture, nonaccreting disks to values as high as 60 au for accretionrates of 10 − M (cid:12) yr − (Harsono et al. 2015). There could also bewater emission from the remnant warm inner envelope (§ 8). Themaximum water emission can then be computed by consideringoptically thick line emission within a ∼
100 au radius at 100 Kand then diluted by the
Herschel beam. This gives an integratedline intensity of ∼
20 mK km s − for the 557 GHz line for disks at150 pc, which is a factor of 5 below the upper limit for TMC1Aof ∼
100 mK km s − . For low accretion rates, the model in-tensities drop steeply with the reservoir of photodesorbed water(region 2) providing a minimum floor.In order to further quantify the young disk’s water contribu-tion to the observed spectra, a good physical model of the enve-lope + disk system inside of 100 au is required. This limits suchan analysis to two Class I sources, TMC1A and L1489. Herethe case of TMC1A is illustrated, for which the physical modelis taken from Harsono et al. (2018). In these models, the waterabundance in region 2 is assumed to vary from 10 − − − .Region 3 is assumed to have an abundance of 10 − and is takento lie at T gas >
250 K and A V between 0.5 and 2 mag. The UVradiation is typically G > A V < . − M (cid:12) yr − is included in the mid- plane. The predicted H O 1 − peak line temperatures fromsuch a model are low, ranging from 0.15 to 2 mK for the highestabundance in region 2, well below the noise level of 8–13 mKrms in 0.3 km s − bin in the Herschel spectra. In these models,most of the emission comes from region 3 ( T gas >
250 K) dueto the combined high excitation and column density that extendsup to 100 au. Region 2 only contributes if the water abundanceis as high as 10 − .In conclusion, to detect cold water vapor emission fromyoung disks with Herschel the observations would need to havebeen much deeper than the typical 15–30 min integration timesused in WISH + for Class I sources. For Class II disks, muchlonger integration times of >
10 hr have been adopted, but eventhen only few detections have been found (see § 9.2). Given thatClass 0 and I disks are overall warmer and more massive thantheir Class II counterparts with dust less evolved, they could havehigher water fluxes but future far-infrared space missions will beneeded to test this.
Exploratory searches for warm water in young disks have re-cently been undertaken with NOEMA and ALMA through the(nonmasing) H
O 3 , − , ( E up =
203 K) 203 GHz and4 , − , ( E up =
322 K) 390 GHz lines (e.g., Jørgensen & vanDishoeck 2010b; Persson et al. 2016; Harsono et al. 2020). Bytargeting H
O, any outflow emission is minimized, as evidencedby the narrow lines observed for Class 0 sources. In contrast withClass 0 sources, however, no H
O lines are detected for Class Isources. The current limits for four Class I sources suggest thatany warm water must be located inside of 10 au (region 1) (Har-sono et al. 2020). No ALMA 203 GHz data have yet been takentoward Class I sources, so future ALMA observations that are upto an order of magnitude more sensitive and higher angular res-olution than was possible with NOEMA are needed. For region3, the H
O column is too low to contribute, in contrast with thecase for the main water isotopolog discussed in § 9.1.1.The nondetections of warm water with ALMA are signifi-cant for other reasons as well. It is well known that the infall ofhigh velocity envelope material onto the low velocity disk cancause a shock, raising temperatures of gas and dust at the disk-envelope interface to values much higher than those providedby heating by stellar photons alone. These accretion shocks arewidely found in models and simulations (Neufeld & Hollenbach1994; Li et al. 2013) but have not yet been unambiguously ob-served. In the early stages of disk formation, the accretion occursclose to the star at high velocities, but that material ends up inthe star, not the disk. With time, the accretion quickly moves tolarger disk radii (tens of au) with typical shock speeds of < − impacting gas with densities > cm − . For these con-ditions, molecules in the infalling gas largely survive the shockbut gas temperatures just behind the shock reach up to 4000 K.These high temperatures drive most of the oxygen in OH andH O and should result in bright far-infrared lines. The dust tem-perature is also raised, but only to about 50 K (assuming 0.1 µ m grains), enough to release weakly bound molecules but notstrongly bound species such as H O (Visser et al. 2009; Miuraet al. 2017). Some sputtering of ice mantles can also occur. Thefact that no warm water vapor emission is detected with ALMAnear the young disk on 50 au scales for either L1527, Elias 29 orTMC1A (Harsono et al. 2020) limits the importance of any suchshocks.Interestingly, warm water emission in the H
O 10 , − , ( E up = Article number, page 29 of 58 & A proofs: manuscript no. wish_final_Jan2021 −
20 0 20 υ - υ source (km s − ) T M B ( m K ) H O - TW HyaHD100546HD163296AA Tau + MWC480+ LkCa15 + DM Taux5 −
20 0 20 υ - υ source (km s − ) H O − TW HyaHD100546HD163296
Fig. 28.
Overview of detected emission of H O 1 –1 (left) and 1 –0 (right) with Herschel -HIFI to TW Hya, HD 100456, HD 163296and the stacked results of AA Tau, MWC 480, LkCa 15, and DM Tau.Data from Hogerheijde et al. (2011); Du et al. (2017),
Herschel archiveand Hogerheijde et al. (unpublished).
SMA toward the Class I source HL Tau, known to have a largedisk structure seen in the continuum (100 au radius) with ALMA(ALMA Partnership et al. 2015). The nonmasing part of this lineis blueshifted by 20 km s − , has a FWHM of 25 km s − , andextends over 3 − (cid:48)(cid:48) , or ∼
500 au. These characteristics suggestthat the bulk of the water emission originates in the protostel-lar jet or wind, not in the rotating disk itself. Similarly, Watsonet al. (2007) found highly excited water lines in the
Spitzer mid-infrared spectra of the Class 0 source NGC 1333 IRAS4B whichthey attributed to an accretion shock onto the young disk. How-ever, Herczeg et al. (2012) concluded based on
Herschel- -PACSdata (taken as part of WISH), that both the mid- and far-infraredwater emission is consistent with an outflow origin.For high-mass protostars, bonafide Keplerian disks of sev-eral hundred au radius have been identified using the high spatialresolution of ALMA (e.g., Johnston et al. 2015; Moscadelli et al.2019; Izquierdo et al. 2018). The large beam of
Herschel makesit di ffi cult to isolate the water coming from young disks for theseobjects, so they will not be considered here. Also, these disks arewarm and are therefore better targets for ALMA. Indeed, vibra-tionally excited water lines ( v =
1, 5 , − , , E up = < ff erent environments. Herschel -HIFI was unique in probing the coldwater vapor reservoir in planet-forming disks, through velocity-resolved observations of the two ground-state rotational transi-tions of ortho-H O 1 − line at 557 GHz and the para-H O1 − line at 1113 GHz. With upper-level energies of 53and 61 K, and critical densities of ∼ × and ∼ × cm − , respectively, emission can be expected from across the disk. This includes the cold ( <
50 K) outer disk regions, al-though sub-thermal excitation conditions need to be taken intoaccount when interpreting the observations. It is these outer diskregions that will dominate the signal in the large
Herschel beamsof 38 (cid:48)(cid:48) and 19 (cid:48)(cid:48) at the respective line frequencies if the water va-por abundance is su ffi ciently high. Since the beam sizes exceedthose of disks (typically diameters of no more than a few arc-sec), the expected main-beam antenna temperatures are small.Although Herschel cannot spatially resolve the emission, HIFI’shigh spectral resolution ( < − ) allows the emission linesfrom suitably inclined disks to be velocity-resolved, thus pro-viding constraints on the radial origin of the emission if standardKepler orbital speeds are assumed.As discussed above, theoretical considerations predict wa-ter vapor to occur in three distinct regions in disks, as illustratedschematically in Figure 25 and in Figure 27 for one specific Her-big disk (e.g., Dominik et al. 2005; Woitke et al. 2009; Cleeveset al. 2014). The small radial extent of the region 1 results in anegligible contribution of the expected flux in the Herschel beamin the ground-state lines ( <
10% of total emission). In contrastwith young disks, region 3 also contributes little to the emissionin the ground-state lines in the
Herschel beam ( < density that precludese ffi cient excitation of the lines. Outside these two regions, wateris generally expected to be frozen out onto cold ( <
150 K) dustgrains, as evidenced by the handful of detections of water icein planet forming disks (e.g., Chiang et al. 2001; McClure et al.2015; Min et al. 2016; Honda et al. 2016). Region 2 consists ofthe cold water vapor reservoir produced by UV photodesorptionof water ice into the gas phase (e.g, Andersson & van Dishoeck2008; Öberg et al. 2009), with water vapor abundances of up to10 − relative to H predicted at intermediate disk scale heightsout to large radii. In this region, densities of a few times 10 cm − and temperatures of 30–45 K are su ffi cient to generate detectableemission in the Herschel beam.
HIFI results and models.
The main observational result from
Herschel on the ground-state rotational water lines from maturedisks, is that the emission is even weaker than expected on theo-retical grounds.
Herschel -HIFI observed the ortho-H O 1 –1 line toward ten disks around T Tauri stars and four disks aroundHerbig Ae stars as part of the WISH + programs (Table A.1). Asubset of these targets were also observed in the para-H O 1 –0 line with similar results and are not further discussed here.The observations and their results are presented in Hogerheijdeet al. (2011) and Hogerheijde et al. (unpublished), Salinas et al.(2016), and Du et al. (2017).In summary, out of the fourteen observed sources, only twoyielded detected lines (TW Hya and HD100546; Fig. 28) inspite of long integrations and rms levels down to 1.2–2.0 mKin 0.27 km s − channels. A stacked spectrum of AA Tau, DMTau, LkCa 15 and MWC 480 also yielded a significant detec-tion of the line. These detections correspond to the data with thelongest integration times (typically >
10 hr) and thus lowest noiselevels. For all other sources, observed to various rms levels, onlyupper limits were obtained. Figure 29 shows the detections andupper limits in terms of velocity-integrated line luminosities, us-ing the most up-to-date distances of the sources obtained fromGAIA.Figure 29 plots the observed line luminosities and the up-per limits as function of the gas and dust outer radii of the diskas measured in CO emission and through the mm-continuumemission. It also compares the observations to the expected emis-
Article number, page 30 of 58.F. van Dishoeck et al.: WISH: physics and chemistry from cloud to disks probed by
Herschel spectroscopy R CO (au) l og ( R T M B d υ a t c )( m K k m / s ) O 1 -1 R mm (au) l og ( R T M B d υ a t c )( m K k m / s ) O 1 -1 Fig. 29.
Line luminosities of models and observations of ortho-H O1 –1 for planet-forming disks observed with Herschel -HIFI. Solidblue and red lines show predicted line luminosities for a Herbig diskmodel (blue) and a T Tauri disk model (red); see text for details. Notethe logarithmic scale. The e ff ects of non-unity filling factors are illus-trated by shaded lines at 50% and 25%. The observational data arefrom Hogerheijde et al. (2011); Du et al. (2017) and Hogerheijde etal. (unpublished). Filled circles show detections, blue for the Herbigstar HD100546 and red for the T Tauri star TW Hya. The gray circlesconnected by the solid gray line show the detected line luminosity ofthe stacked detection of AA Tau, LkCa15, MWC 480 and DM Tau.Down-pointing triangles show upper limits (blue: Herbig disks; red: TTauri disks). Data are plotted versus the observed size of the disk asmeasured through CO emission (upper panel) and mm-wavelength con-tinuum emission (lower panel); for the stacked detection (gray line) therange of disk sizes is plotted. Observed disk sizes are taken from Walshet al. (2014); Andrews et al. (2012); Loomis et al. (2017); Bergin et al.(2016); Kudo et al. (2018); Jin et al. (2019); Qi et al. (2003); Liu et al.(2019); Fedele et al. (2017); Huang et al. (2016); Dutrey et al. (2003,2016); Hughes et al. (2009); Macías et al. (2018); Isella et al. (2016);Boehler et al. (2018); Huélamo et al. (2015); Cleeves et al. (2016). Lineluminosities and disk sizes are scaled to the latest GAIA distance esti-mates. sion for a “typical” T Tauri disk (red) and a “typical” Herbig Aedisk (blue). For the T Tauri disk, the TW Hya model of Hoger-heijde et al. (2011) and Salinas et al. (2016) is used: a disk massof 0.04 M (cid:12) , disk outer radius of 200 au, and a disk inclinationof 7 ◦ . For the Herbig disk, a model of the HD 100546 disk isadopted (Hogerheijde et al. unpublished): a disk mass of 0.01 M (cid:12) , a disk outer radius of 400 au, and a disk inclination of 42 ◦ .In both cases, the disk temperature is calculated self-consistently and the disk chemistry includes photodesorption due to stellarUV photons. The H O excitation and emission line luminosityis calculated using LIME version 1.9.5 e (Brinch & Hogerhei-jde 2010) and H O collisional rates from Dubernet et al. (2009)(taken from the LAMDA database Schöier et al. 2005). Whilethe T Tauri and Herbig disk models have a fixed outer radius, thee ff ect of confining the emission from water to a smaller regionis mimicked by artificially setting the water abundance to zerooutside a certain radius. Figure 29 plots the predicted emissionfrom these models against this radius. Predicted fluxes for Her-big disks are higher than for T Tauri disks because of the highertemperature of the former; due to the subthermal excitation ofthe water molecule, a factor of two di ff erence in temperaturestrongly a ff ects the excitation.Comparison of the model curves with the two detected disks(TW Hya and HD 100546) and many of the upper limits con-firms the earlier reported results, namely that the H O ground-state rotational emission from disks is weaker than expected byfactors of 5–10. Bergin et al. (2010a) and Hogerheijde et al.(2011) suggested that this indicates a reduction of water ice atintermediate disk heights where stellar ultraviolet photons canstill penetrate. These authors hypothesize that settling of (icy)grains can, over time, remove significant fractions of the waterice reservoir (see also Kama et al. 2016; Krijt et al. 2016, 2020).A similar overall reduction of volatile material was inferred byDu et al. (2017) for the full sample.Salinas et al. (2016) put forward a di ff erent scenario for TWHya, where the observed fluxes are reproduced when the icygrains are taken to spatially coincide only with the millimeter-sized grains as probed by the submm continuum emission. An-drews et al. (2012) find that in the TW Hya disk, these grainsare confined to ∼
60 au, much smaller than the 200 au foundfor the gaseous outer disk radius as seen in CO emission. Fig-ure 29 shows that much better agreement between the TW Hyadetection and the model curve is found when the former is plot-ted at the detected radius of the mm emission. Although in thiscase the model underpredicts the line luminosity of TW Hya, itshould be remembered that these model curves are produced bysetting the water abundance to zero outside the plotted radius,and do not include transporting the water ice from larger radii tosmaller radii. No exact agreement would therefore be expectedin this simple minded comparison. In contrast to TW Hya, forHD 100546 reducing the emitting area to that of the millimeter-sized grains ( ∼
265 au, Walsh et al. 2014) does not su ffi cientlyreduce the emission.Instead of reducing the size of the emitting region radially,introducing a non-unity filling factor can also produce a reduc-tion of the line luminosity. Figure 29 shows that scenarios whereas little as 25% of the disk out to the CO outer radius is filledwith H O emitting material, reproduces the observed line fluxesand many of the upper limits. Such filling factors may be appro-priate for several of the disks in the sample, if most of the wa-ter ice is spatially coincident with the mm-continuum emittinggrains. High-resolution ALMA studies show that such grains areconfined in many (bright) disks to rings or bands covering onlysmall fractions of the disks (e.g., Andrews et al. 2018; Long et al.2019; van der Marel et al. 2015). For the particular case of TWHya, the disk is found to be relatively uniformly filled with mm-continuum emitting grains out to ∼
60 au, while for HD100546the disk’s mm emission is dominated by a 15–45 au ring (Pinedaet al. 2019). While such a ring takes up only 1% of the total diskarea (out to 400 au), it does leave open the possibility that the e https: // github.com / lime-rt / lime Article number, page 31 of 58 & A proofs: manuscript no. wish_final_Jan2021 H O emission originates from only limited radial ranges insidethe disk.The three scenarios outlined above can all explain the weakH O emission found in the spatially unresolved
Herschel -HIFIobservations: (i) overall reduction of volatile material in the ver-tical regions of the disk subjected to UV radiation; (ii) a uniformradial confinement to the disk area rich in mm-continuum emit-ting grains; and (iii) H O emission regions carved up in rings orbands reminiscent of the resolved submm substructures seen inmany bright disks.For one disk, it is possible to go one step further. The de-tected spectrally resolved line toward the suitably inclined diskof HD 100546 (Fig. 28) holds information on the radial originof the emission. A simple model with a radial power-law de-pendency for the line emission suggests that the H O emissionoriginates from radii between 40 au out to 250–300 au (Hoger-heijde et al. unpublished). While the outer radius is not very welldetermined due to the adopted radial power-law drop o ff of theemission, the inner radius is robust and indicates that the emis-sion originates outside the bright ring of mm continuum emis-sion seen with ALMA (Pineda et al. 2019). Likely, this dustring is optically thick at the frequencies of the observed waterlines (557 and 1113 GHz), obscuring the water emission. Theinferred region of water vapor emission actually overlaps withthe 40–100 au region where water ice absorption has been de-tected (Honda et al. 2016). Summary and future prospects.
