emission observed, and in particular H O emission, is a peculiarity of the environments of class 0

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1 THE ASTROPHYSICAL JOURNAL, 555:4È57, 1 July 1 ( 1. The American Astronomical Society. All rights reserved. Printed in U.S.A. V FAR-INFRARED INVESTIGATION OF CLASS SOURCES: LINE COOLING1 TERESA GIANNINI, BRUNELLA NISINI, AND DARIO LORENZETTI Osservatorio Astronomico di Roma, via Frascati 33, I-4 Monte Porzio Catone, Italy Received December ; accepted 1 March 5 ABSTRACT We have investigated with the Long Wavelength Spectrometer (LWS) of the Infrared Space Observatory (ISO) the far-infrared spectra (43È197 km) of a sample of 17 class sources and their associated outñows. In addition to [O I] 63km, the pure rotational lines of abundant molecules such as CO, H O, and OH are frequently observed in these sources, at variance with more evolved young stellar objects. We found, in agreement with previous studies conducted on individual sources, that the molecular line excitation arises from small regions, with typical sizes of 1~9 sr, of warm ( \ T \ K) and dense gas (14 \ n \ 17 cm~3), compressed after the passage of shocks. In particular, we found slow, nondissociative shocks H as the main mechanism at the origin of the molecular gas heating, while the bulk of the [O I] 63km line emission is due to the dissociative J-shock component arising from the Mach disk at the head of the protostellar jet, as testiðed by the fact that this line emission happens to be a good tracer of the source mass-loss rate. Large abundances of gas-phase H O are commonly estimated, with values that appear to be correlated with the gas temperature. The total far-infrared (FIR) line cooling L \ L (O I) ] L (CO) ] L (H O) ] L (OH), which amounts to D1~ to 1~1 L, is roughly equal to FIR _ the outñow kinetic luminosity as estimated by means of millimeter molecular mapping. This circumstance demonstrates that the FIR line cooling can be a valid direct measure of the power deposited in the outñow, not a ected by geometrical or opacity problems like the determination of L or by extinc- kin tion problems like the near-infrared shocked H emission. We Ðnally remark that the strong molecular emission observed, and in particular H O emission, is a peculiarity of the environments of class sources. The present analysis shows that the ratio between FIR molecular line luminosity and bolometric luminosity (L /L ) is always larger than D1~3 in class objects. We suggest that this parameter mol bol could be used as a further criterion for identifying future class candidates. Subject headings: infrared: ISM È ISM: jets and outñows È ISM: molecules È stars: formation È stars: mass loss On-line material: machine-readable table 1. INTRODUCTION Class sources are highly obscured objects showing a strong submillimeter continuum emission compared with their bolometric luminosity (Andre, Ward-Thompson, & Barsony 1993). They are intrinsically very cold, with blackbody-like spectral energy distribution at dust temperatures of only È3 K. These characteristics, together with their inferred short lifetime of only about 14 yr, have led to the suggestion that class objects are very young protostars, which have already developed a hydrostatic core but still have an envelope mass larger than the already accumulated stellar mass (Andre & Montmerle 1994). Because of their short lifetime and their invisibility even at mid-infrared wavelengths, these sources are very difficult to detect. In fact, from the discovery of the class prototype VLA 163 by Andre et al. (1993), only about 4 class candidates have been so far identiðed (Andre, Ward- Thompson, & Barsony ). A large part of the known class sources have been indirectly discovered through the detection of their powerful bipolar outñows. It has been indeed well established that these objects drive strong Ñows, usually more energetic than those from class I sources of the same luminosity (Bontemps et al. 1996; Saraceno et al. 1996). In these works it has been suggested that this may be a direct consequence of their 1 Based on observations with Infrared Space Observatory (ISO), an ESA project with instruments funded by ESA Member States (especially the PI countries: France, Germany, the Netherlands, and the United Kingdom) and with the participation of ISAS and NASA. 4 larger mass accretion rates with respect to more evolved protostars. Since mass accretion and ejection are strictly related, the relatively large L /L in class sources can be kin bol reasonably due to a larger rate of mass loss with respect to older sources. This property demonstrates that the study of outñows can be a very powerful way to retrieve information indirectly on the elusive exciting source. The Ñows from class sources expand in a very dense medium (n [ 15 cm~3) where the inñuence of the magnetic Ðeld on the streaming gas may be high. Therefore, jets from class sources are expected to drive strong MHD shocks (continuous or C-shocks; Draine 198) where the medium does not reach temperatures high enough to dissociate the molecules. In fact, outñows from these early sources are on average more likely associated to H jets than to optical HH objects (Eislo ffel et al. ), which trace a lowexcitation ionized medium. H rovibrational lines observable from the ground can, however, trace only a fraction of the shocked gas; a signiðcant contribution to the cooling is also expected from pure rotational H,HO, and high-j CO transitions, emitting in the mid- and far-infrared (FIR) parts of the spectrum, whose relative importance is a strong function of the shock parameters (e.g., Kaufman & Neufeld 1996). At the same time, the high velocity attained at the head of the jets can be larger than the Alfve n ion velocity, thus dissociative J-shocks can also be produced along the Ñow. The best tracer of such shocks in dense environments is the [O I] 63km line, which represents the main cooling channel in the postshocked region.

