Ethynyl (C2H) in massive star formation: Tracing the initial conditions?

0.21 MB·English

by  H. Beuther

Accepted forAstrophysicalJournalLetters,draftversionfromFebruary4,2008
PreprinttypesetusingLATEXstyleemulateapjv.14/09/00
ETHYNYL (C H) IN MASSIVE STAR FORMATION: TRACING THE INITIAL CONDITIONS?
2
H. Beuther, D. Semenov, Th. Henning, H. Linz
MaxPlanckInstitute forAstronomy,Ko¨nigstuhl17,69117Heidelberg,Germany
Accepted for Astrophysical Journal Letters,draft versionfrom February 4, 2008
ABSTRACT
APEXsingle-dishobservationsatsub-millimeterwavelengthstowardasampleofmassivestar-forming
regionsrevealthatC Hisalmostomni-presenttowardallcoveredevolutionarystagesfromInfraredDark
2
8
Clouds via High-Mass Protostellar Objects to Ultracompact Hii regions. High-resolution data from the
0
SubmillimeterArraytowardonehot-corelikeHigh-MassProtostellarObjectshowashell-likedistribution
0
of C H with a radius of ∼9000AU around the central submm peak position. These observed features
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are well reproduced by a 1D cloud model with power-law density and temperature distributions and a

n gas-grain chemical network. The reactive C H radical (ethynyl) is abundant from the onset of massive
2
a
star formation, but later it is rapidly transformed to other molecules in the core center. In the outer
J
cloud regions the abundance of C2H remains high due to constant replenishment of elemental carbon
9
fromCObeingdissociatedbythe interstellarUVphotons. WesuggestthatC Hmaybeamoleculewell
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suited to study the initial conditions of massive star formation.

] Subject headings: stars: formation – astrochemistry – stars: individual (IRAS18089-1732)
h
p
- 1. introduction and interferometer observations with chemical modeling.
o
Although C H was previously observed in low-mass cores
r Spectralline surveyshaverevealedthathigh-massstar- 2
t and Photon Dominated Regions (e.g., Millar & Freeman
s forming regions are rich reservoirs of molecules from sim-
1984; Jansen et al. 1995), so far it was not systematically
a plediatomicspeciestocomplexandlargermolecules(e.g.,
[ investigatedintheframeworkofhigh-massstarformation.
Schilke et al. 1997; Hatchell et al. 1998; Comito et al.

1 2005; Bisschop et al. 2007). However, there have been 2. observations
v rarely studies undertaken to investigate the chemical
3 evolution during massive star formation from the ear- The 21 massive star-forming regions were observed
9 liest evolutionary stages, i.e., from High-Mass Starless with the Atacama Pathfinder Experiment (APEX) in the
4 Cores (HMSCs) and High-Mass Cores with embedded 875µm window in fall 2006. We observed 1GHz from
4 low- to intermediate-mass protostars destined to become 338 to 339GHz and 1GHz in the image sideband from
1. massive stars, via High-Mass Protostellar Objects (HM- 349 to 350GHz. The spectral resolution was 0.1kms−1,
0 POs) to the final stars that are able to produce Ul- but we smoothed the data to ∼0.9kms−1. The average
8 tracompact Hii regions (UCHiis, see Beuther et al. 2007 system temperatures were around 200K, each source had
0 for a recent description of the evolutionary sequence). on-source integration times between 5 and 16 min. The
v: The first two evolutionary stages are found within so- datawereconvertedtomain-beamtemperatureswithfor-
i called Infrared Dark Clouds (IRDCs). While for low- ward and beam efficiencies of 0.97 and 0.73, respectively
X mass stars the chemical evolution from early molecu- (Belloche et al.2006). Theaverage1σ rmswas0.4K.The
r lar freeze-out to more evolved protostellar cores is well mainspectralfeaturesofinterestaretheC Hlinesaround
2
a
studied (e.g., Bergin & Langer 1997; Dutrey et al. 1997; 349.4GHz with upper level excitation energies Eu/k of
Pavlyuchenkov et al. 2006; Jørgensenet al. 2007), it is 42K(line blends ofC2H(45,5−34,4) &C2H(45,4−34,3)at
far from clear whether similar evolutionary patterns are 349.338GHz, and C2H(44,4−33,3) & C2H(44,3−33,2) at
present during massive star formation. 349.399GHz). The beam size was ∼18′′.