In summary, the
Herschel -HIFI observations of the water rotational ground state lines to-ward fourteen planet-forming disks show that water vapor ispresent in at least five of these for which the lowest noise levelswere obtained. Theoretical models suggest that this water vaporlikely originates from photodesorption of icy grains, but that thereservoir of available water ice is smaller than expected. Radialconfinement in a single region or several radial rings or bands,or vertical confinement at heights below where UV photons canpenetrate, can all explain this. Whichever mechanism is at work,a close relation with the settling, growth and radial transport ofdust grains seems implicated. The low gas-phase abundance ofcold water vapor is also in line with the inferred low gas-phaseoxygen and carbon abundances in the TW Hya outer disk (Favreet al. 2013; Du et al. 2015, 2017; Bergin et al. 2016; Bergneret al. 2019). If low oxygen and carbon abundances go together,then the low volatile carbon abundances inferred from weak COemission in many other mature disks (e.g., Miotello et al. 2017;Long et al. 2017) suggests that this is a common feature.The prospects of further investigating the distribution of coldwater vapor across disks requires spatially resolved observa-tions. ALMA observations of lines of water isotopologs providea possible avenue. However, available H
O transitions such asthe 203 GHz line (§ 8) have lower-level energies exceeding 200K, and do not trace the cold reservoir of water ice ( <
150 K).Several low-lying HDO transitions are observable with ALMAin Bands 8 and 10, but their interpretation is complicated by therequired knowledge of the deuteration fractionation.A highly promising avenue to constrain the radial distribu-tion of gas-phase water was pioneered by Zhang et al. (2013),who analyzed multiple transitions of water from HIFI and PACSin TW Hya to derive a radial abundance profile and snowline lo-cation by adopting a temperature structure for the disk. SPICAand the
Origins Space Telescope (25–588 µ m) o ff er the possibil-ity to extend such studies to a statistically significant sample ofdisks. .
930 2 . Wavelength ( µ m) N o r m a li ze d fl ux + o ff s e t OHH O Herbig AeAS205
65 66 67
Wavelength ( µ m)OH H OHerbig AeT Tauri
Fig. 30.
Water observations of protoplanetary disks at near-infraredwavelength with VLT-CRIRES and far-infrared wavelengths by
Herschel -PACS in typical T Tauri and Herbig disks. The VLT-CRIRESspectra are for individual disks, AS 205 and HD 250550, whereas thePACS spectra are obtained by stacking multiple spectra. Data fromFedele et al. (2011, 2013).
The
Spitzer -IRS detected a wealth of highly-excited pure ro-tational lines of warm water at 10–30 µ m in disks around asignificant fraction of T Tauri stars (Carr et al. 2004; Carr& Najita 2008; Salyk et al. 2008; Pontoppidan et al. 2010b;Salyk et al. 2011; Carr & Najita 2011; Salyk et al. 2015), withline profiles consistent with a disk origin (Pontoppidan et al.2010a; Salyk et al. 2019). Typical water excitation tempera-tures are T ex ≈
450 K. Spectrally resolved groundbased near-IRvibration-rotation lines around 3 µ m show that in some sourcesthe hot water originates in both a disk and a slow disk wind(Salyk et al. 2008; Mandell et al. 2012).Abundance ratios extracted from the Spitzer observations areuncertain because the lines are highly saturated and spectrallyunresolved. Nevertheless, within the more than an order of mag-nitude uncertainty, H O / CO ≈ −
10 has been inferred for emit-ting radii up to a few au (Salyk et al. 2011; Mandell et al. 2012).This indicates that the inner disks of those sources where water isprominently detected have high water abundances of order 10 − and are thus not dry, at least not in the surface layers down towhere the dust becomes optically thick at mid-IR wavelengths.In contrast with the case for young disks in the Class 0 and Istages, no NOEMA or ALMA observations of warm H O linesin Class II disks have yet been published. Emission of morehighly excited lines, however, has been detected using
Herschel -PACS (e.g., Fedele et al. 2012, 2013; Rivière-Marichalar et al.2012) and these transitions, when combined with longer wave-length HIFI and shorter wavelength
Spitzer / VLT / Keck infraredlines, provide insight into the gas-phase water abundance insidethe water snowline (Fig. 30).Interestingly, the HD 100546 Herbig disk, which has promi-nent HIFI lines pointing to cold water in the outer disk (Fig. 28),does not show any detection of warm water lines from the innerdisk with PACS or VLT. In contrast, the HD 163296 disk hasno detected HIFI lines, but does show a (stacked) detection ofwarm water with PACS (Fedele et al. 2013). These two stars andtheir luminosities are similar, so the global disk thermal structurecannot be the explanation. Instead, disk substructures in gas anddust likely are. One possibility is that water ice is trapped in theouter disk of HD 100546 by the bright dust ring at 40 au (Pinedaet al. 2019) (see discussion in § 9.2.1), whereas the icy pebbles
Article number, page 32 of 58.F. van Dishoeck et al.: WISH: physics and chemistry from cloud to disks probed by
Herschel spectroscopy have drifted inward for HD 163296. This would be consistentwith the high CO abundances detected inside the CO snowlinein the HD 163296 disk (e.g., Booth et al. 2019).Another illustrative case of the importance of dust traps inregulating the inner warm water reservoir is provided by theTW Hya disk, the only other disk for which HIFI ground-statecold water lines are clearly detected. Based on
Spitzer spectraof warm water combined with H , the inner few au of the TWHya disk are found to be oxygen poor by a factor of ∼
50. Thispoints to an e ffi cient water ice dust trap just outside the snowlinearound 2.4 au (Bosman & Banzatti 2019).More generally, Banzatti et al. (2017) find a correlation be-tween water line fluxes and size of an inner disk gas cavity,as measured from CO ro-vibrational line profiles for a set of TTauri and Herbig disks. Water emission first disappears at near-infrared wavelengths (hot water) and then at mid-infrared wave-lengths (warm water) as the radius of gas emission expands outto the water ice line. This suggests that infrared water spectraare a good tracer of inside-out water depletion in regions 1 and 3from within to outside the snow line. Banzatti et al. (2020) alsofind an anticorrelation between the water line fluxes in Spitzer data and the radius of the mm-sized dust disk, suggesting thatsmall disks where icy dust grains have drifted inward have higherwarm water abundances.
10. Water and the oxygen budget
Sections 5–9 have quantified the abundances of water vapor andother oxygen-containing molecules such as OH and O in dif-ferent components of protostellar sources: outflows and shocks,pre-stellar cores, cold and warm envelopes around protostars,and in young and mature disks. This section looks at the questionwhether we have now accounted for all of the oxygen.The broader issue of the oxygen budget from the di ff use ISMto dense clouds and comets was introduced in § 3.1 and is sum-marized in much more detail in Appendix G. The overview Ta-ble 3 and Figure 32 are included here in the main text of thepaper. As discussed in § 3.1, some fraction of oxygen appears tobe locked up in an unknown form called UDO (Unidentified De-pleted Oxygen) even in the di ff use ISM, assuming a total oxygenabundance of [O] = . × − or 575 ppm. The HIFI lines have allowed a detailed gaseous water abundanceprofile through the envelope to be derived (§ 7), but this chem-ical analysis ultimately has to be consistent with observationsof other oxygen carrying species, most notably water ice and O gas (§ 7.4). There are only a limited number of sources for whichboth water ice and gas have been observed along the same lineof sight, since most deeply embedded Class 0 sources were toofaint at mid-infrared wavelengths for ice absorption studies with Spitzer . In fact, there is no source for which all three species —water gas, water ice and deep O gas limits — are available.Nevertheless, surveys of water ice exist for large samples ofinfrared-bright low- and high-mass protostars, revealing typicalwater ice column density ratios N s − H2O / N H ∼ × − (e.g.,Gibb et al. 2004; Pontoppidan et al. 2004; Boogert et al. 2008;Öberg et al. 2011a; Whittet et al. 2013; Boogert et al. 2015).Similar values of (2 − × − are measured toward stars behinddark quiescent clouds (Boogert et al. 2013). The column densityof hydrogen nuclei, N H ≈ N (H ), is inferred either from the silicate optical depth, or from the color excess toward the star, orfrom a combination of the two. Observed values show a range ofa factor of two around this mean.The important implication of these water gas + ice obser-vations is that water gas contains a negligible fraction of oxy-gen and that even water ice locks up only a moderate fraction, (cid:46) . × − , that is, the oxygen that is notlocked up in silicates (§ 3.1). In other words, not all volatile oxy-gen observed in di ff use clouds ends up as water ice in the densecloud phase. A similar conclusion was reached by comparing gasand ice for high-mass protostellar envelopes by Boonman et al.(2003a). As shown in § 7.4, O contains at most a few %.There are four possible explanations for this conundrum: (i)water ice formation is suppressed, with most oxygen locked upin other volatile species; (ii) water ice or gas is destroyed afterformation, with oxygen driven into other species; (iii) the wa-ter ice abundance inferred from observations is underestimated;and (iv) the fraction of oxygen locked up in refractory UDO in-creases from di ff use to dense clouds. The following analysis fo-cuses on low-mass protostars, but the same arguments hold forhigh-mass protostellar envelopes. Regarding option (i), our modeling of both the water gas and iceusing the SWaN network demonstrates that the only way to geta low water ice column is to start with a low water ice abun-dance in the pre-stellar stage (Schmalzl et al. 2014). Water icecannot be e ff ectively destroyed in protostellar envelopes at tem-peratures lower than 100 K, but some additional water ice canbe made in the coldest outer parts where the dust temperature isbelow (cid:46)
15 K, which could increase the ice column by up to afactor of ∼ ∼
15 K is set by the binding energies of atomicand molecular oxygen, either of which is needed to make wa-ter ice. Once the temperature is above this critical temperature,which holds for the bulk of the protostellar envelope, no newwater ice can form and the water ice abundance is ‘frozen in’ atthe value given at the start of the protostellar phase, t =
0. TheSchmalzl et al. models use E b =
800 K for atomic O, but the morerecent higher values of E b = ff use cloud phase, water ice for-mation is fast (Cuppen & Herbst 2007). Current models thus re-quire a short pre-stellar phase, typically only 0.1 Myr at densities > cm − , to limit water ice formation (Schmalzl et al. 2014).This timescale is shorter than the estimated duration of ∼ and H O. Figure 31 presentsthe abundances in gas and ice of major oxygen-bearing speciesas functions of position in the IRAS4A envelope, together withtheir sum, using the three-phase chemical network of Furuyaet al. (2016, 2017). Initial abundances are identical to those inSchmalzl et al. (2014), with [O vol ] = . × − . The temperatureand density structure of the source is taken to be the same as thatused to analyze the H O and HDO spectra in § 7.1 (Fig. 22).After a short cold pre-stellar phase of 0.1 Myr, most oxygenis still in atomic form with an O abundance of 1.6 × − (Fig. 31,left panels). The water ice abundance varies with depth, but its Article number, page 33 of 58 & A proofs: manuscript no. wish_final_Jan2021 − − − − − A bund a n ce H O gasO O atomH O ice t pre = 0.1 Myr − − A V (mag) . . . F r ac ti ono f t o t a l O O atom+H O+O +CO+CO O atom + H O t pre = 10 Myr − − Fig. 31.
Full multilayer gas-grain model results for water and related chemical species for the temperature and density structure of the NGC 1333IRAS4A envelope model using the Furuya et al. (2016) chemical network, for two di ff erent timescales of the pre-stellar phase: 0.1 and 10 Myr.Top panel: gaseous O, O and H O (full lines) and H O ice (dashed lines). O is seen to be strongly reduced in the 10 Myr case. The bottom figureincludes the sum of the abundances of other major O-bearing species in the models in both gas + ice. The remaining oxygen is in H CO, CH OHand other larger organic species. integrated abundance along the line of sight is only 0.5 × − with respect to total H, consistent with the ice observations. Incontrast, for a long pre-stellar phase of 10 Myr (right panels),all oxygen is in water ice at the start of the protostellar phasewith no change with depth. Velocity resolved observations ofcold atomic [O I] in emission or absorption could in principledistinguish between these scenarios, but may be di ffi cult to in-terpret because of large optical depths.In what form is the remaining oxygen in these models? Asexpected, the O abundance is indeed strongly a ff ected by theduration of the pre-stellar phase: at 10 Myr, all the oxygen isin water ice with very little gaseous O , in contrast with the 0.1Myr case. As the bottom panels of Figure 31 show, some fractionis in CO and CO gas or ice. At most 20% of the volatile oxy-gen, about 60-70 ppm, is in other species in either model, mostlyH CO, CH OH and minor oxygen-containing species. All ofthese results are consistent with those presented in Schmalzlet al. (2014), which used a much simpler chemical model.The models presented in Fig. 31 assume a static physicalstructure with time. Figure 24 (top) shows the result of the waterchemistry in a collapse model such as described by Furuya et al.(2016) in one dimension and Visser et al. (2009) in two dimen-sions. There a streamline is followed from the outer edge of theenvelope inward, assuming a pre-stellar phase of 10 Myr. Con-sistent with the static case, the H O ice abundance stays higheven though di ff erent conditions are experienced “en route”.In conclusion, the only way to suppress H O ice formation isthrough a short pre-stellar stage.
For option (ii), episodic accretion and heating during the pro-tostellar phase has been proposed as a possible way to de-crease the water ice abundance. It is now well established thatmost protostars undergo multiple luminosity outbursts duringthe embedded phase of star formation, increasing the luminos-ity by up to a factor of 100 for a short period of time, typi-cally <
100 yr (e.g., Evans et al. 2009; Dunham & Vorobyov2012; Audard et al. 2014). The enhanced luminosity heats theenvelopes to higher temperatures, sublimating water ice out tolarger radii and re-freezing it once the envelope has cooled backdown again, but with some delay due to the fact that the freeze-out time is longer than the cooling time (Visser & Bergin 2012;Frimann et al. 2017). If water gas could be driven into otheroxygen-containing molecules during a luminosity burst beforere-freezing, this could suppress the water ice abundance com-pared with the pre-stellar stage and lessen the requirement on itsshort time scale. However, detailed chemical models have failedto make this scenario work in practice under protostellar condi-tions (e.g., Taquet et al. 2016; Eistrup & Walsh 2019).
Another solution (iii) is that the water ice absorption measure-ments underestimate the true water ice column. This could bethe case if some fraction of grains along the line of sight arelarge enough (typically a few µ m) that they do not show anyabsorption feature. This situation has been proposed for the dif-fuse ζ Oph cloud where also some fraction of oxygen is unac-counted for (Poteet et al. 2013). For our much denser protostellarsources, only a small fraction of the grains needs to be large. As
Article number, page 34 of 58.F. van Dishoeck et al.: WISH: physics and chemistry from cloud to disks probed by
Herschel spectroscopy
Table 3.
Oxygen budget of various species in ppm assuming overall [O] =
575 ppm
Form Species Di ff use clouds Low-mass YSO High-mass YSO Shocks CometsDust Refr. Dust 140 140 140 140 140Refr. Organics − − − − − H O < ≤
62 80 − CO 1 100 100 100 − CO − < < < − O-other − − − − −
Ice H O < − < − < − OH < − < − − See Appendix G for references to the adopted numbers in this table and their uncertainties. The ice abundances for protostars are asobserved, without hiding ices in large grains.a toy-model example, assume two dust populations: a standardone with grain size a s = µ m, and a large-grain population with a (cid:96) = µ m. The latter population would not show water ice ab-sorption even if covered with ice. The relative abundances areadjusted so that the overall grain abundance is 6.5 × − withrespect to hydrogen if all grains have a = . µ m (correspondingto a gas / dust mass ratio of ∼ f large . The equations in the SWaNnetwork by Schmalzl et al. (2014) have been adjusted for oxy-gen freeze-out onto the two di ff erent grain sizes. The fraction oftraceable water ice by infrared absorption is then taken to be theamount of ice on small grains relative to the total amount of wa-ter ice on both small and large grains. This analysis shows thata fraction f large = .