2 FIR LINE COOLING IN CLASS SOURCES 41 From the above arguments, it appears clear that FIR spectroscopy is fundamental for understanding the nature of the mass accretion/ejection in these sources and to study the chemical modiðcations induced by the formation of the new protostar in the surrounding medium. Detailed studies, conducted on the FIR spectra obtained by the Infrared Space Observatory (ISO) on few class sources, have indeed demonstrated the richness of these spectra, characterized by strong emission lines mainly from O I, CO, and H O (e.g., Nisini et al. 1999a, 1999b; Benedettini et al. ). Motivated by these results, we have undertaken a systematic analysis of the FIR spectra of all the class sources observed with the Long Wavelength Spectrometer (LWS) of ISO, with the aim of deðning the global properties of this class based on their FIR line emission. The layout of the present paper is the following. We deðne our sample in and describe the observations and the data reduction techniques in 3. In 4 and 5 we give the results and the line emission analysis. In 6 we discuss the origin of the observed emission and investigate the energy budget. Our conclusions are summarized in 7.. DEFINITION OF THE SAMPLE The investigated sample has been selected on the basis of the list of the known class sources compiled by Andre et al. (). Out of the 4 reported objects, 1 sources were observed, along with their associated outñow, during the ISO mission (Kessler et al. 1996) with the LWS (Clegg et al. 1996). We include in our study all these targets, with the exception of four objects, namely, NGC 4, FIR 3, FIR 5, and VLA 163, whose spectra show a clear contamination by strong gas emission from nearby H II regions and photodissociation regions (PDRs), which are encompassed by the relatively large LWS beam (Ðeld of view [FOV] of D8A; Giannini et al. ), and SSV 13b, which is a lowluminosity unresolvable companion of the source SSV 13 (Molinari et al. ). The remaining 17 sources are listed in Table 1 (identiðcation number in col. [1] and source name in col. []), along with their coordinates (cols. [3] and [4]) and some useful information, such as the bolometric luminosity (L ; col. [5]) and distance (D; col. [6]). We note that they are bol all low-luminosity sources (L ¹ 75 L ), located in nearby star-forming regions (D ¹ 75 bol pc). These _ conditions guarantee both to investigate low-mass objects and to minimize distance-dependent e ects. In more detail, our Ðnal sample is constituted by six objects, whose LWS spectra have been recently published in the literature, namely, L1448ÈIRS 3, L1448-MM (Nisini et al. 1999b, ), HH 4ÈMMS, HH 5ÈMMS (Benedettini et al. ), IRAS 1693[4 (Ceccarelli et al. 1998), and B335 (Nisini et al. 1999a), as well as 11 objects (along with their outñow) presented here for the Ðrst time. In fact, an investigation of the water emission from NGC 1333 IRAS and IRAS 4 was done by Ceccarelli et al. (1999), but, to give more homogeneous and complete information regarding all the line emission in the FIR, we reduced and analyzed these spectra from scratch. 3. OBSERVATIONS AND DATA REDUCTION The LWS data were retrieved from the public ISO Data Archive. Together with the on-source observations, we investigated all the available spectra covering the outñow lobes, in order to study the total emission of the whole source outñow system. Table shows the journal of observations. We report here the source names with the relative pointed coordinates3 (cols. [1]È[3]), the observer code (col. [4]), the observation date (col. [5]), the total dedicated time (TDT; col. [6]), and the number of scans that compose the observation (n ; col. [7]). Most of the objects belong to scan the guaranteed time program PSARA and thus constitute a homogeneous sample in which the TDTs of the on-source and outñow observations are similar, thus allowing a meaningful comparative analysis of the emission along all the outñow extent. In addition to these data, we included other us.html. 3 Observations of the outñow positions are identiðed with the source name followed by red ÏÏ or blue ÏÏ for the redshifted and blueshifted outñow lobes, respectively. TABLE 1 OUR SAMPLE OF CLASS OBJECTS L D bol IdentiÐcation Number Source a(.) d(.) (L ) (pc) _ (1) () (3) (4) (5) (6) 1... L1448-IRS L1448-MM NGC 1333ÈIRAS NGC 1333ÈIRAS IRAS 38] HH 11ÈMM L L1641ÈVLA [ HH 4ÈMMS [ HH 5ÈMMS [ IRAS 1693[ [ L483-MM [ IRAS 1873] L73-MM B L1157-MM CepE-MM NOTE.ÈFrom Andre et al.. Units of right ascension are hours, minutes, and seconds, and units of declination are degrees, arcminutes, and arcseconds.

3 4 GIANNINI, NISINI, & LORENZETTI Vol. 555 TABLE JOURNAL OF OBSERVATIONS Source a(.) d(.) Observer Code Date (s) n scan (1) () (3) (4) (5) (6) (7) NGC 1333ÈIRAS PSARA 1998 Feb NGC 1333ÈIRAS red N PSARA 1998 Feb NGC 1333ÈIRAS blue S PSARA 1998 Feb NGC 1333ÈIRAS red E PSARA 1998 Feb NGC 1333ÈIRAS blue W PSARA 1998 Feb NGC 1333ÈIRAS PSARA 1998 Mar NGC 1333ÈIRAS 4 red PSARA 1998 Mar NGC 1333ÈIRAS 4 blue PSARA 1998 Mar IRAS 38] PSARA 1997 Aug IRAS 38]335 blue PSARA 1997 Aug IRAS 38]335 red PSARA 1997 Aug HH 11ÈMM JCERNICH 1997 Aug HH 11ÈMM blue MSMITH 1997 Sep HH 11ÈMM red MSMITH 1997 Sep L EVDISHOE 1997 Sep L1641ÈVLA [ PSARA 1997 Oct HH1C [ PSARA 1997 Oct HHC [ PSARA 1997 Oct L483-MM [ PSARA 1997 Mar IRAS 1873] PSARA 1997 Mar L73-MM PSARA 1996 Oct L73 Ñowa PSARA 1996 Oct L1157-MM RBACHILL 1997 Feb L1157-MM blue RBACHILL 1997 Feb L1157-MM red RBACHILL 1996 Aug CepE-MM blue JEISLOEF 1997 Jun CepE-MM red JEISLOEF 1997 Jun NOTE.ÈUnits of right ascension are hours, minutes, and seconds, and units of declination are degrees, arcminutes, and arcseconds. a Raster map of Ðve pointings along the Ñow. Raster parameters: center on L73-MM; P.A. \ 16 ; spacing between each position is 1A. The total dedicated time and the number of scans refer to each raster observation. TDT observations belonging to di erent programs that have followed di erent observational guidelines: this is the case of the CepE-MM observation, in which the class source was not directly observed but was encompassed by both of the outñow lobe pointings that partially overlap, or the case of HH 11ÈMM, whose on-source and outñow measurements, pertaining to two di erent programs, were performed with n values that di er by about a factor of 4. scan All the spectra were obtained with the LWS AOT1 full grating scan mode (i.e., 43È197 km wavelength range, resolution j/*j D ) adopting the fast scanning ÏÏ option and an oversampling of 4 times the spectral resolution element. The raw data were reduced with the O -Line Processing (OLP), Version 7. The Ñux calibration was derived from the Uranus spectrum (accuracy of about 3%; Swinyard et al. 1996), while the uncertainty in the wavelength calibration is about 5% of the resolution element (.9 and.6 km for the short-wavelength [SW: 43È9 km] and the long-wavelength [LW: 9È197 km] ranges, respectively). Postpipeline processing, which consists in glitch removal, averaging of the di erent spectral scans, and interference fringe correction, was performed using the ISO Spectral Analysis Package (ISAP), Version 1.6a.4 4 The ISO Spectral Analysis Package is a joint development by the LWS and SWS Instrument Teams and Data Centers. Contributing institutes are CESR, IAS, IPAC, MPE, RAL, and SRON. 4. RESULTS As an example of the commonly observed FIR spectra, we show in Figures 1aÈ1c the spectra of NGC 1333ÈIRAS 4, CepE-MM (blue lobe), and HH 11ÈMM (continuum subtracted by means of Ðrst- or second-order polynomial Ðtting), which can be considered as representative of the whole sample. The most prominent feature is the copious presence, both on source and along the outñow lobes, of pure rotational lines of CO, H O, and OH (detection rate of 94%, 7.5%, and 41%, respectively), which are rarely observed in more evolved objects (see, e.g., Giannini et al. 1999; Lorenzetti et al. ). In addition, the atomic [O I] 63 km and [C II] 158 km lines were observed everywhere, while the [O I] 145 km line was discovered in 65% of the sources. In Figure we show some examples of the [O I] 63 km line detected at various intensity levels. No trace of higher ionization lines was found. The line Ñuxes were measured from the defringed, single-detector spectra, by Ðtting the spectral proðles with a single or double (in case of blending) Gaussian function. For all the lines, the FWHM is comparable with the instrumental resolution element width except in the case of blending and for the OH lines, whose FWHMs are broader because of the unresolved contribution of the Ðne-structure lines due to the " doubling splitting of the levels. For all the lines the distance between observed and rest wavelength is comparable with onequarter of the spectral resolution element. Table 3 shows,