To better understand the chemical evolution of high- The original Submillimeter Array (SMA) C H data
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massstar-formingregionsweinitiatedaprogramtoinves- toward the HMPO18089-1732 were first presented in
tigatethechemicalpropertiesfromIRDCstoUCHiisfrom Beuther et al. (2005b). There we used the compact and
an observational and theoretical perspective. We start extended configurations resulting in good images for all
withsingle-dishlinesurveystowardalargesampleobtain- spectrallinesexceptofC H.Forthisproject,were-worked
2
ing their basic characteristics, and then perform detailed on these data only using the compact configuration. Be-
studiesofselectedsourcesusinginterferometersonsmaller causethe C H emissionis distributedonlargerscales(see
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scales. Theseobservationsareaccompaniedbytheoretical §3), we were now able to derive a C H image. The inte-
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modeling of the chemical processes. Long-term goals are gration range was from 32 to 35kms−1, and the achieved
thechemicalcharacterizationoftheevolutionarysequence 1σ rms of the C H image was 450mJybeam−1. For more
2
in massive star formation, the development of chemical details on these observations see Beuther et al. (2005b).
clocks, and the identification of molecules as astrophysi-
cal tools to study the physical processes during different 3. results
evolutionary stages. Here, we present an initial study of
The sources were selected to cover all evolutionary
the reactive radical ethynyl (C H) combining single-dish
2 stages from IRDCs via HMPOs to UCHiis. We derived
1
2 Beuther et al.
our target list from the samples of Klein et al. (2005); 1D cloud model with a mass of 1200M⊙, an outer ra-
Fontani et al. (2005); Hill et al. (2005); Beltr´an et al. dius of 0.36pc and a power-law density profile (ρ ∝ rp
(2006). Table 1 lists the observed sources, their coordi- with p = −1.5) is the initially assumed configuration.
nates, distances, luminosities and a first order classifica- Three cases are studied: (1) a cold isothermal cloud with
tionintotheevolutionarysub-groupsIRDCs,HMPOsand T = 10K, (2) T = 50K, and (3) a warm model with a
UCHiis based on the previously available data. Although temperature profile T ∝ rq with q = −0.4 and a tem-
this classification is only based on a limited set of data, perature at the outer radius of 44K. The cloud is illumi-
here we are just interested in generalevolutionary trends. natedbythe interstellarUVradiationfield(IRSF, Draine
Hence,thedivisionintothethreemainclassesissufficient. 1978) and by cosmic ray particles (CRP). The ISRF at-
Figure 1 presents sample spectra toward one source of tenuation by single-sized 0.1µm silicate grains at a given
each evolutionary group. While we see several CH OH radius is calculated in a plane-parallel geometry follow-
3
lines as well as SO and H CS toward some of the HM- ing van Dishoeck (1988). The CRP ionization rate is as-
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POs and UCHiis but not toward the IRDCs, the surpris- sumed to be 1.3×10−17 s−1 (Spitzer & Tomasko 1968).
ing result of this comparison is the presence of the C H The gas-grain chemical model by Vasyunin et al. (2008)
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linesaround349.4GHztowardallsourcetypesfromyoung with the desorption energies and surface reactions from
IRDCs via the HMPOs to evolved UCHiis. Table 1 lists Garrod et al. (2007) is used. Gas-phase reaction rates are
the peak brightness temperatures, the integrated intensi- taken from RATE06 (Woodall et al. 2007), initial abun-
ties and the FWHM line-widths of the C H line blend at dances,wereadoptedfromthe“lowmetal”setofLee et al.