01% of the grain population in 10 µ m-sizedgrains (by numbers) would catch 50% of the atomic oxygen dur-ing freeze-out and make it invisible. Assuming a (cid:96) = µ m wouldmove the 50% point to f large = µ m size in denseclouds (e.g., Pagani et al. 2010; Boogert et al. 2013; Steinackeret al. 2015) and protostellar envelopes (e.g., Chiang et al. 2012;Miotello et al. 2014). These µ m-sized grains likely coagulatedfrom ice-coated small particles (since they carry the largest sur-face area), thereby natively locking up water ice within largergrains. This would have the added benefit that there does notneed to be a major fraction in atomic oxygen but there is stillroom for some fraction of UDO. For dark quiescent clouds priorto star formation, which also show low water ice abundances, theobserved low water ice abundances could be an early evolution-ary e ff ect with water ice still being formed, although even theregrain growth could start in the denser parts. The fourth possibility is that oxygen is locked up in increasingamounts in some refractory form called UDO as the density in-creases (Table 3) (Jenkins 2009; Whittet 2010). This scenario isfurther discussed in § 10.4.In summary, the combined observations of gaseous and solidwater in protostellar envelopes allow the water gas and ice abun-dances to be constrained quantitatively as functions of radius. The initial pre-stellar ice abundance is a crucial parameter thatsets the ice abundance profile in protostellar envelopes. Unlessthe pre-stellar stage is very short, it is di ffi cult to prevent allvolatile oxygen being turned into water ice. Alternatively, thepre-stellar stage could be longer, and the observed low ice abun-dances are then most easily explained by some fraction of grainshaving grown to at least µ m size. The final option is that thefraction of oxygen in UDO has increased from di ff use to denseclouds. The above models and discussion concerned the outer part ofthe protostellar envelope where water is frozen out. Section 8shows that the gaseous water abundance in the inner hot coreregion inside the water ice line is also generally lower than theexpected value of H O / H = × − . If most hot cores are in-deed dry, where is the remaining oxygen? Any water ice lockedup in large grains, as in option (iii), should also have sublimated.Photodissociation by UV radiation from the young star-disk ac-creting boundary layer (LM) or the stars themselves (HM) canrapidly destroy water, but only in a narrow region of visual ex-tinctions or in outflow cavities through which UV photons canescape (Stäuber et al. 2004; Visser et al. 2012). X-rays may bemore e ff ective in destroying water, but only if the gas tempera-ture is less than ∼
250 K (Stäuber et al. 2006, Notsu et al. unpub-lished). Once the temperature is higher, water persists thanks tothe rapid gas-phase reactions that drive oxygen back into waterin regions shielded from high energy radiation (see also Fig. 16).Which oxygen species are left as reservoirs? The lack of
Herschel -HIFI O detections from the hot core region in low-and high-mass protostars limits its contribution to less than afew %. CO gas is also known to be a minor contributor, atleast for high-mass protostars (van Dishoeck 1998; Boonmanet al. 2003b). This leaves atomic O as the major unknown in hotcores, even though models expect any atomic O to be rapidlytransformed into either H O or O under dense warm condi-tions. There are currently no observations that can constrain theamount of atomic O in hot cores. The only alternative would beto have a large fraction of oxygen locked up in UDO which isapparently not vaporized or atomized above 100 K. Article number, page 35 of 58 & A proofs: manuscript no. wish_final_Jan2021
Fig. 32.
Partitioning of oxygen in di ff erent forms of gas, ice and dust in (a) di ff use clouds; (b) protostellar envelopes (low- and high-mass); (c)shocks; and (d) comet 67P. The amount of oxygen in refractory dust is kept constant at 140 ppm in all cases. Section 5 and Table 2 have shown that the abundances of H Oin the di ff erent velocity components probed by Herschel are wellbelow the expected value of H O / H = × − . Even if enhanceddue to UV radiation, OH is also not a major oxygen reservoir(§ 5.2). Is the bulk of the oxygen in warm outflows and shocksin atomic oxygen, if not in H O or OH? Irradiated shock modelsindeed predict that atomic oxygen may be a primary reservoir, al-though published results are for pre-shock densities of only 10 cm − (Godard et al. 2019). There are numerous low resolutionspectra of the [O I] 63 and 145 µ m lines taken toward protostars(e.g., Green et al. 2013; Watson et al. 2016; Mottram et al. 2017),but those data are generally used to determine mass outflow ratesassuming an oxygen abundance with emission dominated by thejets (Nisini et al. 2015). Without spectrally and / or spatially re-solved data, no independent atomic oxygen abundance can bederived.Velocity-resolved observations of the [O I] 63 µ m line withthe GREAT instrument on SOFIA show that most of the [OI] emission in star-forming regions indeed originates in shocks(Leurini et al. 2015; Kristensen et al. 2017a; Gusdorf et al. 2017).However, for the NGC 1333 IRAS4A shock position R1, atomicoxygen accounts for only 15% of the oxygen budget in the highvelocity hot gas. CO appears to be the dominant oxygen carrierat an abundance of CO / H ≈ × − ; H O is a minor com-ponent, as is OH at this position (Kristensen et al. 2017a). Forthe distant high-mass source G5.89-0.39, Leurini et al. (2015)use spectrally resolved [O I] data to derive atomic oxygen col-umn densities which could be as large as those of CO, or evenlarger depending on assumptions, but no hydrogen columns areavailable to quantify abundances. H O and OH are again minorcomponents.Thus, even in high temperature (shocked) gas, some signif-icant fraction of oxygen may still be unaccounted for since it isalso not observed in [O I]. Only in the very hottest gas could theH O abundance approach the value required to contain most ofthe oxygen (Melnick et al. 2008; Neill et al. 2013). The impor-tant conclusion from this and related work (see Appendix G) is that this fraction of UDO must be locked up in a highly refrac-tory form (such as refractory organic material) that is not vapor-ized or atomized even for shock velocities up to 50 km s − andtemperatures up to several hundred K. Figure 32 summarizes the partionizing of oxygen between dust,gas and ice in the various regions discussed above. The adoptednumbers are summarized in Table 3 with detailed references andmotivation for choices presented in Appendix G. Uncertaintiesin each entry are di ffi cult to quantify and are discussed in theAppendix, but they are such that they do not change the overallpicture.Interestingly, the amount of unaccounted oxygen in the formof UDO is comparable between low-mass protostars, high-massprotostars and shocks at ∼
225 ppm, even though each individ-ual UDO number is uncertain by ∼
50% and very di ff erent tech-niques and instruments have been used in all three cases, witha mix of emission and absorption lines involved. It is roughlydouble the amount of UDO compared with di ff use clouds, con-sistent with the finding that the amount of UDO increases withdensity (Jenkins 2009; Whittet 2010). The absolute amount ofUDO depends in all cases strongly on the adopted overall [O]budget at 575 ppm; if the lower solar abundance of 490 ppm isused, the fraction of UDO may be halved but it does not disap-pear. In fact, Draine & Hensley (2020) argue for a higher ISM[O] abundance of 682 ppm with 66 ppm in UDO unaccountedfor in di ff use clouds.Table 3 and Figure 32 include the oxygen reservoirs incomet 67P / Churyumov-Gerasimenko, which has been studied inexquisite detail by the
Rosetta mission. Rubin et al. (2019) pro-vide an overview of the abundances in volatiles measured by theROSINA instrument, and those in refractory solids (dust) mea-sured primarily by the COSIMA instrument. This table assumesan ice:dust mass ratio of 1:1. By definition, there is no room forUDO in these measurements. However, some fraction of oxy-
Article number, page 36 of 58.F. van Dishoeck et al.: WISH: physics and chemistry from cloud to disks probed by
Herschel spectroscopy gen is found to be in refractory organic material, about 9% or 50ppm.There has been some speculation that UDO could consistmostly of refractory organic material that is unobserved in dif-fuse and dense clouds (Whittet 2010). The measurements forcomet 67P suggest that this component is not enough to explainall of the UDO in protostars and shocks. However, its amount isquite close to that found locked up in complex organic moleculesin large chemical models of protostellar envelopes (§ 10.1). Thissuggests that the chemistry may proceed from the more volatilesmall organic molecules to larger more refractory species as ma-terial is exposed to higher temperatures and UV radiation “enroute” from cloud to the comet-forming zones of disks (e.g.,Greenberg & Hage 1990; Drozdovskaya et al. 2016).For di ff use clouds, the measurement uncertainties are largeenough that refractory organics of the type seen in comet 67Pcould be the major UDO reservoir, especially if the di ff use cloudnumbers of Draine & Hensley (2020) are adopted.In summary, the observations of protostars and shocks dis-cussed in §5–9 indicate that a significant fraction of oxygen, oforder 100-200 ppm, is not accounted for. This section providespossible solutions to this conundrum. In all cases, some smallfraction of order 50 ppm may be in refractory organics, assumingthat this fraction is similar in the ISM as in comets. For cold pro-tostellar envelopes, an elegant solution for the remaining miss-ing oxygen is that ice columns are much larger than assumedbecause grains have grown to larger sizes suppressing the icefeatures. For hot cores and shocks, this solution does not worksince the water ice in large grains should have been sublimatedor sputtered so the situation is more puzzling. If gaseous atomicoxygen is indeed negligible, then the oxygen must be in somerefractory form that does not vaporize or atomize, even in strongshocks up to 1000 K.Two sets of observations are urgently needed. First, moremeasurements of velocity resolved atomic oxygen lines (both [OI] 63 and 145 µ m) in protostars and shocks can quantify its con-tribution in more sources. The current numbers hinge on just acouple of heterodyne [O I] spectra. Only the upGREAT instru-ment on SOFIA is capable of such measurements (Risacher et al.2018).Second, deep observations of mid-infrared lines of H O andH probing the hottest gas in shocks are warranted to determinewhether the UDO material is vaporized at T > O abundance starts to approach H O / H = × − . So far,only the data for NGC 2071 and Orion-KL point in this direc-tion (§ 5). This will be possible with the JWST-MIRI instrument(Wright et al. 2015) at shock positions o ff set from the protostarto avoid large extinctions. The simultaneous measurement of H and other atomic and molecular lines should provide better con-straints on shock density, which is one of the largest uncertaintiesin the analysis.
11. Discussion
Based on the detailed analyses for individual physical compo-nents and evolutionary stages, we can now come back to thethree questions that WISH aimed to address (§ 1), as well asthe two axes of WISH: from low- to high mass, including impli-cations for extragalactic water observations, and the evolution ofwater from cloud to disk. The focus is on
Herschel ’s contribu-tions to these questions. Finally some current and future oppor-tunities for observing water will be summarized together withlessons for future far-infrared missions.
Herschel observations, in particular the velocity-resolved HIFIdata, have allowed the main chemical processes leading to waterin star- and planet forming regions to be identified and quanti-fied (Fig. 33). Of the three routes illustrated in Fig. 4, the waterchemistry in star-forming regions is clearly controlled by a closecoupling between the gas and ice phases. Water ice builds up inthe cold dilute phase with some gaseous water vapor producedby ion-molecule reactions and destroyed by UV photodissocia-tion. In cold dense cores, the bulk of the water ice has formed,with a small fraction released into the gas by nonthermal pro-cesses such as photodesorption (see § 6 and 7 for more details).The new
Herschel observations have highlighted the importanceof cosmic-ray induced photodesorption deep inside clouds. Also,the multilayer nature of ice formation is key to interpreting dataon deuterated water.Once the protostar turns on, heating occurs both activelythrough shocks and winds associated with the outflows and pas-sively by warming up the dust and gas in the surrounding en-velope through the protostellar luminosity. Shocks can sputterices and produce copious water vapor through high-temperaturegas-phase reactions: the
Herschel data provide evidence thatboth processes are taking place with high temperature gas-phasechemistry dominating at the high velocities. Envelope heatingresults in water ice sublimation at its iceline around 100 K. Oncethe star has formed, it emits strong UV radiation setting up aPDR on the neighboring cloud, envelope or disk surface. X-raysmay also play a role in destroying water in the inner envelope andhot core. Figure 33 (left) summarizes the chemical processes,whereas § 11.3 discusses the abundances.
Table 1, Fig. 2 and 10 summarize the main water vapor linecharacteristics at each of the physical stages, for low and high-mass sources. Water vapor in cold low-density gas is primarilyseen in narrow ( ∼ − ) absorption lines against a brightfar-infrared continuum. Narrow emission lines are seen only incold high density quiescent gas such as those found in pre-stellarcores, in PDRs, and in the outer parts of protoplanetary disksnear the midplane. Relatively narrow water emission lines (a fewkm s − ) originate in warm, dense gas in protostellar envelopes(isotopologs, away from outflows).The water vapor emission seen by Herschel is, in contrast,mostly in broad lines associated with various types of kineticheating (cavity shocks, disk winds, turbulent entrainment) andmedium-broad o ff set lines (spot shocks). Regions of strong wa-ter vapor emission are generally compact, with an extent com-parable to or less than one Herschel beam ( (cid:46) (cid:48)(cid:48) , § 4.1). Evenwithin this beam, the emitting sizes are generally small, of order100 au for low-mass protostars.Three possible heating mechanisms are suggested for thebroad warm component involving kinetic energy dissipation(component 1, § 4.3, Fig. 34). First, cavity shocks heat and com-press the envelope along the cavity walls via non-dissociative,15–20 km s − , C -type shocks (Mottram et al. 2014). Second,dusty MHD disk winds (Yvart et al. 2016) launch dense gas fromthe inner disk at 0.2 out to 5–20 au, with heating taking place viaion-neutral drift in the accelerating flow. This results in a broadrange of outflow velocities, spanning a few to 80 km s − . Finally,turbulent entrainment heats gas within a mixing layer flowing Article number, page 37 of 58 & A proofs: manuscript no. wish_final_Jan2021
Cold, dilute • ice build-up • ion-molecule reactions Cold, dense • Photodesorption of ice
Warm, dense • thermal sublimation of ice • high- T neutral-neutral reactions Cold dense gas • prestellar core • protostellar envelope • disk midplane Warm dense gas • protostellar envelope • disk surface Outflows and shocks • cavity shocks • MHD disk wind • turbulent entrainment • spot shocks PDR • molecular cloud skin • outer envelope skin • outflow cavity wall • disk surface X gas (H O)10 -8 –10 – –7 –7 – –4 X gas (H O)10 –10 – –7 –6 –8 – 10 –7 –7 – 10 -4 Cold dilute gas • translucent clouds Chemistry 10 -8 Physics
Fig. 33.
Physical components dominating water emission with evolutionary stage (right) and principal chemical processes at each temperatureregime (left). Typical water abundances with respect to H associated with each evolutionary or temperature regime are indicated. between the wide-angle protostellar wind and the infalling enve-lope (Liang et al. 2020) and also provides a natural explanationfor broad line profiles peaking at systemic velocity. In addition,the mixing-layer model allows for a static outflow cavity thatdoes not expand beyond observed sizes over the typical Class0 phase duration. It is possible that all three mechanisms oper-ate simultaneously. Spatially resolved observations of the warmdense gas together with mid-infrared searches for shock tracers(e.g., atomic lines of Fe, S or Si) will be required to determinewhich dominates.Overall, the conclusion from the Herschel data is that wa-ter emission is a highly sensitive probe of the physics of dif-ferent star-formation phases with unique absorption or emis-sion line characteristics that are not easily seen by any othermolecule. Viewed from a more global perspective: strong wa-ter emission with large line luminosities clearly points to activestar-formation sites.
Herschel has also provided insight into the physical componentstraced by other molecules and their lines, with a summary pro-vided in Figure 34 (see also Fig. 8 in San José-García et al.2016 and Fig. 5 in Mottram et al. 2014). High- J CO lines with J up ≥
14 trace the same regions as water. In contrast, the low- J CO lines probe the entrained outflow gas physically separatedfrom the actively heated and shocked material.OH shows broad line profiles consistent with those of H Oand thus probes the same components (§ 5.2). Its abundance isenhanced because the warm gas is exposed to UV radiation dis-sociating H O into OH and changing the overall shock struc-ture. Hydrides such as OH + , H O + and CH + have narrower lines (but not as narrow as those of CH) and originate primarily in theUV irradiated outflow cavity walls and spot shocks (Benz et al.2016).Other molecules that are abundant in ices, such as NH andCH OH, mostly trace the hot core region and the low-velocitypart of the shock. HCO + , whose abundance anticorrelates withthat of H O, largely avoids the outflow region. SiO can trace thefast EHV part of the shock similar to H O, but its abundance andline profiles show a more rapid evolution across sources than thatof H O. Water lines are a significant coolant of gas in warm gas, butnot dominant. For low-mass Class 0 and I sources, water con-tributes typically at the 20% level, whereas CO does so at 30-45% of the total far-infrared line cooling. The remainder of thefar-infrared cooling budget is provided by OH and [O I]. Forhigh mass sources, CO becomes more important, up to 70% oftotal, whereas the contribution from water lines drops becausemore water lines occur in absorption rather than emission.Significant gas cooling is also provided by the H mid-infrared lines (Maret et al. 2009). These are not considered herebut will be further quantified by future JWST observations. Figure 33 summarizes the typical water vapor abundances in-ferred for each of the physical components. Only average orderof magnitude values for each type of source are given, with de-tails described in §5–9. For individual sources, abundances areaccurate to typically a factor of a few or better.