4 No. 1, 1 FIR LINE COOLING IN CLASS SOURCES 43 Fig. 1a Fig. 1b Fig. 1c FIG. 1.È(a) LWS continuum-subtracted spectrum of NGC 1333ÈIRAS 4. Lines detected at intensity level greater than p are indicated. (b) As in(a), but for the blue lobe of CepE-MM. (c) Portions of the LWS continuum-subtracted spectrum of HH 11ÈMM. for each investigated region, the list of rest vacuum wavelengths (Ðrst column) of the identiðed lines, along with the integrated Ñuxes with the associated 1 p statistical uncertainty, estimated from the rms in the adjacent baseline (second column). Whenever deblending was impossible, we have indicated the lines that contribute to the whole Ñux. The reported values refer to detections with a signal-tonoise ratio (S/N) º3, unless explicitly indicated with a Ñag (which corresponds to \ S/N \ 3). In all the cases in which at least one CO line was detected, a circumstance that allows us to retrieve indications on temperature and density (see 5), we have also calculated the 3 p upper limits for the Ñuxes of the ortho-h O 179 km and OH 119 km lines, which are the lines closest to their fundamental levels in the LWS range, and which have been subsequently used to estimate the line cooling of H O and OH. In a few cases other upper limits that were useful for the analysis are reported in the table. From a Ðrst inspection of the line Ñuxes we note that (1) the [C II] intensity level appears quite similar in most of the pointed positions (F[C II] B 1~19 Wcm~), () a larger variability among the various sources is observed in the [O I] 63km emission, and (3) roughly the same molecular lines, with similar intensity levels, are seen both on source and along the outñow lobes. The implications of these points will be discussed later. As we have already anticipated, water emission from IRAS and IRAS 4 was discussed in a dedicated paper by Ceccarelli at al. (1999). The comparison with the newly extracted Ñuxes shows a general agreement as far as the IRAS spectrum is concerned (their unique 3 p detection is the ortho-h O at 179 km line, which we consider here as an upper limit). On the contrary, we found in the IRAS 4 spectrum, in addition to their 1 water lines, three other lines (at 8.3, , and km) with S/N [ 3 and four lines (at least, not considering possibly blended lines) with \ S/

5 44 GIANNINI, NISINI, & LORENZETTI FIG..ÈObserved [O I]63 km line toward class sources (continuum subtracted) N \ 3, some of them emitted from levels at high-excitation energy (T [ 5 K). As we will see in the following, this ex evidence is in favor of a gas temperature higher than their estimate of D K. 5. LINE EMISSION ANALYSIS The physical parameters of the gas were determined by analyzing the molecular emission from CO, H O, and OH. For this purpose the radiation transfer was solved simultaneously with the statistical equilibrium equations, under the large velocity gradient (LVG) conditions in plane-parallel geometry. We developed the LVG code for the Ðrst 45 rotational levels for both ortho- and para-h O (Nisini et al. 1999b) and for 6 and levels for CO and OH, respectively (Giannini et al. 1999). In the water model we assume an ortho-to-para ratio of 3. In the adopted approximation the free parameters, which depend on the line ratios, are the gas kinetic temperature, its density, and the ratio between column density and intrinsic line width (N/*V ). Having these parameters Ðxed, the angular extent of the emitting region depends only on the absolute line intensities. In principle, independent determinations of all the parameters could be derived by modeling the emission of each species. In fact, CO lines are optically thin and, as such, can constrain only temperature and density, while H O and OH lines, having in general opacities of the order of unities or more, depend also on the ratio N/*V (Nisini et al. 1999b), but they are usually too few to constrain all the parameters. This difficulty can be circumvented by noting (see Table 3) that sources exhibiting higher (lower) J CO transitions are those where also the H O lines come up from levels at higher (lower) excitation. This can be clearly seen in the spectrum of IRAS 4 (Fig. 1a), which exhibits simultaneously the highest J CO (J \ 9) and H O (at 65. and 8.3 km, with J \ up 6) observed up lines. Given this circumstance, it is reasonable up to assume that all the molecular species are emitted from the same gas component. This allows us (1) to derive temperature and density from the CO Ðt and () to use the H O and OH lines to estimate both the respective N/*V ratios (or absolute column densities, if the velocity dispersion is independently known) and the extendedness of the emitting area. The last step consists in deriving the CO column densities from the absolute Ñuxes and, once the CO abundance with respect to H is assumed [we adopted the standard value X(CO) \ 1~4], in calculating also X(H O) and X(OH). Although the number of detected lines is not