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349.399GHz. Theseparationofthetwolinesof1.375MHz (1998).
alreadycorrespondstoaline-widthof1.2kms−1. Wehave Figure 3 presents the C H abundances for the three
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three C H non-detections (2 IRDCs and 1 HMPO), how- models at two different time steps: (a) 100yr, and (b)
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ever, with no clear trend with respect to the distances or in a more evolved stage after 5×104yr. The C H abun-
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the luminosities(the lattercomparisonisonlypossiblefor dance is high toward the core center right from the be-
the HMPOs). While IRDCs are on average colder than ginning of the evolution, similar to previous models (e.g.,
more evolved sources, and have lower brightness temper- Millar & Nejad 1985; Herbst & Leung 1986; Turner et al.
atures, the non-detections are more probable due to the 1999). During the evolution, the C H abundance stays
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relatively low sensitivity of the short observations (§2). approximately constant at the outer core edges, whereas
Hence, the data indicate that the C H lines are detected it decreases by more than three orders of magnitude in
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independentoftheevolutionarystageofthesourcesincon- the center, except for the cold T = 10 K model. The
trasttothesituationwithothermolecules. Whencompar- C H abundance profiles for all three models show similar
2
ing the line-widths between the different sub-groups, one behavior.
finds only a marginal difference between the IRDCs and Thechemicalevolutionofethynylisdeterminedbyrela-
the HMPOs (the average ∆v of the two groups are 2.8 tiveremovalratesofcarbonandoxygenatomsorionsinto
and3.1kms−1). However,theUCHiisexhibitsignificantly molecules like CO, OH, H O. Light ionized hydrocarbons
2
broader line-widths with an average value of 5.5kms−1. CH+ (n=2..5) are quickly formed by radiative association
n
Intrigued by this finding, we wanted to understand the of C+ with H and hydrogen addition reactions: C+ →
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C H spatial structure during the different evolutionary CH+ → CH+ → CH+. The protonated methane reacts
2 2 3 5
stages. Therefore,wewentbacktoadatasetobtainedwith with electrons, CO, C, OH, and more complex species at
the Submillimeter Arraytowardthe hypercompactHiire- later stage and forms methane. The CH molecules un-
4
gion IRAS18089-1732 with a much higher spatial resolu- dergo reactive collisions with C+, producing C H+ and
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tionof∼1′′ (Beuther et al.2005b). Albeitthishypercom- C H+. An alternative way to produce C H+ is the dis-
2 3 2 2
pactHiiregionbelongstotheclassofHMPOs,itisalready sociative recombination of CH+ into CH followed by re-
5 3
inarelativelyevolvedstageandhasformedahotcorewith actions with C+. Finally, C H+ and C H+ dissociatively
a richmolecular spectrum. Beuther et al.(2005b) showed 2 2 2 3
recombine into CH, C H, and C H . The major removal
2 2 2
the spectraldetectionofthe C Hlinestowardthis source,
2 forC Hiseitherthedirectneutral-neutralreactionwithO
2
but they did not present any spatially resolved images.
thatformsCO,orthe samereactionbut withheaviercar-
To recover large-scale structure, we restricted the data to
bon chain ions that are formed from C H by subsequent
2
thosefromthecompactSMAconfiguration(§2). Withthis
insertion of carbon. At later times, depletion and gas-
refinement, we were able to produce a spatially resolved
phasereactionswith morecomplex species mayenter into
C Hmapofthelineblendat349.338GHzwithanangular
2 this cycle. At the cloudedge the interstellarUV radiation
resolution of 2.9′′×1.4′′ (correspondingto an averagelin-
instantaneously dissociates CO despite its self-shielding,
earresolutionof7700AUatthegivendistanceof3.6kpc).
re-enriching the gas with elemental carbon.