Article number, page 38 of 58.F. van Dishoeck et al.: WISH: physics and chemistry from cloud to disks probed by
Herschel spectroscopy
Fig. 34.
Cartoon illustrating in which physical components most of the H O, CO and other molecular line emission originates for the di ff erentscenarios presented in the text. These include the outflow cavity shocks and spot shocks (top left), the disk wind (top right), the turbulent entrain-ment (bottom left), or a combination of the three (bottom right). Each panel shows the jet with bullets (green), the outflow cavity wall (red / yellowgradient), disk and envelope (gray-scale). Arrows are meant to indicate gas motion. This figure illustrates that the gaseous water abundancevaries significantly and is thus a good diagnostic of di ff erentphysical regimes. However, it never reaches the maximum abun-dance of H O / H ≈ × − expected if all volatile oxygenis driven into water. Only the warmest gas has abundances ap-proaching 10 − , but in most cases the gas components probed by Herschel have abundances of 10 − − − with respect to hydro-gen.In cold gas, water is mostly frozen out as ice so water vaporabundances are very low, down to 10 − . Interestingly, the mostaccurate constraints on water abundance profiles have been madefor cold cores or outer protostellar envelopes, where the veloc-ity profile permits reconstruction of the water abundance as afunction of position (see § 7.1). At the edge of the core, H Ogas is produced by photodesorption of water ice at an abundanceof ∼ − ; deeper into the core it drops to less than 10 − . Thesame holds for protoplanetary disks, but then in the vertical di-rection: water vapor is UV photodesorbed in layers at intermedi-ate heights, dropping to very low abundances near the midplane. Section 4, together with Figures 2 and 10 and Table 1, demon-strate the large similarity in water line profiles from low- to high-mass YSOs. As summarized in San José-García et al. (2016), theprofiles can be decomposed in similar physical components as-sociated with kinetically-heated gas (warm outflows and shocks)and warm dense envelopes, but with the relative contributionfrom the envelope increasing from low- to high mass (see Ap-pendix C). The same physical components can also be seen inhigh- J CO lines, but only for J up >
14. Even then, they are lessprominent than in water lines. The similarity in profiles meansthat common physical mechanisms are at work in the outflows ofsources across more than five orders of magnitude in luminosity and that the gas cooling structure appears universal. The waterlines probed by
Herschel and the high- J CO emission are al-ways dominated by material at similar conditions: temperaturesof a few hundred K and densities from 10 − cm − .The water line luminosities show a strong, near-linear cor-relation with bolometric luminosities of the sources, L bol . Foroptically thick, but e ff ectively optically thin lines such as thoseof water, this empirical relation simply implies that there is moregas at the required conditions for more luminous sources: mate-rial is heated further from the source and deeper into the cav-ity wall. This is consistent with the larger emitting areas foundfor high mass sources from the water radiative transfer analysiscompared with that for their lower mass counterparts (few thou-sand au versus 100 au).Overall, the intensity ratios of the various water lines probedby Herschel -HIFI are also very similar across the luminosityrange (and even as function of velocity, modulo absorption fea-tures), pointing to the similar physical conditions inferred above.Surprisingly, however, one line ratio stands out: the 752 GHz2 − /
988 GHz 2 − ratio is generally less than 1 for awide range of sources, but lies above 1 for high-mass protostars(San José-García et al. 2016). One possible explanation is thatradiative pumping plays a more significant role for higher-massprotostars than for their lower- and intermediate-mass counter-parts. Alternatively, absorption in the 988 GHz line a ff ects theline ratio in high-mass sources (van der Tak et al. 2019). In manyaspects, the intermediate mass sources look more similar to low-mass protostars, or an unresolved cluster of low-mass protostarswithin the Herschel beam, than a scaled-down version of high-mass sources.The water abundances in the various physical componentsare also found to be similar across the luminosity range. Innerand outer envelope water abundances, when analyzed in a simi-lar manner, show a comparable range of values (Fig. 21). Water
Article number, page 39 of 58 & A proofs: manuscript no. wish_final_Jan2021 − L FIR ( L (cid:12) )10 − − L ( H O − )( L (cid:12) ) Low-massHigh-massExtragalacticHigh- z SMG
Fig. 35.
Line luminosity of the H O 2 —1 transition at 988 GHz ver-sus the FIR luminosity for low- and high-mass sources and extragalacticsources. The gray line is the best-fit power-law function, with an indexof 0.95 ± abundances in warm outflowing gas are low by at least two or-ders of magnitude due to UV irradiation in all sources, althoughthere is some debate to what degree shocks associated with high-mass sources have even lower abundances than their low-masscounterparts due to more intense UV (Table 2). The inferred linear relations of line luminosity with bolometricluminosity for Galactic protostars hold over many more ordersof magnitude when observations of extragalactic sources fromYang et al. (2013) are included, both for water and for mid- tohigh- J CO lines (San José-García et al. 2013, 2016). Figure 35shows an updated version of this correlation for the 988 GHz lineincluding the low-mass sources from the WILL survey (Mot-tram et al. 2017) and low- to intermediate-mass sources in theCygnus star-forming cloud (San José-García 2015). This strongrelation with FIR emission has been interpreted di ff erently bythe two communities, however (Kristensen & Bergin, subm.). InGalactic sources, it is well known from the Herschel data thatwater and high- J CO are associated with shocked gas where thelines are collisionally excited (e.g., Mottram et al. 2014). Theemission is spatially compact, not extended over entire molec-ular cloud scales (§ 4.1). In fact, away from the immediate sur-roundings of protostars and their outflows, the water emissionand abundance drops steeply, within a single
Herschel beamfor high-mass sources ( (cid:46) ff ect the water excitation: there is no direct causal relation.In contrast, the extragalactic community uses the tight cor-relation between the H O line luminosity and FIR luminosity toinfer that the water excitation is dominated by FIR pumping, asopposed to collisional excitation in outflows. In the extragalacticscenario, water emission thus originates in the entire molecu-lar cloud in which stars form, and is not directly associated andco-located with star formation (e.g., van der Werf et al. 2011;González-Alfonso et al. 2014; Yang et al. 2016; Liu et al. 2017;Jarugula et al. 2019). As shown by Kristensen & Bergin (subm.), the low- and mid- J water line ratios ( E up <
300 K) do not change significantlybetween Galactic and extragalactic sources. Most notably, thepeculiar 752 /
988 GHz line ratio found for high-mass sources isnot found for extragalactic sources (Yang et al. 2013; San José-García et al. 2016). Thus, the baseline interpretation is that thesame mechanism that is responsible for water emission in Galac-tic sources also holds for extragalactic sources: water emissionarises in outflows and shocks associated with (clusters of) cur-rently forming stars. Thus, any water line that does not su ff ersignificant self-absorption is a good and clean tracer of buriedstar formation activity.By inference, the same interpretation would hold for mid- tohigh- J CO lines. For high J up ≥
14, there is an observed closerelation between CO and water lines with both tracing currentlyactive outflow heating (§ 4.1 and 4.3). For mid- J CO lines therecan also be a contribution from the warm envelope. For example,for CO J = −
9, the fraction of envelope emission is around50% for high-mass sources (San José-García et al. 2016). Sinceenvelope emission also traces current star formation, it does notnecessarily a ff ect the interpretation, albeit that there could be ex-tra scatter due to the variation in outflow versus envelope contri-bution for each source.Is all water emission associated with current star formationactivity? A subset of extragalactic sources are known to driveoutflows on much larger galaxy scales, often seen perpendicu-lar to the galactic disk where they can escape freely. The jetsand flows are thought to be driven by Active Galactive Nuclei(AGN) buried in the center of the galaxy (see Blandford et al.2019, for a review). These large galaxy-scale outflows are welltraced by OH lines, not by H O lines, observed by
Herschel -PACS (Sturm et al. 2011; González-Alfonso et al. 2017). OH ismuch more abundant than H O due to the intense UV radiationin the outflows and cavity walls away from the shielded regions.Thus, H O is the better probe of star formation activity whereasOH may be the better probe of AGN activity.A big advantage of extragalactic observations is that evenfor small redshifts, the water lines shift into frequency rangesthat are not obscured by the Earth’s atmosphere and can be wellobserved from the ground. ALMA is already providing a richdatabase on water lines in extragalactic sources and will continueto do so in the coming decades. The
Herschel mission studyingwater in Galactic star-forming regions was needed to properlyinterpret these ALMA observations, thereby revealing the phys-ical origin of emission. The next step will be to calibrate wateremission with measured star-formation rates on galactic scales,such that water emission can be used to quantify star-formationrates in distant high-redshift galaxies.
The third goal of the WISH + program was to follow the trail ofwater from its formation in molecular clouds to the sites wherenew planets are formed. Addressing this question requires obser-vations of both water vapor and water ice using a wide range ofinstruments. The broader picture of water from clouds to plan-ets is reviewed elsewhere (van Dishoeck et al. 2014; Hartmannet al. 2017; Morbidelli et al. 2018). Here the unique contribu-tions that Herschel has made are highlighted, in concert withrecent millimeter interferometry (ALMA, NOEMA) results forwarm water.
Article number, page 40 of 58.F. van Dishoeck et al.: WISH: physics and chemistry from cloud to disks probed by
Herschel spectroscopy
Herschel -HIFI data of cold clouds and outer protostellar en-velopes coupled with detailed chemistry and radiative transfermodeling have demonstrated that most water is formed as ice inthe early stages of cloud evolution (§ 6 and 7). This is consistentwith direct observations of water ice at infrared wavelengths inquiescent and star-forming clouds.
Herschel studies, most no-tably those as part of WISH, have demonstrated a good quantita-tive understanding of the cold water chemistry and ice formationin the pre-stellar stage. The amount of water ice that can be pro-duced at subsequent warmer protostellar stages is only a smallfraction of the total ice.It is likely that a significant fraction of available volatile oxy-gen (i.e, not locked up in silicates) is contained in water ice incold clouds. The fact that not all of this ice is detected in infraredabsorption spectra either points to a short pre-stellar stage, or toa small fraction of grains that have grown to at least microm-eter size by the dense cloud stage. Some ill-determined frac-tion of oxygen is contained in other forms, either in volatileO-containing species or more refractory material called UDO(§ 10). While the UDO fraction may change with evolutionarystage, it is not expected to a ff ect the bulk of the water trail.How much water ice is made in dense clouds? A typical 1 M (cid:12) cloud forms about 6 × water molecules assuming a water iceabundance of 10 − with respect to H , over a period of about 10 yr. This amounts to more than 1 million oceans of water availablefor new planetary systems that are formed in the cloud (1 Earthocean = gr of water = × water molecules). Thislarge number of “oceans” implies that a significant water reser-voir is available to be supplied. However, there are a number ofoverall loss factors, both in the supply to the natal disk and thento a nascent planet (see below). Thus this number should not betaken to imply a direct link to Earth’s water (e.g., van Dishoecket al. 2014). Herschel data have confirmed the rapid production of water inwarm kinetically heated gas associated with protostellar out-flows and winds and their impact on the surrounding envelope.Using the observed mass outflow rates and water abundances, oforder 1–5 oceans of water vapor per year can be produced in asingle low-mass outflow over the period of the lifetime of theembedded phase, which is about 10 yr (Kristensen & Dunham2018). Thus, the amount of water molecules formed in outflowsper year is comparable to that in cold clouds. However, all ofthis water is carried away by the outflows and thus lost to space.There water is either photodissociated or frozen out back ontothe grains as ice (Bergin et al. 1998). The transport and survival of water ice from the collapsing coreonto the forming disk is a crucial step in the water trail. Howand where material falls into the disk and with what speed isstill poorly understood and ill constrained observationally. Di-rectly linked to this question is the extent to which the chemicalcomposition is preserved from cloud to disk (“inheritance”), orwhether it will be modified en route (“reset”) due to the changingtemperatures and UV radiation along the infall paths (e.g., Visseret al. 2009; Hincelin et al. 2013; Drozdovskaya et al. 2016). Inthe most extreme case of strong accretion shocks, reset implies complete vaporization of all molecules and dust grains back toatoms, with subsequent reformation; milder versions of reset in-clude sputtering of ices, and gas and ice chemistry modifyingabundances.To test the inheritance versus reset scenarios, key observa-tional diagnostics of these processes are needed. There is someobservational evidence for chemical changes near the centrifugalbarrier in the form of strong SO and SO emission (e.g., Sakaiet al. 2014; Artur de la Villarmois et al. 2019), but implied tem-peratures are low, ∼
50 K, and their origin and implications forany shock are still unclear.The results for young disks found in § 9.1 argue for sig-nificant inheritance of water ice since no water vapor, warm orcold, is observed to be associated with young disks in the Class Iphase, as would be expected when accretion shocks onto the diskwould sputter the ices or produce abundant water gas. Also, awater-vapor rich warm envelope is ruled out by the current
Her-schel and ALMA + NOEMA data, although deeper and higherspatial resolution ALMA data could put up to an order of mag-nitude stronger constraints on any warm water produced by ac-cretion shocks. Even the disks themselves are cold enough bythe Class I stage that they have very low water vapor abundancesin the outer disk and surface layers, suggesting that the bulk ofthe water is locked up as ices onto grains at an early stage of diskformation.Models by Visser et al. (2009) and Furuya et al. (2017) thatfollow the water chemistry along many infall trajectories havefound that most water in disks is indeed inherited as ice. Asmall fraction may be sublimated during a luminosity burst ofthe young (proto)star moving the water ice-line outward, as of-ten happens in the Class 0 and I phases, but this water ice quicklyrefreezes again without chemical alteration (Taquet et al. 2016).Only trajectories that come close to the outflow cavity wall resultin water dissociation due to exposure to UV radiation or X-rays.
Figure 25 provides a useful framework to illustrate the di ff er-ent water reservoirs in disks. Herschel has shown that gas anddust evolution as well as dynamics can alter this simple picture,as discussed in § 9.2. Once in the disk, the icy grains can growfurther and settle to the midplane (Krijt et al. 2016). There isobservational evidence based on the spectral index at mm wave-lengths, as well the required dust opacity for hiding lines in theinner disk, that grains have already grown to mm or cm size indisks during the embedded phase (e.g., Kwon et al. 2009; Testiet al. 2014; Harsono et al. 2018). These large icy grains will driftinward or be trapped in pressure bumps where they can grow fur-ther to water-rich planetesimals (e.g., Raymond & Izidoro 2017;Schoonenberg & Ormel 2017). The
Herschel -HIFI data of disksindicate that water vapor likely follows the large icy dust grainsthat have settled and drifted inward.Once the drifting icy pebbles cross the ice line, they will re-lease their water into the gas resulting in an inner disk that isrich in water vapor, with abundances that can even exceed thecanonical value of 10 − by more than two orders of magnitudewithin 1 Myr (Ciesla & Cuzzi 2006; Bosman et al. 2018a). Ac-cretion onto the star will however remove this oxygen-rich gasand there is some disagreement on the relative importance of re-plenishment versus removal. If removal dominates, it may drainthe inner disk of water-rich gas in as little as 0.1 Myr (Ali-Dibet al. 2014). Tracing the water vapor in this inner disk requiresmid-infrared observations. Spitzer data have indeed provided ev-idence for both high and low water abundances in the inner disk
Article number, page 41 of 58 & A proofs: manuscript no. wish_final_Jan2021
Fig. 36.
HDO / H O (lower points) and D O / HDO (two top points)abundance ratios measured for warm gas in low-mass protostellar re-gions and in comets (JFC = Jupiter Family Comets; OCC = Oort CloudComets). The HDO / H O data are taken from the summary in Jensenet al. (2019) (brown: isolated protostars; blue: protostars in clusteredregions), the D O / HDO star-forming cloud ratio from Coutens et al.(2014) and the HDO / H O and D O / HDO ratios in comet 67P (JFC highpoints) from Altwegg et al. (2017). The much higher HDO / H O ratiosfound in the top layers of multilayer ice models are indicated with thelight green zone and are similar to the HDO / H O values of 0.025 mea-sured with
Herschel -HIFI in cold outer envelope gas. linked to dust disk size and structures (e.g., Banzatti et al. 2017,2020) (see § 9.2.2).What do the HIFI data imply for the availability of water innew planetary systems? The HIFI detection of gaseous H O inthe TW Hya disk implies a reservoir of water ice of a few thou-sand oceans, which suggests plenty of water to build a habitableplanet such as Earth. However, it is unclear whether all of thiswater can reach the habitable zone close to the star. The forma-tion of a single Jupiter-like planet in a disk could lock up thebulk of this available oxygen reservoir. Also the flow of icy peb-bles from outer to inner disk can be stopped at a dust trap outsidethe ice line. Indeed, the inner TW Hya disk is known to be poorin oxygen and carbon, likely due to a dust trap outside of 20 aulocking up CO and one at 2.4 au locking up H O (Zhang et al.2013; Schwarz et al. 2016; Bosman & Banzatti 2019; McClureet al. 2020) (see also § 9.2.2). In such cases, scattering of ice-rich planetesimals from the outer to the inner disk could happenonly at later evolutionary stages when the gas has disappeared,and in that way deliver water to forming terrestrial planets in thehabitable zone (Morbidelli et al. 2012).