6 TABLE 3 OBSERVED LINES NGC 1333ÈIRAS j (km) vac IDENTIFICATION On-Source Red N Blue S Red E Blue W [O I] 3P ] 3P 9.^ ^. 38.8^.3 38.^ ^ OH % ] % 1.^3.7a @,1@ 3@,3@ OH % ] % \1.5 \.1 \.1 \3. \1.8 3@,5@ 3@,3@ [O I] 3P ] 3P \ ^.7 3.5^ CO 17] ^ [C II] 3P ] 3P 1.^.6 8.4^1.7 1.^ ^ ^.8 3@ 1@ CO 16] 15 7.^. 4.7^ ^1..7^ CO 15] ^ ^ o-h O ] \6 \7 \6 7.8^.5 8.6^. NGC 1333ÈIRAS 4 j vac (km) IDENTIFICATION On-Source Red Blue o-h O4 3 ] 3 1.9^1.a [O I] 3P 1 ] 3P 4.3^ ^3.8.6^ *... o-h O6 5 ] *... OH % 3@,9@ ] % 3@,7@ 3.9^1.5a, b o-h O3 3 ] 1 3.7^1.5a *... p-h O3 31 ] 67.7*... o-h O3 3 ] ^1.5b *... p-h O5 4 ] *... OH % 1@,7@ ] % 1@,5@ 6.7^.5a, b o-h O3 1 ] ^3. 14.^5a c... o-h O4 3 ] ^ c... OH % 1@,1@ ] % 3@,3@ 1.4^ o-h O6 16 ] ^ p-h O6 6 ] ^.5a OH % 3@,7@ ] % 3@,5@ 1.4^ *... p-h O3 ] *... CO 9] 8 15.^.1b CO 8] 7 3.3^1.3a p-h O5 15 ] 4 4 \ CO 7] 6 3.8^1.4a o-h O5 6 ] ^ *... CO 6] *... o-h O5 14 ] ^.b *... p-h O ] CO 5] 4 4.4^1.7a c... o-h O 1 ] ^.5 5.4^1.8.6^ c... CO 4] 3 6.^ *... CO 3] *... o-h O4 14 ] ^.b 3.5^1.b \.6b c... CO ] 1 5.6^ c... OH % 3@,5@ ] % 3@,3@ 9.^ ^.6 \ o-h O4 3 ] 4 3 \ CO 1] 6.1^ p-h O4 4 ] ^ CO ] ^1.5.6^ o-h O4 3 ] ^1.5a c... o-h O3 3 ] 3 1 \ c... CO 19] 18 9.^1.6 \ p-h O3 13 ] 14.8^ c... CO 18] ^1. 3.5^ c... [O I] 3P ] 3P 1 \3.6 \ CO 17] ^ ^.7 3.5^ [C II] 3P 3@ ] 3P 1@ 14.7^ ^.8 13.^ c... CO 16] ^ ^.7 3.3^ c... OH % 1@,3@ ] % 1@,1@ 3.5^1.4a c... CO 15] ^ ^ c... o-h O3 3 ] ^ ^....

7 TABLE 3ÈContinued j vac (km) IDENTIFICATION On-Source Red Blue c... o-h O ] 1 1.^ ^. 6.1^ c... o-h O ] 5.5^ CO14] ^ ^.3... IRAS 38]335 j (km) vac IDENTIFICATION On-Source Blue Red [O I] 3P ] 3P 5.6^ ^.7 4.^1.9a OH % ] % \. \. \. 3@,5@ 3@,3@ [O I] 3P ] 3P....6^ CO 17] ^ [C II] 3P ] 3P 8.9^.5 1.4^.8 8.^1.5 3@ 1@ CO16] 15 1.^.4a.9^.9 7.5^ o-h O ] \.9 6.^. 8.5^.8 HH 11ÈMM j (km) vac IDENTIFICATION On-Source Blue Red [O I] 3P ] 3P 38.1^ ^.6 3.^ o-h O ] ^ o-h O4 ] 3 \ ^ OH % ] % \3.5 \3.1 \3.4 3@,5@ 3@,3@ CO] 19.7^ CO19] ^ c... CO 18] ^.9.3^.6.3^ c... [O I] 3P ] 3P \3 1.9^.6.^ CO17] ^.8 4.^.8 4.9^ [C II] 3P ] 3P 5.^1. 53.^ ^.9 3@ 1@ CO16] ^1.1 6.^ ^ CO15] ^.8 4.^.8 6.8^ o-h O ] 1 \1 6.^ ^ CO14] 13 \13 6.7^ ^1.6 L157 j (km) IDENTIFICATION On-Source vac [O I] 3P ] 3P 13.4^ [O I] 3P ] 3P 4.8^ [C II] 3P ] 3P 6.^. 3@ 1@ L1641ÈVLA 1 j (km) vac IDENTIFICATION On-Source HH 1C HH C [O I] 3P ] 3P 175.1^.4 13.^3. 15.^ CO5] 4 4.5^ c... o-h O ] 1.6^.7... \ c... CO 4] 3 4.4^ OH % ] % \3.6 \3.1 \.5 3@,5@ 3@,3@ p-h O3 ] 3 6.^ CO] ^1. 3.^ CO19] ^1.1.9^1.a c... CO 18] ^1. 3.8^.9 3.7^ c... [O I] 3P ] 3P 5.4^1. 3.9^.4 7.^ CO17] ^1. 4.^ ^ [C II] 3P ] 3P 83.5^ ^ ^.1 3@ 1@ CO16] ^1. 6.^ ^ CO15] ^. \ o-h O ] 1 \6.3 \ ^ CO14] 13 4.^.a \6.3 5.^.6a 46

8 TABLE 3ÈContinued L483-MM j (km) vac IDENTIFICATION On-Source [O I] 3P ] 3P 18.8^ *... CO 6] *... p-h O ] 1 8.6^1.b CO 5] 4 7.^ *... o-h O ] *... CO 4] 3 7.3^1.4b OH % ] % 5.9^1. 3@,5@ 3@,3@ CO ] ^1.7a CO 19] ^ [O I] 3P ] 3P 3.7^ [C II] 3P ] 3P 18.4^1.5 3@ 1@ CO 16] ^ o-h O ] \4.3 IRAS 1873]113 j (km) IDENTIFICATION On-Source vac [O I] 3P ] 3P 11.^ *... o-h O6 ] *... OH % ] % 14.3^3.5b 3@,9@ 3@,7@ o-h O3 ] 6.^3.1a p-h O3 ] 6.8^3.1a o-h O3 ] 8.^ OH % ] % 41.3^8.7 1@,1@ 3@,3@ o-h O6 ] ^ OH % ] % 48.3^6.7 3@,7@ 3@,5@ CO7] ^5.3a OH % ] % 7.1^6.6 1@,5@ 1@,3@ o-h O5 ] ^6.6a *... CO 6] *... o-h O5 ] ^5.4b *... p-h O ] CO 5] 4 17.^ c... o-h O ] 1 8.8^ c... CO 4] 3 1.1^ *... CO 3] *... o-h O4 ] ^3.1b CO ] 1 \ OH % ] % \13.5 3@,5@ 3@,3@ CO 1] 4.8^ p-h O4 ] 3 8.4^3.3a CO ] ^ CO 19] ^ c... CO 18] ^ c... [O I] 3P ] 3P 9.5^3.5a CO 17] 16.^ [C II] 3P ] 3P 6.7^3.7 3@ 1@ CO 16] ^ c... CO 15] ^ c... o-h O3 ] 1.8^ o-h O ] 1.^ CO 14] ^. L73-MM j (km) IDENTIFICATION On-Source Blue Blue 1 Red 1 Red vac [O I] 3P ] 3P 14.5^3.6 \8.3 \1.6 \ ^ OH % ] % \ \ @,5@ 3@,3@ [O I] 3P ] 3P 3.1^.4.3^ CO 15] ^ [C II] 3P ] 3P 9.5^ ^.8 11.^.8 13.^ ^1. 3@ 1@ 47