Figure2presentsthe integratedC Hemissionwithacon-
2 The transformation of C H into CO and other species
2
touroverlayofthe860µmcontinuumsourceoutlining the
proceeds efficiently in dense regions, in particular in the
positionofthemassiveprotostar. Incontrasttoalmostall
“warm” model where endothermic reactions result in rich
othermolecularlinesthatpeakalongwiththedustcontin-
molecular complexity of the gas (see Fig. 3). In contrast,
uum (Beuther et al. 2005b), the C H emission surrounds
2 inthe“cold”10Kmodelgas-graininteractionsandsurface
the continuum peak in a shell-like fashion.
reactions become important. As a result, a large fraction
of oxygen is locked in water ice that is hard to desorb
4. discussion and conclusions (Edes ∼ 5500 K), while half of the elemental carbon goes
to volatilemethane ice (E ∼1300K).UponCRP heat-
To understand the observations, we conducted a sim- des
ing of dust grains, this leads to much higher gas-phase
ple chemical modeling of massive star-forming regions. A
C H in massive star formation 3
2
abundance of C H in the cloud core for the cold model effects appear the more plausible explanation.
2
comparedtothewarmmodel. Theeffectisnotthatstrong The fact that we see C H at the earliest and the later
2
for less dense regions at larger radii from the center. evolutionarystagescanbeexplainedbythereactivenature
Since the C H emission is anti-correlated with the ofC H:itisproducedquicklyearlyonandgetsreplenished
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dust continuum emission in the case of IRAS18089-1732 atthecoreedgesbythe UVphotodissociationofCO.The
(Fig.2), we do not have the H column densities to quan- inner “chemical” hole observed toward IRAS18089-1732
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titatively compare the abundance profiles of IRAS18089- can be explained by C H being consumed in the chemi-
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1732 with our model. However, data and model al- cal network forming CO and more complex molecules like
low a qualitative comparison of the spatial structures. larger carbon-hydrogencomplexes and/or depletion.
Estimating an exact evolutionary time for IRAS18089- The data show that C H is not suited to investigate
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1732 is hardly possible, but based on the strong molec- the central gas cores in more evolved sources, however,
ular line emission, its high central gas temperatures and our analysis indicates that C H may be a suitable tracer
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the observed outflow-disk system (Beuther et al. 2004b,a, of the earliest stages of (massive) star formation, like
2005b), an approximate age of 5×104yr appears reason- N H+ orNH (e.g.,Bergin et al.2002;Tafalla et al.2004;
2 3
able. Although dynamical and chemical times are not Beuther et al. 2005a; Pillai et al. 2006). While a spatial
necessarily exactly the same, in high-mass star formation analysis of the line emission will give insights into the
they should not differ to much: Following the models by kinematicsofthegasandalsotheevolutionarystagefrom
McKee & Tan(2003)orKrumholz et al.(2007),thelumi- chemicalmodels,multipleC Hlineswillevenallowatem-
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nosityrisesstronglyrightfromthe onsetofcollapsewhich perature characterization. With its lowest J = 1 − 0
canbe consideredasastartingpointforthe chemicalevo- transitionsaround87GHz,C Hhaseasilyaccessiblespec-
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lution. At the same time disks andoutflows evolve,which tral lines in several bands between the 3mm and 850µm.
should hence have similar time-scales. The diameter of Furthermore, even the 349GHz lines presented here have
the shell-like C H structure in IRAS18089-1732 is ∼ 5′′ stillrelativelylowupper levelexcitationenergies(E /k ∼
2 u
(Fig.2), or ∼9000AU in radius at the given distance of 42K), hence allowing to study cold cores even at sub-
3.6kpc. This value is well matched by the modeled re- millimeter wavelengths. This prediction can further be
gion with decreased C H abundance (Fig.3). Although proved via high spectral and spatial resolution observa-
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in principle optical depths and/or excitation effects could tions of different C H lines toward young IRDCs.