An alternative way to probe water ice processing en route fromcloud to disk is ice deuteration, which is particularly sensitiveto physical parameters (Ceccarelli et al. 2014): if most water iceis inherited, the HDO / H O ice ratio set in the cold pre-stellarstage should be preserved in subsequent stages. The situation ishowever not as simple as initially thought because the layeredstructure of ices needs to be taken into account. Also, HDO hasnot yet been detected in mature Class II disks.Ice formation models that include a multilayer structure haveshown two stages: the initial cloud formation stage in which most of the H O ice is formed, and the later dense core phasein which the bulk of the HDO and D O ice are made, with onlylittle additional H O ice (Taquet et al. 2014; Furuya et al. 2016)(Fig. 5). A typical grain of ∼ µ m has nearly 100 ice layersif all oxygen is converted to ices. The bottom 75–95% of thoselayers have HDO / H O around 10 − whereas the top 25–5% haveHDO / H O as high as 10 − − − (Fig. 24). This results in over-all bulk HDO / H O ice ratios of ∼ − within a factor of two(Taquet et al. 2014; Furuya et al. 2016). A characteristic featureof this cold, layered ice chemistry is that doubly deuterated wateris even more enhanced: the bulk D O / HDO >> bulk HDO / H Oby about a factor of 10 (Fig. 24).What do the
Herschel and new NOEMA + ALMA mm inter-ferometry data tell us? Figure 36 summarizes the observationalresults for protostellar envelopes compared with comets, usingdata summarized in Jensen et al. (2019). First, the derived highgaseous HDO / H O of 0.025 in the cold outer core (§ 7.3) fromthe HIFI data is entirely consistent with the multilayer ice for-mation (green shaded area), and is much higher than the valuesfound in hot cores. The high value reflects both the fact that onlythe top, deuterium-enriched ice layers are photodesorbed, withlow-temperature ion-molecule chemistry further enhancing thedeuteration (Fig. 24) (Furuya et al. 2016). Also, Coutens et al.(2013b) found D O / HDO > HDO / H O in the cold outer layersof the IRAS16293-2422 core from
Herschel -HIFI D O, HDOand H
O lines consistent with the models.Second, the “hot core” HDO / H O abundances of (0 . − × − (§ 8), which presumably represent thermally sublimatedices, are consistent with the bulk ice deuteration found in mod-els of the cold pre-stellar phase (Fig. 24). They also cover thesame range as the cometary data shown in Fig. 36. The threebrown HDO / H O ratios for star-forming clouds that lie some-what higher in this figure refer to isolated cores that are gener-ally somewhat colder than the clustered regions (Jensen et al.2019). Since our Solar System is thought to be born in a cluster(Adams 2010), the lower blue points are more appropriate forcomparison with comets.The “smoking gun” for the inheritance scenario comes fromthe observed high abundance of D O, with D O / HDO ≈ / H O for the NGC 1333 IRAS2A hot core (Coutens et al.2014). Similarly, the detection of D O in comet 67P / ChuryumovGerasimenko provides an important clue (Altwegg et al. 2017).The inferred D O / HDO ratio for 67P is a factor of 17 higherthan the HDO / H O ratio. Both values (high points in Fig. 36) areconsistent with pre-stellar multilayer water ice formation and notwith water reformed in the outer part of the solar nebula disk.Could the deuteration fractions actually be modified en routefrom cloud to disks or be established within disks? The deliv-ery of layered ice mantles from cloud to disk has been inves-tigated by Furuya et al. (2017). These models confirm that themajority of pre-stellar water ice is retained upon delivery to thedisk without significant UV processing and ice sublimation. TheHDO / H O ratio is somewhat lowered (by up to a factor of two)because of selective loss of the upper ice layers but is largely re-tained when averaged over the entire disk, as is D O / HDO. Lo-cally, the HDO / H O can deviate more from the original value,especially in the upper layers of the outer disk.How can this inheritance case be distinguished from the fullreset case in disks? Cleeves et al. (2014) show that starting fromthe elemental [D] / [H] ratio, mature Class II disks cannot produceenough HDO to explain the measured level of water deuterationin comets of HDO / H O ≈ (0 . − × − (see also Willacy &Woods 2009), a conclusion that is strengthened if cosmic raysare excluded from disks by stellar winds. There have been some Article number, page 42 of 58.F. van Dishoeck et al.: WISH: physics and chemistry from cloud to disks probed by
Herschel spectroscopy claims that disk models with vertical turbulent mixing can reachthe required HDO / H O levels (Furuya et al. 2013; Albertssonet al. 2014), but the recent observational limits on turbulent mix-ing (e.g., Flaherty et al. 2018, 2020; Teague et al. 2018) makethose models less realistic.In any case, there would be two clear ways to distinguish theinheritance and reset models: (i) in-situ disk chemistry modelspredict an increasing HDO / H O ratio with radius due to the de-creasing temperature profile, whereas this is not always the casefor the inheritance models; (ii) reformed ices in the disk haveD O / HDO ∼ × HDO / H O whereas the original pre-stellar iceshave ∼ × HDO / H O (Furuya et al. 2017).Although the results for D O / HDO hinge on only one pro-tostar and one comet observed so far, the HDO / H O ratios aswell as the models discussed above point to the same conclu-sion. Further observations of D O in protostellar sources shouldhave high priority, and new ALMA data indeed strengthen theinheritance case (Jensen et al., unpublished).In conclusion, several lines of observational and modelingevidence on deuterated water, together with the results for youngdisks (§ 9.1), point toward the bulk of the water being formed inthe cold pre-stellar stage and incorporated largely unaltered intodisks and comets, where they are locked up quickly into largergrains and planetesimals.
A general lesson from WISH + is that most of the new sciencecomes from the deepest and longest integrations, either to detectweak lines or to get high signal-to-noise on strong line profiles,pushing the instruments to the limits. For WISH, such long in-tegrations were possible to plan because of the freedom that theGTO time o ff ers. In contrast, time allocation committees, in par-ticular on Galactic science, are often reluctant to grant significanttime per source. Future missions should allow for a mechanismfor deep integrations on Galactic sources early in the mission.The importance of future mid- and far-infrared observationsand millimeter interferometry to trace water has been high-lighted near the end of each of the sections, §5–10. Summarizedper facility or instrument, each has its strength in providing (partof the) answers on the water trail from cloud to disk. In fact, it isclear from the discussions in §5–10 that the full picture can onlybe pieced together by the combination of many di ff erent observa-tions, with the far-infrared wavelength range providing a centralrole. Roughly in order of when instruments become available,the opportunities for adding to our understanding of water are asfollows. – Millimeter interferometers, most notably ALMA, will bevery powerful to image warm water on subarcsec scales ( < O 203 and390 GHz and various HDO and D O lines. Mm arrays canobserve optically thin water isotopologs at low frequenciesat which the dust continuum emission is optically thin, incontrast with space-based data at high frequencies. So far,ALMA has observed just a few sources; large surveys areneeded. ALMA can also put constraints on water releasedor produced at the accretion shock. By observing HDO andD O in the same beam, direct and model-independent probesof HDO / H O and D O / HDO in warm gas are obtained. Foroutflow shocks, higher excitation (nonmasing) H
O linesthat are not obscured by the atmosphere can be imaged. – The VLT-CRIRES + and SOFIA-EXES high-resolution R ≈ near- to mid-infrared spectrometers are well suited to ob-serve hot water absorption lines at 3 and 6 µ m arising in theinner envelope or outflow toward protostars. By spectrallyresolving the lines, both the origin as well as accurate opti-cal depths and thus column densities can be determined. – VLT-CRIRES + and Keck-NIRSPEC, and in the future ELT-METIS (planned ∼ Ovibration-rotation emission lines at 3 µ m in disks at R ≈ and constrain their location through systematic velocity pat-terns. Moreover, ELT-METIS can spatially resolve the linesdown to a few au and distinguish a disk surface layer ver-sus disk wind origin. These 3 µ m vibration-rotation linesare also bright in comets, and provide an accurate way todetermine cometary HDO / H O when combined with near-simultaneous observations of HDO.Individual mid-infrared lines in the 10 µ m region can be tar-geted with VLT-VISIR at R ≈ in bright sources; higherresolution instruments such as TEXES on Gemini are onlyo ff ered sporadically. ELT-METIS will o ff er high spatial res-olution at N-band, but with a spectral resolving power of onlya few hundred, limiting the line-to-continuum ratio. – The SOFIA-GREAT instrument is unique in observations ofspectrally resolved lines above 1 THz, most notably of the[O I] 63 and 145 µ m lines. Such observations will be of cru-cial importance to nail down the oxygen budget in shocksand other warm regions. OH, NH and high- J CO can alsobe observed, albeit with reduced sensitivity compared with
Herschel . SOFIA-4GREAT now also allows observations ofthe H
O ortho-ground state line in bright sources, includ-ing comets. Also, velocity resolved [C II] line mapping as atracer of UV radiation on molecular cloud scale is possiblewith SOFIA-UPGREAT. – The JWST-MIRI instrument will probe hot and warm waterabundances in shocks and inner disk surface layers, throughthe 6 µ m vibration-rotation band and the pure rotational linesat > µ m. For disks, this combination of lines will allowmore accurate assessments whether the inner regions insidethe ice line are enhanced or depleted in water. For shocks,the combination of observed H O and H lines will be key todetermine whether water accounts for the full oxygen bud-get in hot gas, or whether any UDO is still required. JWST-NIRSPEC can probe the hottest water emission at 3 µ m, aswell as the primary H O, HDO and CO ice absorption bandsat 3–5 µ m, even spatially resolved on 10-20 au scales. – The approved SPHEREx mission with a launch date around2024 will obtain low resolution absorption spectra at 0.75–5 µ m toward > lines of sight with strong back-ground sources and thus measure H O ice and other oxygen-containing ices such as CO , OCN − , CO, and OCS along thelines of sight, further testing and quantifying ice formationand ice chemistry models. – Future far-infrared missions such as the proposed
OriginsSpace Telescope and the proposed but recently cancelledSPICA mission are particularly well suited to trace the watertrail by observing a very wide range of mid- and far-infraredwater lines (Fig. 1), covering both hot (1000 K), warm ( ∼ ∼
30 K) gas with a resolving power of up to3 × for Origins . This includes the critical 179 µ m lineconnecting with the ground state. For protostellar envelopesand shocks, the analysis would have to take into account thatthere are likely multiple physical components (as found byWISH) that are not spectrally resolved. Also, the flux of low-lying water lines and [O I] may be significantly a ff ected byabsorption. Article number, page 43 of 58 & A proofs: manuscript no. wish_final_Jan2021
For disks, far-infrared missions would be a very powerfultool to locate the water snowline and quantify the water vaporin the di ff erent disk water reservoirs in statistically signifi-cant samples of disks. Indeed, models show that the largestsensitivity to the location of the snow line is provided bylines in the 40–60 µ m region, which is exactly the wave-length range without recent observational facilities. More-over, the 44 and 62 µ m (crystalline) water ice features canbe observed in emission, allowing to quantify the bulk of thewater and oxygen in the emitting layer at intermediate diskheights. In the current design Origins is four orders of magni-tude more sensitive to the mass of cold water vapor emissionin disks compared to
Herschel . Observing a thousand disksystems across the range of stellar masses will enable mean-ingful comparison to the exoplanet composition and inven-tory.None of these planned space missions include a THz het-erodyne capability. Only the Russian-led
Millimetron proposalhas such an instrument in its baseline plan, whereas the
Origins -HERO instrument is an upscope option. This lack of heterodynefacility will strongly limit studies of cold quiescent water gas inclouds and outer envelopes as well as in protoplanetary disks forwhich R > is needed to retrieve abundance profiles, whetherfrom absorption or emission line profiles. For example, relatingthe water abundance to substructures in disks, such as appearsthe case for the HD 100546 disk (§ 9.2), cannot be done withoutfull velocity resolution. All three water lines connecting with theground-state – 557, 1113, 1670 GHz (179.5 µ m) — are impor-tant, as are their isotopolog lines. ALMA cannot do such studiesdue to the blocking of these lines by the Earth’s atmosphere. Todetermine HDO / H O in cold gas, the 893 GHz 1 − HDOline is particularly useful. For warmer regions such as protostarsand shocks, spectrally resolved profiles of the 988, 1097 and1153 GHz lines should be added to retrieve physical parametersof individual velocity components.