9 TABLE 3ÈContinued L73-MM j (km) vac IDENTIFICATION On-Source Blue Blue 1 Red 1 Red CO 16] 15 1.^ ^ o-h O ] \ \ L1157-MM j (km) vac IDENTIFICATION On-Source Blue Red [O I] 3P ] 3P 1.^.6 5.^ ^.a o-h O4 ] ^.9a OH % ] % 5.1^1..1^.7 \.9 3@,5@ 3@,3@ CO ] ^.5 1.1^ CO 19] ^.5 1.1^.5a p-h O3 ] ^ c... CO 18] 17.4^.6 1.^.35a c... [O I] 3P ] 3P... 1.^ CO 17] 16.1^.6 \ [C II] 3P ] 3P 3.8^.6 3.6^.7.7^.5 3@ 1@ CO 16] ^ CO 15] ^.a o-h O ] ^.4 13.^.3 5.^1. CepE-MM j (km) IDENTIFICATION Blue Red vac [O I] 3P ] 3P 13.6^ ^ *... o-h O6 ] *... OH % ] % 5.3^.a, b... 3@,9@ 3@,7@ o-h O3 ] 18.^5.5.^ OH % ] % 17.7^5.5.5^4. 1@,1@ 3@,3@ o-h O6 ] ^ OH % ] % 3.1^7..^7.7a 3@,7@ 3@,5@ 89.99*... p-h O3 ] *... CO 9] ^4.1b o-h O5 ] 5 \.6 \ *... CO 6] *... o-h O5 ] 4 4.9^4.5b *... p-h O ] CO 5] 4 4.9^ c... o-h O ] 1 14.^.1 9.^ c... CO 4] 3 5.^.1a o-h O4 ] ^.1 16.^ c... CO ] 1 6.4^.6a \ c... OH % ] % 1.9^.6 \1.5 3@,5@ 3@,3@ CO 1] 9.1^.7 7.1^ CO ] ^1. 7.6^ CO 19] ^ ^ p-h O3 ] 6.3^ c... CO 18] ^ ^ c... [O I] 3P ] 3P 6.1^ ^1.3a CO 17] 16.^ ^ [C II] 3P ] 3P 81.1^ ^1.5 3@ 1@ CO 16] 15.4^ ^ CO 15] ^.1 1.^ o-h O3 ] 19.4^.1 14.^ c... o-h O ] 1 9.^.1 7.9^ c... o-h O ] 7.5^ CO 14] ^3.5.9^. NOTE.ÈTable 3 is also available in machine-readable form in the electronic edition of the Astrophysical Journal. a S/N marginally less than 3. b Integrated Ñux of those lines marked with asterisks (impossible deblending). c Deblended lines.

10 FIR LINE COOLING IN CLASS SOURCES 49 large enough to apply the above described procedure to all the sources, we tried to derive temperature and density estimates whenever at least one CO line was observed. This happens for all the sources but L157, where only atomic emission is present. Well-restricted ranges of temperature and density can be e ectively obtained from CO Ðts only if the peak of the line Ñux distribution as a function of J is traced by the detected lines; this is the case of about half up of the 7 considered spectra. As an example, we plot in Figures 3aÈ3c the CO line Ñux distribution Ðts of IRAS 4, CepE-MM (blue lobe), and HH 11ÈMM: in the Ðrst two sources, where the emission is detected in several lines with high signal-to-noise ratios, temperature and density are well constrained, while in the HH 11ÈMM case the possible range of the parameters Ðtting the data is indicated by the dashed and solid lines. Temperatures and densities derived from our CO Ðts, together with the determinations retrieved from the literature for the six sources previously studied, are reported in Table 4 (cols. [] and [3]). We note that the temperatures we derive range from D (HH 11ÈMM) to more than 1 K (NGC 1333ÈIRAS 4, L483-MM, IRAS 1873]113, CepE-MM); this result is in general agreement with the few hundreds of kelvin already found in several sources (i.e., B335 and HH 5ÈMMS) and somehow enlarges the occurrence of higher temperatures, indicating that the value of 1 K found in L1448-MM is not an Fig. 3a Fig. 3b Fig. 3c FIG. 3.È(a) CO line Ñuxes measured toward NGC 1333ÈIRAS 4 as a function of the rotational quantum number J. The error bars indicate the 1p uncertainties. The solid line is the best Ðt to the data, whose parameters are indicated as well. (b)asin(a), but for the blue lobe up of CepE-MM. (c)asin(a), but for HH 11ÈMM. The downward arrow refers to a 3 p upper limit. The dashed and solid lines give the two extreme model Ðts to the data.

11 5 GIANNINI, NISINI, & LORENZETTI Vol. 555 TABLE 4 PARAMETERS FROM THE LVG FIT Temperature Density N(CO) N(H O) Size X(H O) Source (K) (15 cm~3) (116 cm~) (116 cm~) (1~1 sr) (1~4) (1) () (3) (4) (5) (6) (7) L1448ÈIRS 3a... 5È6È L1448-MMa L1448 red lobea... 6È65.1È.3 D1 D1 NGC 1333ÈIRAS b NGC 1333ÈIRAS 4 red lobe... 7È1.3È1. 1È4 17b D.6È1.7 HH 11ÈMM... 35È95È HH 11ÈMM blue lobe... 5È65È L1641ÈVLA È95È HH1C... 65È18.È HH 4ÈMMSc... 65È14 5È HH 5ÈMMSc... 15È55 5È8 9È6.9È.3È.1 IRAS 1693[4d... È4 4È5.5È L483-MM... 85È18 È IRAS 1873] e B335f L1157-MM... 5È18.È L1157 blue lobe... 35È8 5È CepE-MM blue lobe e CepE-MM red lobe e a Nisini et al.. b *v \ km s~1 is assumed. c Benedettini et al.. d Ceccarelli et al e *v \ 1 km s~1 is assumed. f Nisini et al. 1999a. isolated case. On the contrary, our density range is inside the limits (from 14 to a few times 16 cm~3) already established in previous studies. In all the remaining CO spectra, showing very few (but at least one) lines (all with 15 ¹ J ¹ 17), only a poor estimate of temperature up (5 \ T \ K) and density (14 \ n \ 17 cm~3) can be given by considering all the possible H conðgurations (e.g., McKee et al. 198) in which the observed lines constitute the peak of the line Ñux distribution. Detailed H O Ðts could be made only for three sources, namely, IRAS 4 (on-source and red lobe), IRAS 1873]113, and CepE-MM (blue and red lobe), where several lines from both ortho and para forms were detected. As an example, the H O Ðt of IRAS 4 is shown in Figure 4. The indicated temperature and density (T \ 18 K and n \ 3 ] 14 cm~3) were Ðxed to the values derived from CO, H while the N(H O)/*V ratio is optimized in order to reproduce the observations. The ortho and para predicted LVG model for NGC 1333ÈIRAS 4. The and predicted intensities are represented by triangles and squares, FIG. 4.ÈH O ortho-h O para-h O respectively, while the observations with the relative 1 p uncertainties are represented by open circles. Downward arrows indicate 3 p upper limits.