12. Conclusions
Here we summarize in abbreviated form the main conclusionsthat water observations with
Herschel have taught us on thephysics and chemistry in star-forming regions, and their evolu-tion from cloud to disk. We also reiterate some of the broader im-plications of the WISH + program, for example for the interpre-tation of extragalactic data on water. Key points are that Herschel has revealed new physical components, has tested and confirmedbasic chemical networks, and has pointed the way for the inter-pretation of subsequent (ground-based) data of other molecules,including the various transitions of CO itself, both in Galacticand extragalactic sources. – Water traces active star formation sites: the bulk of the ob-served gaseous water emission with
Herschel in star-formingregions is associated with warm kinetically heated gas in out-flows and shocks that trace gas of several hundred K. Thissame gas is traced by CO lines with J up >
14, but not bylower- J CO lines. – Water spectral profiles are complex: at least two di ff erentphysical components are universally seen in water and high- J CO lines for low-mass sources: a broad component (cavityshock, disk wind, turbulent entrainment) with T rot (CO) ≈
300 K, and a medium-broad o ff set (slightly blue-shifted) disso-ciative J − type ‘spot shock’ with T rot (CO) ≈
700 K. The lat-ter category includes extremely high velocity features withvelocities out to ±
100 km s − seen in a small fraction ofsources. For higher-mass sources, only the broad componentis seen. Very few cloud positions show narrow (FWHM < fewkm s − ) water emission lines. – Water profiles probe small motions: inverse P-Cygni pro-files due to infall are seen in a fraction of protostellar sources,indicating mass infall rates from cloud onto envelope of10 − − − M (cid:12) yr − for low-mass protostars, and increasingup to 10 − − − M (cid:12) yr − for high-mass sources. A slightlysmaller fraction of sources show regular P-Cygni profiles in-dicating expansion. – Water emission is compact: for most low-mass sources, thewater emission, and thus the mechanisms that produce it,is limited to ∼ < < ff source, the emitting areas are larger as thewarm gas can expand in more directions. – Water traces warm, dense gas: the inferred physical condi-tions in the water emitting gas are high density ( > cm − )and kinetic temperature (300–1000 K) with small emittingareas on source, of order 100 au for low-mass protostars and1000 au for high-mass protostars. Herschel was not sensitiveto a possible very hot component of several thousand K. – Water points to UV-irradiated outflow cavities andshocks: the data on H O, OH and other hydrides indicateUV fields up to 10 − times the general interstellar ra-diation field in both low- and high-mass sources in outflowcavity walls on scales of the Herschel beam ( ∼ (cid:48)(cid:48) ). A newclass of UV-irradiated outflows or shocks is required to ex-plain the data. – Water is a significant but not dominant coolant: Far-infrared line cooling in low-mass protostars is dominated byCO and H O in the earliest low-mass Class 0 and I stages,with [O I] becoming relatively more important in the laterstage. The total H O line cooling does not change from Class0 to Class I, whereas that of CO decreases by a factor of 2.The absolute [O I] cooling is similar from Class 0 to ClassII, but its fraction increases as the jet changes from beingmostly molecular to being primarily atomic. For high-masssources, line cooling is dominated by CO, with a much lowerfraction of H O and OH cooling than for low-mass sources. – Water in extragalactic sources traces buried star forma-tion activity , originating from scales much smaller than en-tire molecular clouds. – Water abundance in outflows and shocks is low:
The waterabundance in warm outflows and shocks is universally foundto be low, only ∼ − , much less than the H O / H abun-dance of 4 × − expected if all volatile oxygen is driveninto water. Only very hot gas ( > − . Simplified chemical models con-firm that such high temperatures are needed to drive the bulkof the oxygen into water in dense gas with high UV fields. – High temperature routes to water confirmed: comparisonwith line profiles of other ice species such as CH OH andNH shows that ice sputtering is only significant at low ve-locities (out to ±
10 km s − ). High temperature gas-phaseformation dominates water production at high velocities. Article number, page 44 of 58.F. van Dishoeck et al.: WISH: physics and chemistry from cloud to disks probed by
Herschel spectroscopy – Water abundance profiles in cold cores retrieved: velocityresolved HIFI line profiles have allowed the water abundanceto be derived as a function of position across pre-stellar coresand protostellar envelopes, even if the sources are not spa-tially resolved or mapped. A relatively high gaseous waterabundance of ∼ − is found in the outer layers, decreasingroughly inversely with density deeper into the core. – Simplified water chemistry networks explain data: smallgas-grain chemical models including freeze-out of O, forma-tion of H O ice, (photo-)desorption of ice, and photodissoci-ation of gaseous H O, work well to explain water profiles ina wide range of low-mass sources. The best fits imply a rangeof the external radiation field G ISRF , cosmic-ray induced field G CR and pre-stellar core lifetimes. – Water ice forms early: Water ice is primarily formed inthe cold pre-stellar stage. Some small enhancement (but nodestruction) is possible in the cold outer envelopes up to ∼ −
20 K. – Bulk HDO / H O ice is lower than cold HDO / H O gas:
HDO ground-state line profiles obtained with
Herschel arewell fitted by a constant HDO / H O abundance ratio of ∼ / H O abundance ratio in cold gas and in the outer-most layers of the water ice which are enriched in deuterium,not that of the bulk of the ice. Indeed, observed HDO / H Ovalues in hot cores, where the bulk ices sublimate, are up toan order of magnitude lower. A high D O / HDO ratio com-pared with HDO / H O is a signature of cold gas and icechemistry in the dense pre-stellar phase. – NH does not follow H O: In stark contrast with H O andHDO, the gaseous NH abundance profile inferred from Her-schel data is flat with radius throughout pre-stellar cores andprotostellar envelopes. This strengthens the conclusion thatthe bulk of the gaseous NH does not result from photodes-orption of ammonia ice, but from cold gas-phase processes. – Inner hot cores are dry: inner hot core abundances derivedwith a step-function retrieval analysis show a large variationfrom 10 − to a few × − for low- and high-mass sources,with only a few sources at the upper end of this range. Thissuggests that most hot cores are “dry”, although not as dry asoriginally thought in analyses in which the small scale phys-ical structure is ignored. Herschel data are not well suited toconstrain hot core abundances due to high optical depths inTHz continuum and lines. ALMA and NOEMA can makemajor advances on this topic. – Similarity of low- and high-mass sources: all conclusionson abundances and chemistry hold across the luminosityrange. – Puzzling oxygen budget strengthens evidence for UDO: the combined analysis of water gas, water ice and O limitsin cold clouds indicates that a large fraction of the oxygenbudget is unaccounted for. Within the simple water chem-istry models, the only solution is to have a short pre-stellarstage of only 0.1 Myr to prevent all O being turned into H O.An alternative option is for dense cores to have a small frac-tion of large ( > µ m) grains which hide more than 50% ofthe water ice from being observed through infrared absorp-tion. This solution does not apply to hot cores and shocks,where the large icy grains should have sublimated and wherea large fraction of oxygen is also missing. Another option istherefore that oxygen is in some refractory form called UDO,whose abundance increases from di ff use to dense clouds, andwhich consists of material that does not vaporize or atomizeeven in strong shocks up to 1000 K. – The water ice reservoir in disks is spatially confined andencompasses less of the disk than expected: the weak wa-ter vapor emission from protoplanetary disks suggests eitherradial confinement of large icy grains in the inner disk or inseveral radial rings, or vertical settling of icy grains to themidplane below heights where UV photons can penetrate forphotodesorption. This finding suggests that low gas-phaseoxygen abundances in outer disks are common. In conclusion, what have we learned about water in star-formingregions? – Most water forms prior to star formation: most water isformed as ice in dense molecular clouds before they collapseto form stars. However, not all oxygen may be turned intowater ice, possibly implying a short pre-stellar phase (§ 6,7). – Water formed through high- T chemistry is mainly lostto space: significant amounts of water are also produced orreformed by high temperature chemistry in warm outflowsor shocks (§ 5). However, most of this water is lost to spacein outflows and does not contribute to the water reservoir inplanet-forming disks. – Water is mostly transported as ice during protostellarcollapse and infall: water enters the (forming) disk at largeradii, with no observational evidence yet that it is a ff ected byaccretion shocks (§ 9.1). – Water locked up early in large grains in outer disks: thewater vapor content in the outer regions of protoplanetarydisks indicates that the reservoir of available water ice issmaller than expected, being either radially or vertically con-fined or both. The weakness of the HIFI lines implies thatice-coated grains grow quickly, settle to the midplane anddrift inward when grown to mm / cm sizes. This process likelystarts already in embedded phase. It also implies that diskswith a low gas-phase oxygen abundance are common, in linewith the findings of low volatile carbon abundances (§ 9.2). – Inner disk water reservoir still to be probed: warm waterin the disk surface layers has been observed by Herschel and
Spitzer but the inner disk midplane inside the water iceline isstill invisible (§ 9.2.2). – New planetary systems are likely to be born with su ffi -cient water to become habitable: the Herschel data haveshown that water gas and ice are commonly associated withstar-forming regions and that this conclusion is independentof “environment” or location in our Galaxy. A key require-ment is that the cloud out of which the star and disk areformed is cold enough ( (cid:46)
25 K) to enable water ice for-mation and has a long enough lifetime prior to collapse( (cid:38)
Herschel has shown that water has lived upto its reputation of being a particularly interesting molecule forstudying the physics and chemistry of star-forming regions. Thelarge abundance changes of gas-phase water between cold andwarm regions –now fully quantified– as well as the sensitivity ofits line profiles to small motions, make water a unique diagnosticprobe among the suite of interstellar molecules.
Herschel has
Article number, page 45 of 58 & A proofs: manuscript no. wish_final_Jan2021 left a legacy for the analysis and interpretation of future waterobservations in Galactic and extragalactic sources. Also, even if
Herschel has found that not all available oxygen is locked upin water gas or ice, this does not diminish its importance as akey ingredient for habitability on other planetary systems: wateris found to be abundantly present in star- and planet-formingregions.
Acknowledgements.
The authors would like to thank all WISH team membersover the years for their seminal contributions to this project, as well as the en-tire HIFI science consortium for two wonderful decades of working togetherto make this happen. They are particularly grateful to Malcolm Walmsley, whostimulated the WISH program from the very beginning and helped with numer-ous projects. He left a great legacy. Fruitful collaborations with the DIGIT andWILL teams, and with members of the HDO team, are also acknowledged. De-tailed discussions with Kathrin Altwegg and Martin Rubin on the oxygen budgetin comets and with Bruce Draine on interstellar clouds are appreciated. Con-structive comments from the referee have helped to improve the paper. A bigsalute goes to the HIFI and PACS instrument teams and to ESA for design-ing, building and operating these two powerful instruments and the
Herschel
Space Observatory.
Herschel was an ESA space observatory with science instru-ments provided by the European-led Principal Investigator consortia and withimportant participation from NASA. HIFI was designed and built by a con-sortium of institutes and university de-partments from across Europe, Canadaand the US under the leadership of SRON Netherlands Institute for Space Re-search, Groningen, The Netherlands with major contributions from Germany,France and the US. Consortium members are: Canada: CSA, U. Waterloo;France: CESR, LAB, LERMA, IRAM; Germany: KOSMA, MPIfR, MPS; Ire-land, NUI Maynooth; Italy: ASI, IFSI-INAF, Arcetri-INAF; Netherlands: SRON,TUD; Poland: CAMK, CBK; Spain: Observatorio Astronómico Nacional (IGN),Centro de Astrobiologıa (CSIC-INTA); Sweden: Chalmers University of Tech-nology - MC2, RSS & GARD, Onsala Space Observatory, Swedish NationalSpace Board, Stockholm University - Stockholm Observatory; Switzerland: ETHZürich, FHNW; USA: Caltech, JPL, NHSC. PACS has been developed by aconsortium of institutes led by MPE (Germany) and including UVIE (Aus-tria); KUL, CSL, IMEC (Belgium); CEA, OAMP (France); MPIA (Germany);IFSI, OAP / OAT, OAA / CAISMI, LENS, SISSA (Italy); IAC (Spain). This de-velopment has been supported by the funding agencies BMVIT (Austria), ESA-PRODEX (Belgium), CEA / CNES (France), DLR (Germany), ASI (Italy), andCICYT / MCYT (Spain) Astrochemistry in Leiden is supported by the Nether-lands Research School for Astronomy (NOVA). JRG thanks the Spanish MICIUfor funding support under grants AYA2017-85111-P and PID2019-106110GB-I00. Part of this research was carried out at the Jet Propulsion Laboratory,California Institute of Technology, under a contract with NASA. DF acknowl-edges financial support from the Italian Ministry of Education, Universities andResearch, project SIR (RBSI14ZRHR) as well as project PRIN-INAF-MAIN-STREAM 2017. AK acknowledges support from the Polish National ScienceCenter grant 2016 / / D / ST9 / / References
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A. 2013, ApJ, 766, 82 Leiden Observatory, Leiden University, PO Box 9513, 2300 RA Lei-den, The Netherlands e-mail: [email protected] Max-Planck Institut für Extraterrestrische Physik (MPE), Giessen-bachstr. 1, 85748 Garching, Germany Niels Bohr Institute & Centre for Star and Planet Formation, Copen-hagen University, Øster Voldgade 5–7, 1350 Copenhagen K, Den-mark Max Planck Institute for Astronomy, Königstuhl 17, 69117, Heidel-berg, Germany Institute for Particle Physics and Astrophysics, ETH Zurich, 8093Zurich, Switzerland Department of Astronomy, The University of Michigan, 1085 S.University Ave., Ann Arbor, MI 48109-1107, USA Lab. d’astrophysique de Bordeaux, Univ. Bordeaux, CNRS, B18N,allée Geo ff roy Saint-Hilaire, 33615 Pessac, France National Research Council Canada, Herzberg Astronomy and As-trophysics, 5071 West Saanich Rd, Victoria, BC, V9E 2E7, Canada Department of Physics & Astronomy, University of Victoria, Victo-ria, BC, V8P 1A1, Canada Department of Space, Earth and Environment, Chalmers Universityof Technology, Onsala Space Observatory, 439 92 Onsala, Sweden INAF - Osservatorio Astronomico di Roma, Via di Frascati 33,00074, Monte Porzio Catone, Italy Observatorio Astronómico Nacional (OAN), Calle Alfonso XII, 3,28014, Madrid, Spain SRON Netherlands Institute for Space Research, PO Box 800, 9700AV, Groningen, the Netherlands Kapteyn Astronomical Institute, University of Groningen, PO Box800, 9700 AV, Groningen, The Netherlands Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69,53121 Bonn, Germany INAF – Istituto di Astrofisica e Planetologia Spaziali, via Fosso delCavaliere 100, 00133 Roma, Italy Division of Geological and Planetary Sciences, California Instituteof Technology, Pasadena, CA 91125, USA LERMA & UMR8112 du CNRS, Observatoire de Paris, PSL Uni-versity, Sorbonne Universités, F-75014, Paris, France Instituto de Fisica Fundamental (IFF-CSIC), Calle Serrano 123,28006, Madrid, Spain Korean Astronomy and Space Science Institute, Daejeon 34055, Ko-rea Anton Pannekoek Institute for Astronomy, University of Amster-dam, Science Park 904, 1098XH, Amsterdam, The Netherlands INAF, Osservatorio Astrofisico di Arcetri, Largo Enrico Fermi 5,50125, Firenze, Italy Department of Physics & Astronomy, University of Waterloo, 200University Avenue, Waterloo, ON, N2L 3G1, Canada National Astronomical Observatory of Japan, 2-21-1 Osawa, Mi-taka, Tokyo 181-8588, Japan Kavli Institute for Astronomy and Astrophysics, Peking University,Yiheyuan Lu 5, Haidian Qu, 100871 Beijing, People’s Republic ofChina Institute of Astronomy, Faculty of Physics, Astronomy and Infor-matics, Nicolaus Copernicus University, Grudziadzka 5, 87-100Torun, Poland Department of Physics and Astronomy, San Jose State University,One Washington Square, San Jose, CA 95192-0106, USA Center for Astrophysics | Harvard & Smithsonian, 60 Garden Street,Cambridge, MA, 02138, USA Leiden Institute of Chemistry, Gorleaus Laboratories, Leiden Uni-versity, PO Box 9502, NL-2300 RA Leiden, the Netherlands Department of Astronomy, Stockholm University, 106 91 Stock-holm, Sweden INAF - Osservatorio Astronomico di Cagliari, via della Scienza 5,09047, Selargius, Italy Jet Propulsion Laboratory, California Institute of Technology, 4800Oak Grove Drive, Pasadena, CA 91109, USA Department of Physics & Astronomy, Johns Hopkins University,Baltimore, MD 21218, USA SRON Netherlands Institute for Space Research, Sorbonnelaan 2,3584 CA Utrecht, Netherlands School of Physics and Astronomy, University of Leeds, Leeds LS29JT, UK Center for Space and Habitability (CSH), University of Bern,Gesellschaftsstrasse 6, 3012, Bern, SwitzerlandArticle number, page 50 of 58.F. van Dishoeck et al.: WISH: physics and chemistry from cloud to disks probed by
Herschel spectroscopy
Appendix A: Overview of Herschel programs
Table A.1 summarizes the various subprograms of the WISH keyprogram as well as related
Herschel open time programs whosedata have been used in the analysis.
Appendix B: Observational details and datareduction
Appendix B.0.1: HIFI
HIFI consisted of a set of seven single-pixel dual-sideband het-erodyne receivers (de Graauw et al. 2010). All observationswere taken in both horizontal and vertical polarizations withsimultaneously the Wide Band Spectrometer (WBS) and HighResolution Spectrometer (HRS) backends providing both wide-band (WBS, 4 GHz bandwidth at 1.1MHz resolution) and high-resolution (HRS, typically 230MHz bandwidth at 250 kHz reso-lution) frequency coverage. The sideband ratio is approximatelyunity (Roelfsema et al. 2012).Observations were taken as single pointings in dual-beam-switch (DBS) mode with a chop throw of 3 (cid:48) using fast chopping.The only exception are some of the H O 1 − observationsof LM sources, which were taken in position-switch mode to anemission-free position (see Kristensen et al. 2012, for details).The Herschel -HIFI beam ranges from 12 . (cid:48)(cid:48) to 38 . (cid:48)(cid:48) over thefrequency range of the various water lines, set by the di ff ractionlimit of the primary mirror.The data were reduced within the Herschel Interactive Pro-cessing Environment (HIPE; Ott 2010; Shipman et al. 2017).After initial spectra production, first the instrumental standingwaves were removed where required. This is especially impor-tant at the highest frequencies, see Kristensen et al. (2017b) fordetails. This was followed by baseline subtraction with a low-order ( ≤ T MB scale using e ffi cien-cies from Roelfsema et al. (2012). Finally, all data were resam-pled to typically ∼ − on the same velocity grid and eitherimported into CLASS or as FITS files into custom-made pythonroutines for further analysis.Comparison of the two polarizations for each source revealedinsignificant di ff erences in line shape or gain in most cases (al-though occasional di ff erences up to 30% were seen, especially athighest frequencies), so these were co-added to reduce the noise.Comparison of peak and integrated intensities between the orig-inal WISH observations and those obtained as part of open timeprograms for the same source indicate that the calibration un-certainty is < (cid:48)(cid:48) (Roelfsemaet al. 2012). The beam positions of the H and V polarizationswere, however, o ff set by up to 6 (cid:48)(cid:48) at the lowest frequencies, soalthough the overall pointing accuracy was good, it did not meanthat HIFI was always perfectly pointed on-source. For specificpeculiarities of individual sources see discussion in Kristensenet al. (2012); Mottram et al. (2014); van der Tak et al. (2013,2019). The HIFI OTF data reduction is described in Jacq et al.(2016). Appendix B.0.2: PACS
The PACS integral field unit (IFU) spectroscopy mode was usedconsisting of two photoconductor arrays with 16 ×
25 pixels, al-lowing simultaneous observations in the red 1st order grating(102–210 µ m) and one of the 2nd or 3rd order blue gratings(51–73 or 71–105 µ m) (Poglitsch et al. 2010). The IFU had 5 × . (cid:48)(cid:48) × . (cid:48)(cid:48) each, which cov-ered a 47 (cid:48)(cid:48) × (cid:48)(cid:48) field of view. It is important to note that the Herschel beam changes with wavelength whereas the spaxel sizestays constant, so that corrections for the wavelength-dependentloss of radiation need to be applied in case of a well-centeredpoint source.The spectral resolving power R increases with wavelengthfrom about 1000 to 2000 (corresponding to velocity resolutionsof 140 to 320 km s − ) in 1st order and from about 1500 to 3000(100 to 210 km s − ) in 2nd order. At the shortest wavelengths,the velocity resolution is as good as 75 km s − , so the [O I] 63 µ mline sometimes shows velocity-resolved line profiles (van Kem-pen et al. 2010; Nisini et al. 2015; Karska et al. 2018).Most WISH observations obtained single footprint spectralmaps in IFU line-scan mode in which deep observations are ob-tained for a narrow wavelength region around selected transi-tions (bandwidth ∆ λ = . − µ m). All observations used achopping / nodding observing mode with o ff -positions within 6 (cid:48) of the target coordinates to subtract the background emission.Data reduction was performed with HIPE including spectralflat-fielding (see Herczeg et al. 2012; Green et al. 2013, 2016;Karska et al. 2018, for full details). The flux density was normal-ized to the telescopic background and calibrated using observa-tions of Neptune. The overall calibration uncertainty in flux den-sities is estimated to be ∼
20% from cross comparisons of sourcesobserved in multiple programs or modes (e.g., Karska et al.2018). 1D spectra were obtained by summing over a number ofspaxels chosen after inspection of the 2D spectral maps (Karskaet al. 2013). For sources with extended line emission, the co-addition of spaxels with detected emission increases the S / N ,smooths the continuum, and enables correction for significantdi ff erences in di ff raction-limited beam sizes over the wide spec-tral range covered by PACS. For sources with point-like emis-sion, only the central spaxel was used, multiplied by wavelength-dependent instrumental correction factors to account for the dif-ferent point source functions (PSF) (see PACS Observers Man-ual).The PACS full range spectroscopy mode, in which the entirefar-infrared spectrum from 50 to 210 µ m is obtained, was usedfor a few WISH sources (Herczeg et al. 2012; Goicoechea et al.2012) and for all DIGIT sources (which partially overlap withthe WISH sample, 8 sources in common) (Green et al. 2013).However, the spectral sampling within a resolution element isabout 3–4 times coarser in the full spectral scan than in line spec-troscopy mode, and integrations are generally less deep. Figure 7shows part of the NGC 1333 IRAS4A spectral scan comparedwith that of the neighboring 4B source (Herczeg et al. 2012),with the latter clearly richer in lines. Figure 3 illustrates the im-portance of spectral sampling for a few PACS lines observed inboth line and range spectroscopy modes. Appendix C: Additional water spectra
Figure C.1 presents the p-H O 1 − spectra of all high-masssources, illustrating the multitude of foreground clouds along thelines of sight. Article number, page 51 of 58 & A proofs: manuscript no. wish_final_Jan2021
Table A.1.