12 No. 1, 1 FIR LINE COOLING IN CLASS SOURCES 51 TABLE 5 FIR LINE COOLING L [O I] L [C II] L (CO) L (H O) L (OH) L (FIR)a Source (1~ L _ ) (1~ L _ ) (1~ L _ ) (1~ L _ ) (1~ L _ ) (1~ L _ ) (1) () (3) (4) (5) (6) (7) L1448ÈIRS 3b c \ L1448-MMb c.5È L1448 red lobeb c... NGC 1333ÈIRAS È4.4 \1.3 \ NGC 1333ÈIRAS N È4. \1.4 \.5 NGC 1333ÈIRAS S È3.8 \1. \.5 NGC 1333ÈIRAS E È4.9.8È1.7 \.7 NGC 1333ÈIRAS W È1.3.9È1.3 \.4 NGC 1333ÈIRAS NGC 1333ÈIRAS 4 NE È9. 4.9È6.1.È.4 NGC 1333ÈIRAS 4 SW È3.3.7È1.3 \.4 IRAS 38] È.6 \.5 \ IRAS 38]335 blue lobe è1..5è.8 \.3 IRAS 38]335 red lobe è.5.7è1.1 \.3 HH 11ÈMM È.9 \.9 \ HH11 blue lobe È.9.9È1.5 \.4 HH11 red lobe È3.4.4È1. \.4 L L1641ÈVLA È È4.6 \ HH1C È5.6 \. \.6 HHC È4.3.9È1.9 \.8 HH 4ÈMMSd È1.6 \ HH 5ÈMMSd È3.7. \ IRAS 1693[4e È IRAS 1693[4 Ñowe f L483-MM È1.5 \.6.È.8.17 IRAS 1873] L73-MM È.8 \.9 \..48 L73-MM Ñow....3g.3f 1.È1.h \.7h \.3h B335i È.5 \.5 \.5.8 B335 Ñowi L1157-MM È È L1157-MM blue lobe È3..9È3.4.3È.4 L1157-MM red lobe D.3j... CepE-MM blue lobe CepE-MM red lobe a L is deðned as the sum over the entire Ñow of L (O I) ] L (CO) ] L (H O) ] L (OH). Where a range of values is given, the averaged FIR value is used for the calculation, while upper limits are not considered. b Nisini et al.. c Revised values with respect to those given in Nisini et al.. d Benedettini et al.. e Ceccarelli et al f Average value along the Ñow. g [O I] 63 and 145 km lines detected only on the positions (1È5) and (1È1), respectively. h Detected only on the position (1È). i Nisini et al. 1999a. j Luminosity of the 179 km line. intensities are represented by triangles and squares, respectively, while the observations are the open circles. In case of blending with lines coming from CO or OH (namely, at 71.7, 89.99, 1.91, 1.98, and km; see Table 3), we Ðtted the di erence between the total observed Ñux and the CO (or OH) line predicted intensity. The agreement between model and observation appears quite satisfactory, conðrming the high temperature of the emitting gas, as it was already suggested by the high-excitation energy of the detected lines. For the three modeled sources, we report in columns (4)È(7) of Table 4 the CO and H O column densities, the size of the projected emitting area, and the H O abundance. In the case of IRAS 4, we have taken the velocity dispersion of km s~1, which is the FWHM of the H O 557 GHz line as measured by the Submillimeter Wave Astronomy Satellite (SW AS) (Neufeld et al. ), while in the other cases a value of 1 km s~1 is assumed. As a general result from this analysis, we underline that the emission arises everywhere from spots ÏÏ of warm and dense material, in which the H O abundance appears enhanced by factors of 13È14 with respect to the values measured by SW AS in quiescent nebulae (Snell et al. ). Meaningful Ðts of the OH emission were derived only in IRAS 4, IRAS 1873]113, and the blue lobe of CepE- MM, again adopting the physical parameters derived from the CO Ðt. Although a good consistency between models and observations exists, the OH column densities needed to reproduce the observed intensities appear quite high,