WISH key program and related OT programsProgram Subprogram (co-)PIs Hours + selected mapsPre-stellar cores P. Caselli 2LM protostars L. Kristensen / J. Mottram 29Outflows R. Lisau / B. Nisini / M. Tafalla 26IM protostars M. Fich / D. Johnstone 6HM protostars F. van der Tak / F. Herpin / F. Wyrowski 24Disks M. Hogerheijde / E. Bergin 12Radiation diag. A. Benz 12WILL E.F. van Dishoeck / J. Mottram 134 49 LM selected HIFI + PACS linesDIGIT N.J. Evans 250 30 LM full PACS scansCOPS-HIFI L. Kristensen 14 24 LM HIFI CO 16-15 / OHLM-deep R. Visser 20 5 Deep HIFI excited H
OOutflows B. Nisini 54 5 LM PACS [O I], CO, H O mapsWater maps R. Liseau 19 6 HIFI H O maps + PACS [O I]IM Cygnus S. Bontemps HIFI selected linesHM ATLASGAL F. Wyrowski HIFI selected linesDisks M. Hogerheijde 135 4 Very deep HIFIWatCH S. Wampfler 19 8 HIFI HCO + , OH, H O + AFGL2591x2.0 DR21(OH)x0.8 G5.89 G10.47x2.0 G29.96x3.0 G31.41x2.0 G34.26 G327 I05358x3.0 I16272x3.0 I18089x4.0 I18151x2.0 N6334-I N6334-I(N) N7538-I1x1.5 − T M B ( K ) W3-IRS5x0.8 υ − υ LSR (km s − ) − W33Ax2.0 − W43-MM1x2.0 − W51-E1
Fig. C.1. H O 1 − spectra toward all high-mass sources in the WISH program. Figure C.2 presents the averaged and normalized water spec-tra and high- J CO spectra for each class of sources, from low-to high mass. The water spectra arise from medium- J levels thatdo not connect with the ground state so that no absorption fea-tures are present. The similarity in profiles within each class of sources and between classes of sources is clearly seen, with onlythe low-mass Class I profiles narrower than the others.Figure C.3 presents deep integrations of the excited 3 − H O transition ( E up =
250 K) toward low-, intermediate-, andhigh-mass sources. The line is barely detected toward the low-and intermediate-mass sources, as opposed to the high-mass
Article number, page 52 of 58.F. van Dishoeck et al.: WISH: physics and chemistry from cloud to disks probed by
Herschel spectroscopy source where it is prominent. However, in spite of the di ff erencein S / N , there is no indication that the line emission traces dif-ferent components toward the di ff erent sources: it probes warmwater in the inner envelope on small scales. Appendix D: Water chemistry: three routes
Appendix D.1: Water and related species
The chemical reactions that form water in interstellar space havebeen well described in various papers, details of which will notbe repeated here (see recent overviews in Hollenbach et al. 2009;van Dishoeck et al. 2013; Lamberts et al. 2013). References topapers are mostly those that have appeared since the 2013 re-views. Three main routes can be distinguished, which are sum-marized in Fig. 4: (i) gas-phase ion-molecule chemistry, drivenby cosmic ray ionization of H and H ; (ii) high temperature gas-phase neutral-neutral chemistry initiated by the O + H reaction;and (iii) ice chemistry.Route (i) dominates the formation of water in cold low-density di ff use and translucent clouds with visual extinctions ofa few magnitudes. Herschel -HIFI has probed this chemistry indetail through far-infrared absorption line observations of H Oas well as intermediates OH + and H O + as part of the PRISMASkey program (Flagey et al. 2013; Gerin et al. 2016). Typical ob-served H O abundances are low, only (0 . − . × − withrespect to total hydrogen, which are well reproduced by mod-els such as those by Hollenbach et al. (2009) within factors of2. In the WISH sample, the intermediate species OH + , H O + and H O + have been detected together with water in foregroundclouds toward many protostars (e.g., Benz et al. 2010; Wyrowskiet al. 2010; Benz et al. 2016). Analysis of these WISH + PRIS-MAS data have resulted in a mean value of the cosmic ray ion-ization rate of ζ H = . × − s − in di ff use clouds (Indrioloet al. 2015b).Route (ii) is much more relevant for star-forming regions,where temperatures close to protostars and in winds can be upto a few hundred K and in shocks as high as a few thousand K.The reaction O + H → OH + H is endothermic, whereas thesubsequent reaction OH + H → H O + H is exothermic but hasa significant energy barrier of ∼ ∼
250 K. Water destruction can occur by thereverse reaction with atomic H, but this reaction has a very highbarrier of ∼ K and is unlikely anyhow in dense moleculargas because of the low abundance of atomic H. This leaves UVphotons as the main route for water destruction, with X-rays onlye ff ective in regions of high extinction. Photodissociation of H Ostarts to be e ff ective shortward of 1800 Å and continues down tothe ionization threshold at 983 Å (12.61 eV), including Ly α at1216 Å. Its lifetime in the general interstellar radiation field, asgiven by Draine (1978), is only 40 yr (Heays et al. 2017). Closeto (proto)stars, the UV radiation field is enhanced by many or-ders of magnitude, making the water vapor lifetime even shorter.Route (iii) dominates the formation of water in cold denseclouds, and is in fact the route that produces the bulk of waterin space (Bergin et al. 2000; van Dishoeck et al. 2013, for re-view) since freeze-out of gaseous water is insu ffi cient (Lee et al.1996). The timescale for an atom or molecule to collide with agrain and stick to it is t freeze = × / n (H ) yr for normal 0.1 µ m size grains and sticking probabilities close to unity. Sincethe freeze-out time scales as 1 / n whereas the free-fall time of acloud (which is a lower limit to its lifetime) scales as 1 / √ n , the former becomes shorter than the latter for densities > cm − so most species will end up on the grains (e.g., Bergin & Tafalla2007). The grain temperature plays a key role in the ability toretain species on grains and for surface reactions to occur. Lightspecies such as H and H are weakly bound so they have a highmobility and short residence time on the grains. In fact, for graintemperatures above ∼
20 K, this residence time is so short thathydrogenation ceases to be e ff ective, except through the directEley-Rideal mechanism. Heavier species such as O and OH aremore strongly bound and less mobile. These rates are usuallyparametrized through binding energies E b and di ff usion energies E d , with the latter often taken to be a fraction of 0.3–0.7 of E b .At the lowest temperatures, tunneling reactions of H and H canbe important as well (Meisner et al. 2017; Lamberts & Kästner2017; Cuppen et al. 2017).Tielens & Hagen (1982) postulated that the formation of wa-ter from O atoms on grain surfaces proceeds through three routesinvolving hydrogenation of s − O, s − O and s − O , respectively,where s − X indicates a species on the surface. All three routeshave subsequently been studied and quantified in various labora-tories (see Linnartz et al. 2015, for summary) (Fig. 4).Water ice can be returned to the gas by various desorptionprocesses. The most e ff ective one is thermal desorption, whichoccurs on timescales of years when the dust temperature risesabove ∼
100 K (precise value depends on pressure). This leads togas-phase abundances of H O as high as the original ice abun-dances, if there is no subsequent destruction of gaseous H O.This process dominates the production of gaseous water in thewarm inner parts of high- and low-mass protostellar envelopes(“hot cores” or “hot corinos”) and inside the snow line in disks.Water ice can also be sputtered by energetic particles in shocks,but not all water ice may be returned to the gas and some of itmay be dissociated (see § 5 for details).At dust temperatures below the sublimation limit, photodes-orption is an e ff ective mechanism to get water back into the gasphase, even though only a small fraction of the UV absorptionsresults in desorption of intact H O molecules (Andersson & vanDishoeck 2008). The e ffi ciency is about 10 − per incident pho-ton within factors of a few, as determined through laboratoryexperiments and theory (e.g., Arasa et al. 2015; Cruz-Diaz et al.2018). Only the top few monolayers of the ice contribute. Thus,icy clouds exposed to radiation from a nearby star or the generalinterstellar radiation field will have a layer of enhanced watervapor at the edge, the so-called “photodesorption layer” (Hol-lenbach et al. 2009).Deep inside clouds, cosmic rays produce a low level ofUV flux, ∼ photons cm − s − , through interaction with H (Prasad & Tarafdar 1983; Shen et al. 2004; Ivlev et al. 2015).This process turns out to be crucial for explaining the water ob-servations of cold dense clouds. Other nonthermal ice desorptionprocesses include cosmic ray induced spot heating (which maywork for CO, but is generally not e ffi cient for strongly boundmolecules such as H O) and desorption due to the energy liber-ated by the formation reaction (called “reactive” or “chemical”desorption). The latter process is less well understood and con-strained than photodesorption and may not be as e ff ective on wa-ter ices as on other surfaces (Dulieu et al. 2013; Minissale et al.2016).Stimulated by the discovery of abundant O ice in comet 67Pbut lack of H O and O ice (Bieler et al. 2015), a new criticallook has been taken at the ice formation network (Taquet et al.2016; Eistrup & Walsh 2019). Important parameters whose val-ues are still uncertain and that are found to have an impact are (i)the di ff usion to binding energy ratios E d / E b ; (ii) the binding en- Article number, page 53 of 58 & A proofs: manuscript no. wish_final_Jan2021 −
50 0 50 v ( km s − ) T m b ( a r b it r a r yun it s ) HMIMLMILM0H O2 − −
50 0 50 v ( km s − ) HMIMLMILM0H O2 − −
50 0 50 v ( km s − ) HMIMLMILM0H O3 − −
50 0 50 v ( km s − ) − CO16 − Fig. C.2.
Averaged and normalized mid- J water and CO spectra for low-mass Class 0 (LM0), low-mass Class I (LMI), intermediate mass (IM)and high mass (HM) sources. From left to right the H O 2 − line at 988 GHz, 2 − line at 752 GHz, the 3 − line at 1097 GHz, andthe CO J = − −
25 0 25 50 υ - υ source (km s − ) . . . . . . . . T M B ( K ) N1333-I2A x 25NGC2071 x 2W3-IRS5 / 130
Fig. C.3.
Deep H
O 3 − spectra obtained toward a low-mass,intermediate-mass and high-mass protostar (from top to bottom) (basedon Visser et al. 2013). ergy of atomic oxygen on ice, E b (O), now found to be ∼ E a of the reactions O + O and H + O . Lamberts et al. (2014) conclude that the O + H reaction isnot important for water ice formation at low temperatures.Taken together, it is gratifying to conclude that the advent of Herschel has stimulated a number of laboratory and theoretical chemical physics studies of fundamental gas-phase and solid-state processes involved in the water network.
Appendix D.2: CO chemistry Another molecule which could potentially be a significant oxy-gen carrier is CO . Its chemistry is loosely connected with thatof H O, primarily through the OH radical.Gas-phase formation of CO takes place mostly through thereaction CO + OH → CO + H, which has a slight activationbarrier of 176 K (Smith et al. 2004). The CO production rate iscontrolled by the availability and fate of OH, since OH can alsobe consumed by the H O formation route, OH + H → H O + H.This route has a higher activation energy, 1740 K, than that lead-ing to CO , so CO production is favored at lower temperatures.However, since H is orders of magnitude more abundant thanCO, the formation of H O dominates at higher temperatures, typ-ically above 150 K (Bosman et al. 2018b). Thus, gaseous CO formation is only e ff ective in a quite narrow temperature range,50-150 K, and then only when OH is present as well. In denseregions, this requires UV photons or X-rays to liberate O andOH from H O or CO.The grain-surface formation of CO is thought to proceedprimarily through the s-CO + s-OH → s-HOCO route, with mostof the s-HOCO subsequently converted into s-CO as found inlaboratory experiments (e.g., Ioppolo et al. 2011) and theoreticalcalculations (Arasa et al. 2013). Again the reaction rate dependson the s-OH abundance, usually created by cosmic-ray inducedphotodissociation of s-H O. Alternative proposed CO forma-tion routes include the energetic processing of ice mixtures (e.g.,Ioppolo et al. 2009) and UV irradiation of water-ice covered hy-drogenated carbon grains (Mennella et al. 2006).The s-CO binding energy is about 2300 K (Noble et al.2012). This implies that the s-CO desorption temperature (itssnow line) is in between that of H O and CO.
Article number, page 54 of 58.F. van Dishoeck et al.: WISH: physics and chemistry from cloud to disks probed by
Herschel spectroscopy
Appendix E: Excitation and radiative transfer ofwater vapor
The populations of the H O and CO energy levels are determinedby the balance between the collisional and radiative excitationand de-excitation rates. A major e ff ort was undertaken by thechemical physics community in preparation for Herschel to cal-culate the collisional rate coe ffi cients of H O and CO with H over a wide range of temperatures. In WISH, the latest valuesfrom Daniel et al. (2011) for H O and Yang et al. (2010) andNeufeld et al. (2010) for CO have been used, as tabulated inLAMDA (Schöier et al. 2005). More recently, rate coe ffi cientsfor CO + H up to high temperatures have become availablewhich play a role in regions with high H / H fractions (Walkeret al. 2015).The water radiative transfer is particularly complex, becausethe large Einstein A values, subthermal excitation, (very) highoptical depths and large numbers of lines that need to be in-cluded. However, many lower J water lines are e ff ectively thin,a regime in which the line brightness still scales with columndensity (e.g., Snell et al. 2000). This assumption holds as longas the collision rate is so slow that every excitation immediatelyleads to a radiative de-excitation and the production of one pho-ton which escapes the cloud, possibly after many absorptionsand re-emissions, before another excitation occurs. In formulae,this implies C < A /τ , where C is the collision rate and τ the op-tical depth (Linke et al. 1977; Keto et al. 2014). To properly treatwater vapor excitation, many energy levels need to be taken intoaccount; while this slows down the computing speed, the advan-tage is that some of the higher-lying lines, especially those in thePACS domain, become optically thin so that they better the exci-tation models. Another complication is that at high frequencies,the continuum can become optically thick as well, thus limitingthe depths for looking into the innermost part of the source (Cer-nicharo et al. 2006; van Kempen et al. 2008; Visser et al. 2013).This far-infrared dust continuum can also e ff ectively couple withthe water excitation (González-Alfonso et al. 2014).WISH team members developed and used a number of ra-diative transfer codes to analyze the Herschel data, ranging fromthe simple constant density and temperature non-LTE excita-tion program RADEX using a local escape probability (van derTak et al. 2007) to more sophisticated 1D nonlocal Monte-Carlocodes RATRAN (Hogerheijde & van der Tak 2000) and MOL-LIE (Keto et al. 2004) that can handle radial temperature anddensity variations as expected in cold cores and protostellar en-velopes and that also follow the transport of continuum photons.For nonspherical geometries, a 2D version of RATRAN and the3D code LIME (Brinch & Hogerheijde 2010) have been used,as well as a fast 3D code using a local source approximationby Bruderer et al. (2010). Also, team members had access tothe Onsala Accelerated Lambda Iteration (ALI) code for arbi-trary source structures that can treat line overlap (Maercker et al.2008). These codes have been tested against each other in a 2004workshop (van der Tak et al. 2005) and subsequent WISH teammeetings.