13 5 GIANNINI, NISINI, & LORENZETTI Vol. 555 exceeding in all cases the value of 116 cm~, which would imply X(OH) º 1~4. This result must be taken with some caution because high column density values could in reality simply reñect the pumping of the population of the levels due to the continuum emission, an e ect that is particularly important in the case of the OH molecule and is not accounted for by our simple model. The last step of our analysis consists in the calculation of the contribution of each species to the total FIR line cooling, which is directly derived by our Ðts by adding up the predicted line intensities coming from all the transitions considered in the model (i.e., 6 for CO, 37 for H O, and 44 for OH). The results are given in Table 5, where we also report the coolings derived for O I and C II obtained from the intensity of the 63 and [O I] 145 km lines and from the 158 km line, respectively. In the sources where only a weak CO emission was observed, we obtained the reported coolings by applying our model in the extreme physical conditions compatible with the CO observations and then assigning for the H O and the OH coolings a range of values or an upper limit depending on whether or not a real line detection exists. We emphasize that, although the physical parameter space is in these cases quite wide, the cooling does not vary by more than 3%, thus a meaningful estimate of this quantity is in any case guaranteed. From the inspection of the reported data it is noticeable that O I, CO, and H O contribute to the whole cooling roughly equally (with a dispersion of less than an order of magnitude and with a mean value of B1~ L ), while L (OH) is _ slightly lower; this means that a strongly prevailing cooling channel does not exist and thus the released energy is approximately shared out among all the species. For this reason, more than the cooling produced by a single species, the total molecular cooling [L \ L (CO) ] L (H O) mol ] L (OH)] seems to be an interesting quantity to be exploited. In particular, the ratio L /L ranges for our mol bol sources (again with the exception of L157) between 1.3 ] 1~3 and 3. ] 1~. This ratio is systematically higher than that relative to class I sources (B. Nisini et al. 1, in preparation); this strongly indicates that FIR line cooling via molecular emission is a peculiar characteristic of the class phase of the protostellar evolution. 6. DISCUSSION 6.1. PDR Di use Emission from External Sources As we have pointed out in 4, the C II line at 158 km is observed in any of the pointed positions. This line is expected to be only a minor coolant in dense shock environments; in dissociative shocks, in particular, [C II] 158 km contributes by no more than 1% of the [O I] 63km emission (e.g., Hollenbach & McKee 1989), while in our measurements it is always brighter or comparable to the 63 km line. On the other hand, C` is a good tracer of PDRs. An origin of the [C II] 158 km line in a local PDR produced by the class itself is very unlikely because of the low luminosity of these sources and the high density of their environments, which prevent any far-ultraviolet (FUV) photons from escaping to distances large enough to create C` in our o -source measurements. In support of this expectation, we have plotted in Figure 5 the measured C II luminosity as a function of the source bolometric luminosity; the mere fact that there is not any correlation between these two quantities reveals that [C II] 158 km is not excited by photons FIG. 5.ÈL (C II) as a function of L. Luminosities are in units of L. Objects located in high-mass star-forming bol regions, close to OB associations, are explicitly _ indicated. coming from the source itself. This plot is, however, instructive in showing that four of the sources have on average larger values of L (C II), namely, HH 4ÈMMS, HH 5ÈMMS, VLA 1, and CepE-MM. Noticeably, these are the sources located in high-mass star-forming regions, close to OB associations. Moreover, the source HH 11ÈMM, which has a C II luminosity comparable to that of HH 4ÈMMS and HH 5ÈMMS, is located at only 7@ from the B1 star HD 4398, belonging to the Perseus OB association (Humphrey 1978), while all the other sources in Perseus are at angular distances of several degrees far from this high-mass association. This occurrence suggests that the origin of C II emission can be found in the di use matter illuminated by the local interstellar Ðeld; this hypothesis is also supported by the fact that wherever we have measurements in genuine o -source positions (i.e., positions far from the outñowing gas as well), we have detected a [C II] Ñux comparable to the value on source. Although not strictly relevant for the discussion about the local excitation mechanisms, it is worthwhile to estimate the strength of the external FUV Ðeld producing strong C II emission in the above four sources. In this way we can also give an estimate of how much of the observed [O I] 63km emission comes from the di use PDR. In the case of HH 4ÈMMS and HH 5ÈMMS, Benedettini et al. () have estimated that the di use [O I] 63km should account for at most 4% of the total O I emission. For the other three sources, a measure of the FUV Ðeld can be made using the relationship I(C II) \ 1~6G ergs cm~ s~1 sr~1 (Hollenbach, Takahashi, & Tielens 1991), where G is the FUV Ðeld intensity expressed in units of 1.6 ] 1~3 (Habing 1968). G values of 5, 55, and 33 are found for CepE-MM, VLA 1, and HH 11ÈMM, respectively. A di erent approach can instead be adopted directly computing the FUV Ðeld produced by the nearest O/B stars at the distance of the considered source: in this way we Ðnd G values of 44, 6, and 3, respectively, i.e., consistent with the estimates given above. Using the PDR

14 No. 1, 1 FIR LINE COOLING IN CLASS SOURCES 53 line ratio predictions given by the Web Infrared Tool Shed5 (Kaufman et al. 1999), we estimate for the above G values a ratio O I/C II ranging from.3 to 1, assuming a density range between 13 and 14 cm~3. Such ratios mean that the considered PDR can account for most of the oxygen observed in HH 11ÈMM and about 5% of that observed in CepE-MM and VLA 1. These estimates will be used in the following discussion about the excitation mechanisms more strictly related with the class sources. 6.. T he L ocal Excitation Mechanisms We can now address the question of the physical environment where the strong observed molecular emission arises. Having excluded the origin in strong PDRs, both locally or externally generated, only two mechanisms can be identiðed as potentially able to excite the observed FIR lines, namely, gas heating by the accretion photons in the infalling envelopes (Ceccarelli, Hollenbach, & Tielens 1996) and shock excitation along the outñowing material (e.g., Hollenbach 1997). Excitation by shock waves was identiðed as the main origin of most of the molecular emission in Ðve out of the six class objects so far studied (Nisini et al. 1999a, 1999b, ; Benedettini et al. ). On the other hand, Ceccarelli et al. (1999) have suggested that the water emission in two of the sources analyzed here (namely, IRAS 1693 and NGC 1333ÈIRAS 4) could have been produced directly in the protostar infalling envelopes. Having enlarged the sample, we can now rely on a more meaningful statistical basis to examine this problem. Observationally, the shock excitation hypothesis is well supported by our detection of emission lines on the outñow lobes at intensity levels comparable to those measured on source. Only in three sources (namely, L1448-MM, IRAS 4, and IRAS 1693) is the emission on source signiðcantly larger than that measured along the outñows, and therefore it could be argued that here an additional contribution from the envelope could be present. In L1448-MM, however, the strong molecular emission observed on source has been associated with a compact molecular jet representing a recent event of episodic mass ejection, considering that the CO lines observed by ISO can be nicely Ðtted with groundbased CO 4È3 and CO 3È observations from this jet (Nisini et al. ). Further evidence in favor of shock excitation has been o ered by the recent SW AS measurements of the line proðle of the 1 ] 1 water transition near 557 GHz. Two of our sources, 1 namely, 1 IRAS 4 and L1157, have been observed by SW AS (Neufeld et al. ), and in both cases broad line proðles, with FWHM of about km s~1, clearly associated with the outñowing gas, were observed. An additional test to ascertain whether the observed emission originates in the outñow is to search for a correlation between the luminosity emitted in the FIR lines and the kinetic luminosity of the outñows. To this aim, we deðne L \ L (O I) ] L (CO) ] L (H O) ] L (OH) as the parameter FIR describing the total radiative line luminosity in the FIR, obtained by summing up the coolings measured along all the outñow pointings (see Table 5). To obtain reliable estimates of L, we do not consider upper limits, and in the cases in which FIR a range of luminosities is given, we adopt the average value. Moreover, in the computation of L we do not include the [C II] 158 km luminosity because, FIR as we 5 have already discussed, this line is mainly emitted in the di use medium rather than in the shocked gas. Now, in the hypothesis of fully radiative shock, L should be somehow related to the kinetic luminosity L FIR associated with the outñowing material. Using the literature kin determinations of L for our sources (Table 6, col. [3]), we searched for such a kin relation by plotting L versus L (Fig. 6). Unfortunately, L is only indirectly FIR measured kin by means of CO millimeter kin mapping, thus su ering from the large uncertainties related to the determination of the angle between the outñow axis and the plane of the sky and to opacity e ects. Despite that, we e ectively Ðnd a rather linear correlation with a unitary slope between L and L. This occurrence has two important implications: FIR on kin the one hand, it represents a further and independent proof that lines originate in the outñow; on the other hand, it also demonstrates that L is really a good representation of L. FIR kin This last statement, in particular, deserves a deeper investigation. Indeed, Davis & Eislo ffel (1996) pointed out that, if most of the kinetic energy is thermalized and radiated away at the shock working surface, and if the velocity of the swept-up gas, which coincides with the CO outñow, roughly equals the shock velocity (V B V ), then the total shock out shock luminosity L has to be equal to L. Given this fact, the rad kin correlation we found can be explained under the condition that the FIR luminosity is a good measurement of the total radiative power deposited along the Ñow. We, however, know that most of the considered sources have strong H emission coming from the shocks along their outñows, which is usually traced by the 1È S(1) transition at.1 km. To investigate how the shock radiative power is shared out between H and FIR emission, we report in column (6) of FIG. 6.ÈFIR line luminosities [L \ L (O I) ] L (CO) ] L (H O) ] L (OH)] plotted vs. the outñow kinetic FIR luminosity L. Luminosities are in units of L. The dotted lines refer to di erent kin L determinations available in the _ literature. The best linear Ðt to the kin data (dashed line) together with the relative parameters are indicated. Numbers associated with the observation points refer to the source name according to the identiðcation code given in Table 1.