Appendix F: Modeling approaches
Appendix F.0.1: Retrieval models
The simplest analysis method is to use a non-LTE excitationand radiative transfer code such as RADEX to derive best-fittingdensities and temperatures from observed line ratios. Also col-umn densities and emitting areas can be derived from the ob- served line intensities, although there is often some degeneracybetween parameters (see van der Tak et al. 2010; Herczeg et al.2012; Mottram et al. 2014; Santangelo et al. 2014b, for exam-ples). This method is most appropriate for shock componentsseen in on-source line profiles and at shock positions o ff set fromthe source.To constrain the envelope abundance, a retrieval method iscommonly used in which a trial water abundance is chosen. Usu-ally a step function with a low outer and high inner abundance(when T >
100 K) is adopted (see also § 7.1.2). The molecularexcitation and radiative transfer in the line is then computed ateach position in the envelope, using a temperature and densityprofile constrained by other data. The resulting sky brightnessdistribution is convolved with the
Herschel beam profile. Theseinner and outer abundances are then adjusted until best agree-ment with observational data is reached. The velocity structurecan be a constant (turbulent) broadening as function of positionin the envelope, and / or some other function such as infall or ex-pansion (see Johnstone et al. 2010; Herpin et al. 2012, for exam-ples).For simplicity, spherically symmetric models are typicallyadopted but Herschel data show that often UV-heated cavitywalls need to be added in a 2D geometry (see Bruderer et al.2009; Visser et al. 2012; Lee et al. 2013, for examples)
Appendix F.0.2: Forward models
The forward modeling approach uses a chemical model as il-lustrated in Fig. 4 coupled with a physical model of the sourceto compute abundances as function of position and time. Theseare then fed into excitation and radiative transfer calculations toobtain the line intensity maps, which are finally convolved withthe beam. These chemistry models can be run either in steady-state or in a time-dependent mode at each position in the cloud.Best fit parameters in this case can be any parameter fed into thechemical network such as overall (initial) abundances or rates ofreactions. Often, the chemical evolution time (“age”) or cosmicray ionization rate are adjusted to obtain the best fit (see Dotyet al. 2006; Caselli et al. 2012, for examples). As noted above,often simplified chemistry networks are used for computationalspeed and to allow key processes to be identified.Some of the above models keep the physical structure of thesource fixed even when the chemistry evolves in time, while oth-ers let the physical structure evolve. Also, a parcel of gas oftenstays fixed at a given location in the source. An alternative ap-proach is to follow the chemistry of a parcel of gas as it fallsin from large radii to smaller distances and eventually enters the(forming) protostellar disk. In the context of WISH, such 2Dmodels have been developed and applied to water by Visser et al.(2009, 2011) and for deuterated water by Furuya et al. (2017).
Appendix G: Oxygen budget
This section summarizes our current understanding of the oxy-gen budget in various regions of the interstellar medium that arepart of the star- and planet formation cycle. The di ff erent oxygenreservoirs are summarized in Table 3 and visualized in percent-ages in Figure 32 in the main text. The following subsectionsprovide more background information on the numbers in this ta-ble and figure. The numbers in the text plus their uncertaintiesare generally as stated in the original papers, with added textto motivate the uncertainties. The number of significant digitstherefore varies and should not be taken to reflect the actual un- Article number, page 55 of 58 & A proofs: manuscript no. wish_final_Jan2021 certainties. Very small numbers have been rounded o ff to ( < )1ppm. Appendix G.1: Overall and refractory oxygen budget
The overall [O] / [H] abundance of oxygen in all forms in the ISMis taken to be 5 . × − or 575 ppm, as measured for early B-type stars (Przybilla et al. 2008). This number, denoted in shortas [O], is somewhat higher than the solar oxygen abundance of4 . × − or 490 ppm (Asplund et al. 2009; Grevesse et al. 2010).A recent compilation by Lodders (2019) of solar and meteoriticabundances gives a recommended value for the present solar sys-tem of 537 ppm (their Table 8), with values ranging from 512–660 ppm. The solar and solar system abundances presumablyreflect the ISM composition as it was 4.6 Gyr ago, whereas thehigher Przybilla et al. (2008) value is representative of hotter andtherefore younger stars. Indeed, the di ff erence in abundances isconsistent with galactic chemical evolution models (Chiappiniet al. 2003). In the following analysis, [O] =
575 ppm is used, butother studies (e.g., Pontoppidan et al. 2014) use the lower [O]abundance in their analysis of the oxygen budget in protoplane-tary disks, decreasing the need for unidentified forms of oxygenby nearly 100 ppm.Oxygen can be partitioned into gas, ice and dust. The gas andice are considered as “volatiles”, the dust as “refractories”, withthe di ff erence related to the sublimation temperature of the ma-terial. The boundary between the two categories is traditionallytaken to be around 600 K, so that refractories include at leastall the silicates, metals (as chemically defined), and metal ox-ides. The fraction of oxygen locked up in refractory dust, [O refr ],then comes from counting the amount of [Mg], [Si] and [Fe] inthe ISM, together denoted as [M]. Here the analysis of Whit-tet (2010) is followed. Depending on the type of mineral (e.g.,olivines or pyroxenes or metal oxides), the oxydation ratios dif-fer slightly. The maximum is [O refr ] / [M] = refr ] / [M] = refr ] = r ] =
140 ppm is assumed here. This number does not includeany oxygen locked up in refractory organics or any other typesof refractory material.
Appendix G.2: Diffuse clouds Di ff use interstellar clouds with extinctions A V < ) as well asthe dominant form of oxygen, neutral atomic oxygen O, can bemeasured directly with high accuracy through ultraviolet spec-troscopy toward background early-type stars. In this paper, theatomic oxygen abundance of O / [H] = . × − or 320 ppm isused, as measured by Meyer et al. (1998) with HST-GHRS for13 lines of sight in the local ISM. The uncertainty of ∼
10% isdue to the spread in measurements as well as a small uncertaintyin the oscillator strength of the optically thin O I 1356 Å line. Asubsequent survey of the 1356 Å line in the Galactic disk withHST-STIS by Cartledge et al. (2004) finds atomic O abundancesranging from 390 ppm for the lowest density clouds to 284 ppmin higher density regions, consistent with the slight decrease withdensity seen by Jenkins (2009) in di ff use clouds.Oxygen-containing molecules such as H O and O haveabundances in di ff use clouds that are at least 3 to 4 orders ofmagnitude lower than that of atomic oxygen (e.g., Spaans et al. 1998; Flagey et al. 2013). Gaseous CO (She ff er et al. 2008) andices (Poteet et al. 2015) are also negligible in di ff use molecularclouds. Thus 320 ppm is taken as the amount of volatile oxygen,[O vol ] that can cycle between gas and ice in dense clouds.For our adopted [O] abundance of 575 ppm, the sum of theidentified refractory and volatile oxygen is 140 + = ff use clouds are consistent withthe detailed analysis of a single well-studied line of sight, thedi ff use molecular cloud toward ζ Oph with A V ≈ refr ] in silicates to be 126 ±
45 ppm,with an additional ≤
19 ppm in iron oxides such as Fe O . Takingthe measured atomic O toward ζ Oph at 307 ±
30 ppm (Jenkins2009), together with their upper limits on H O ice ( ≤ − < Appendix G.3: Dense clouds
Many oxygen-bearing species in gas, ice and dust can be mea-sured directly toward bright infrared sources in or behind denseclouds. Because of their rising spectrum with wavelength, pro-tostars embedded in their natal envelope are particularly suitablefor high S / N spectra. As for di ff use clouds, the advantage of ab-sorption line studies is that all species are measured for the samepencil beam line of sight. Summaries of abundances and oxygenbudget toward protostars based on data from ISO , Spitzer andground-based instruments are presented in Gibb et al. (2004);Öberg et al. (2011a) and Boogert et al. (2015).
Appendix G.3.1: Low mass protostars
For low-mass protostars, the bulk of the envelope is cold so mostof the oxygen budget is in ice and dust. The numbers listed inTable 3 are taken from the Boogert et al. (2015) compilation.The median ice abundances in their Table 2 are listed both withrespect to water ice and with respect to hydrogen. The latter val-ues require a determination of the total hydrogen column N H along the line of sight. This column is obtained either from themeasured optical depth of the silicate feature τ . µ m or from thenear-IR extinction A K , or both, with relations between these ob-servables and the total hydrogen column density benchmarkedfor clouds in which both can be measured (Vuong et al. 2003;Boogert et al. 2013). The amount of oxygen in silicates is takento be the same as for di ff use clouds, 140 ppm (Whittet 2010).The bulk of the oxygen in ices is in H O, CO and CO ices.Of these three species, CO is the most volatile one and can showsignificant variation from cloud to cloud. It is the only majorice species that is not formed on the grains but frozen out fromthe gas, with an ice abundance that depends on temperature anddensity. However, when the CO ice abundance is low, its gas-phase abundance is high, so this does not a ff ect the total volatileoxygen budget. Taken together, the amount of oxygen measuredin ices sums up to only 76 ppm, or 13% of [O].The gas-phase CO column toward several low-mass proto-stars has been measured through high spectral resolution infraredobservations of CO and its isotopologs by Smith et al. (2015).Their inferred CO ice / (gas + ice) abundances are generally muchless than 20% (see their Figure 11). With CO ice at 10 ppm, Article number, page 56 of 58.F. van Dishoeck et al.: WISH: physics and chemistry from cloud to disks probed by
Herschel spectroscopy this implies a gas-phase CO abundance of at least 40 ppm. An-other independent measurement of gaseous CO in cold cloudshas been obtained by observing the infrared absorption lines ofboth CO and H (Lacy et al. 2017). Their inferred CO / H is1 . × − with a variation of ∼
30% between lines of sight, whichwould correspond to CO locking up 83 ppm with some spread.Based on these two studies, we take gaseous CO to account for100 ppm with a 50% uncertainty.No direct observational limits exist for gaseous atomic oxy-gen or other oxygen-bearing molecules, but they should be neg-ligible according to dark cloud models (McElroy et al. 2013).Also, both millimeter interferometry and
Herschel -WISH + stud-ies of H O emission lines show that the warm water abundancein the inner regions of low-mass protostars is generally low,much less than 40 ppm (Visser et al. 2013; Persson et al. 2016,their Figure 11), and that O is negligible (Yıldız et al. 2013b)(see § 7.4).Taken together, the silicate dust, ices and gas account foronly 140 + + + =
356 ppm of oxygen at most. Thisleaves at least 219 ppm or 38% of the oxygen budget unac-counted for. If all in UDO, this would imply a doubling of itsabundance compared with di ff use clouds. Appendix G.3.2: High-mass protostars
The abundances measured toward high-mass protostars listed inTable 3 are also taken from Boogert et al. (2015). Overall, thepattern is similar as for low-mass protostars, except that the iceabundances are even lower. In particular, the CO ice abundance,which counts double for oxygen, is a factor of 2–3 lower in high-mass regions. Taken together, the total amount of oxygen in icesis only 47 ppm.The envelopes around high-mass protostars are more mas-sive and warmer due to their higher luminosities, so gas-phasecolumn densities are larger than for low-mass sources. The COabundance in dense warm gas has been measured directly byLacy et al. (1994) for once source through infrared absorptionlines of both CO and H to be CO / H = . × − , correspond-ing to 130 ppm. We take CO to contain 100 ppm more generally.For high-mass protostars, CO and H O gas have been measureddirectly through infrared absorption studies with
IS O for thesame lines of sight as for the ices. CO gas has been detectedwith typical gas / ice ratios of 0.01–0.08 so it contributes negligi-bly (Boonman et al. 2003b; van Dishoeck 1998). In contrast, theH O gas / ice ratio derived from infrared absorption is more vari-able, with gas / ice ratios ranging from 0.01-2.2 (van Dishoeck &Helmich 1996; van Dishoeck 1998). With H O ice accountingfor 31 ppm, this gives H O gas at most 62 ppm, about 60% ofthat of CO. This number is consistent with the measured gaseousH O column densities relative to those of hot H in the inner re-gion (Boonman & van Dishoeck 2003), with H derived fromCO infrared absorption lines obtained with the CFHT-FTS byMitchell et al. (1990) assuming CO / H = × − . Interestingly,the corresponding ratio H O / CO ≈ O data obtained with SOFIA-EXES and VLT-CRIRES. Assuming that atomic O and O arenegligible, this gives a total of 162 ppm accounted for in the gas.Taken together, the dust, ice and gas account for 140 + + =
349 ppm, leaving 226 ppm for UDO (40% of oxygen).This is very similar to the case for low-mass protostars, eventhough very di ff erent techniques and instruments are involvedfor measuring the gas abundances. Appendix G.4: Shocks
Shocks have the advantage that temperatures are so high that allvolatile oxygen should be returned to the gas phase. The refrac-tory silicate part is again taken to be 140 ppm, and should beviewed as an upper limit since some silicate dust could be sput-tered by the shocks, as evidenced by high gas-phase abundancesof SiO (Guilloteau et al. 1992).The most direct measurements of the gaseous H O and COabundances come from
Herschel -PACS emission line observa-tions of dense shocks at positions o ff set from the protostars forwhich Spitzer data on H are available (Neufeld et al. 2014).There is some degeneracy in the fit, however, so not all columndensities can be determined independently. For the shock associ-ated with the NGC 2071 outflow, H O / CO = / H = . × − (that is, assuming most available carbonin CO). This would imply that CO accounts for 160 ppm andH O for 128 ppm.As discussed in § 5 and Table 2,
Herschel
WISH observa-tions of H O in shocks associated with low-mass protostars gen-erally give low abundances, down to 1 ppm. The best determinedcase is that of HH 54 for which H is obtained from spectrally-resolved VLT-VISIR data, finding H O / H < . × − , lessthan 7 ppm (Santangelo et al. 2014a). Similarly, Kristensen et al.(2017b) find H O / CO ∼ Herschel -HIFI H O and CO 16-15 data. Finally, fitting bothCO
Herschel -PACS and H Spitzer data simultaneously in anumber of shocks associated with low-mass protostars givesCO / H ∼ × − with values up to 10 − , corresponding to10–50 ppm (Dionatos et al. 2013, 2018).OH has been measured in some outflows associated withlow- and high-mass sources with Herschel but remains a smallcontributor to the oxygen budget (Wampfler et al. 2011) (see§ 5.2). The same holds for O (Melnick et al. 2011) (see § 7.4).This leaves atomic oxygen as the only other plausiblegaseous oxygen reservoir. The few existing SOFIA-GREAT dataof spectrally resolved [O I] 63 µ m lines imply low abundances,however, with atomic O accounting for at most 15% of thevolatile oxygen (Kristensen et al. 2017a). For the NGC 1333IRAS 4A shock, atomic O is at most 30 ppm.In summary, for shocks studied with Herschel and taking themaximum CO abundance at 160 ppm, as assumed by Neufeldet al. (2014) for NGC 2071, gives a total oxygen budget in sil-icate dust and gas of 140 + + + =
458 ppm, leav-ing 117 in unidentified form. Taking instead CO at 100 ppmwould give 140 + + + =
350 ppm, leaving 225 ppmfor UDO. For shocks associated with low-mass protostars, theCO and H O abundances may be up to an order of magnitudelower. Both numbers for UDO for NGC 2071 are remarkablysimilar to the range found for di ff use clouds and protostars, eventhough gas and dust in shocks experience much higher temper-atures. Apparently, the UDO material is refractory enough thatit is not sublimated or sputtered under typical shock conditionswith 300–700 K gas probed with Herschel data.There is some evidence that for the hottest gas, > O. Al-though the best fit to the
Herschel data of NGC 2071 givesH O / CO <
1, an acceptable fit can be found with a some-what higher temperature T w above which H O becomes abun-dant (1000 K rather than 300 K), which results in H O / CO = O / H = (0 . − × − using evenhigher excitation H O lines coming from hotter gas observedwith
Spitzer . The upper range of this study would account for
Article number, page 57 of 58 & A proofs: manuscript no. wish_final_Jan2021 most of the volatile oxygen and leave no room for UDO. Simi-larly, Goicoechea et al. (2015) identify a hotter, ∼ J CO and H O Her-schel lines with H O / CO ∼ . O / H (cid:38) × − andCO / H (cid:38) . × − . More such studies using mid-infrared linesprobed with JWST are warranted. Appendix G.5: Comets
The
Rosetta mission to comet 67P / Churyumov-Gerasimenkohas provided a unique opportunity to determine the compositionof one icy planetesimal in exquisite detail. Depending on its for-mation location and history in the solar system, a significant frac-tion of its composition could be inherited from the dense cloudphase (Drozdovskaya et al. 2019). The bulk abundances of themany molecules measured by the