15 54 GIANNINI, NISINI, & LORENZETTI Vol. 555 TABLE 6 LITERATURE PARAMETERS L M kin wind a LL A H V IdentiÐcation Number Source (L ) (M yr~1) References (1~ L ) (mag) References _ (1) () (3) (4) (5) (6) (7) (8) 1... L1448ÈIRS ] 1~6 1, 5 6È9 19,... L1448-MM.3È ] 1~6 1,, 3.3 5È16 19, 3... NGC 1333ÈIRAS NGC 1333ÈIRAS ] 1~7È1.4 ] 1~6 4, IRAS 38] ] 1~7È3 ] 1~6, HH 11ÈMM È L ~3 7.9 ] 1~8, L1641ÈVLA ] 1~7 9 6 D1.5 4, HH 4ÈMMS È19 6, HH 5ÈMMS [ IRAS 1693[4.18È.47 È7 ] 1~6, L483-MM ] 1~7, IRAS 1873] L73-MM.13È ] 1~6 13, 14, 15 [ B ] 1~7, 13 \ L1157-MM.3È ] 1~7 16, \ CepE-MM.43È.18 1È3.4 ] 1~ D3 18, 31 a Wind mass-loss rate derived by the outñow momentum by assuming that M and km s~1. wind v wind \ M out v v \ 1 REFERENCES.È(1) Bachiller et al () Bontemps et al (3) Barsony et al (4) Knee out & Sandell wind. (5) Blake et al (6) Cabrit & Bertout 199. (7) Gueth & Guilloteau (8) Tamura et al (9) Correia, Griffin, & Saraceno (1) Gibb & Davis (11) Mizuno et al (1) Parker, Padman, & Scott (13) Moriarty-Schieven & Snell (14) Avery, Hayashi, & White 199. (15) Goldsmith et al (16) Umemoto et al (17) Gueth, Guilloteau, & Bachiller (18) Ladd & Hodapp (19) Bally, Lada, & Lane 1993b. () Nisini et al.. (1) Garden, Russell, & Burton 199. () Bally et al. 1993a. (3) McCaughrean, Rayner, & Zinnecker (4) Davis & Eislo ffel (5) Chen et al (6) Davis et al (7) Benedettini et al.. (8) Fuller et al (9) Palacios & Eiroa (3) Nisini et al. 1999a. (31) LeÑoch,Eislo ffel, & Lazare Table 6 the literature determinations of L, which is in H general obtained from the Ñux of the 1È S(1) line multiplied by a factor of 1 to take into account the contribution of all the other H rovibrational lines. The so-obtained L values are, for the majority of the cases, only a fraction H of the derived L luminosity. The extinction can, however, FIR signiðcantly reduce the observed H luminosity in the con- sidered embedded sources. In column (7) of Table 6 the A V determinations for our sources retrieved from the literature are listed: they range from a few unities to about mag. An average value of 1 mag produces an increase of a factor of D.7 in the observed L, una ecting the general conclusion that L is a good measure H of the overall shock cooling FIR and that it represents an alternative powerful way to measure directly the energy deposited in the outñow, with the further advantage of being una ected by geometrical or extinction problems C-Shock Components Having established the origin of most of the observed emission in the outñow, we now focus our attention on the possible shock components that can contribute to the line excitation. We can rule out a substantial contribution to the molecular emission from dissociative shocks ( J ÏÏ-shocks; Hollenbach & McKee 1989) because these are able to predict molecular emission as a main coolant only if the preshock density exceeds values of about 15È16 cm~3, whereas we Ðnd similar or lower values for the density of the postshocked gas (see Table 4). This represents an evident inconsistency, only reconcilable by unrealistically assuming a unit compression factor (n /n ), instead of a more realistic value of 5È1. post pre The presence of strong molecular emission, from FIR lines of both CO and H O and from near-infrared H emis- sion in all the sources of our sample (with the exception of L157), suggests that a nondissociative, slow shock (Cshock; Draine 198) should be the prevailing excitation mechanism. An independent way of verifying this hypothesis consists in evaluating the velocity derivable from the shock radiative cooling, once the observed parameters are used. The total intensity in the shock cooling radiation can be expressed as (e.g., Hollenbach A 1997) I \ 9 ] 1~9 n v3 a s B s 14 cm~3 1 km s~1 ] ergs s~1 cm~ sr~1, (1) where n and v are the preshock density and the shock velocity, a respectively. s Considering a preshock density from 14 to 15 cm~3, which are the typical densities in the dense cores hosting these sources, and an emitting region of about 1~9 sr (an average value taken from Table 4), the observed FIR luminosities, which range from 1~ to a few times 1~1 L, imply shock velocities from D5 to4kms~1. These are _ the typical velocities giving rise to nondissociative shocks. The main prediction from C-shock models, which can now be tested for the Ðrst time on a larger statistical basis, is that a large amount of water should be produced as a consequence of the very efficient high-temperature (T Z 3 K) reactions that convert all oxygen not locked in CO into water. The Ðnal H O abundance should be in this case larger than D1~4. Table 4 shows that, when reliable abundance determinations can be obtained, values ranging from 1~5 to a few times 1~4 are found. In Figure 7 we have plotted the estimated abundances as a function of the gas temperature derived from our Ðts to the molecular lines. There is a trend for the water abundance to increase almost linearly with the temperature up to values of a few times