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mimic the C H morphology, we consider this as unlikely
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because the other observedmolecules with many different
transitions all peak toward the central submm continuum H.B. acknowledges financial support by the Emmy-
emissioninIRAS18089-1732(Beuther et al.2005b). Since Noether-Programm of the Deutsche Forschungsgemein-
C H is the only exception in that rich dataset, chemical schaft (DFG, grant BE2578).
2
REFERENCES
Belloche,A.,Parise,B.,vanderTak,F.F.S.,etal.2006,A&A,454, Jørgensen,J.K.,Bourke,T.L.,Myers,P.C.,etal.2007,ApJ,659,
L51 479
Beltra´n,M.T.,Brand,J.,Cesaroni,R.,etal.2006,A&A,447,221 Klein,R.,Posselt,B.,Schreyer,K.,Forbrich,J.,&Henning,T.2005,
Bergin, E. A., Alves, J., Huard, T., & Lada, C. J. 2002, ApJ, 570, ApJS,161,361
L101 Krumholz,M.R.,Klein,R.I.,&McKee,C.F.2007,ApJ,656,959
Bergin,E.A.&Langer,W.D.1997, ApJ,486,316 Lee,H.-H.,Roueff,E.,PineaudesForets,G.,etal.1998,A&A,334,
Beuther,H.,Churchwell,E.B.,McKee,C.F.,&Tan,J.C.2007,in 1047
Protostars andPlanets V,ed.B. Reipurth,D.Jewitt, & K.Keil, McKee,C.F.&Tan,J.C.2003,ApJ,585,850
165–180 Millar,T.J.&Freeman,A.1984,MNRAS,207,405
Beuther, H.,Hunter,T.R.,Zhang,Q.,etal.2004a,ApJ,616, L23 Millar,T.J.&Nejad,L.A.M.1985, MNRAS,217,507
Beuther,H.,Zhang,Q.,Greenhill,L.J.,etal.2005a, ApJ,632,355 Pavlyuchenkov, Y.,Wiebe, D.,Launhardt, R.,&Henning,T.2006,
Beuther, H.,Zhang,Q.,Hunter,T.R.,etal.2004b,ApJ,616, L19 ApJ,645,1212
Beuther, H., Zhang, Q., Sridharan, T. K., & Chen, Y. 2005b, ApJ, Pillai,T.,Wyrowski,F.,Carey,S.J.,&Menten,K.M.2006,A&A,
628,800 450,569
Bisschop,S.E.,Jørgensen,J.K.,vanDishoeck,E.F.,&deWachter, Schilke, P., Groesbeck, T. D.,Blake, G. A.,& Phillips,T.G. 1997,
E.B.M.2007,A&A,465,913 ApJS,108,301
Comito,C.,Schilke,P.,Phillips,T.G.,etal.2005,ApJS,156,127 Spitzer,L.J.&Tomasko,M.G.1968,ApJ,152,971
Draine,B.T.1978,ApJS,36,595 Tafalla,M.,Myers,P.C.,Caselli,P.,&Walmsley,C.M.2004,A&A,
Dutrey,A.,Guilloteau,S.,&Guelin,M.1997, A&A,317,L55 416,191
Fau´ndez, S.,Bronfman,L.,Garay,G.,etal.2004, A&A,426,97 Turner,B.E.,Terzieva,R.,&Herbst,E.1999,ApJ,518,699
Fontani, F.,Beltra´n,M.T.,Brand,J.,etal.2005,A&A,432,921 van Dishoeck, E. F. 1988, in Rate Coefficients in Astrochemistry.
Garrod,R.T.,Wakelam, V.,&Herbst,E.2007,A&A,467,1103 ProceedingsofConferenceinUMIST,UnitedKingdom,1987.Eds,
Hatchell, J., Thompson, M. A., Millar, T. J., & MacDonald, G. H. T.J.Millar,D.A.Williams,ed.T.J.Millar&D.A.Williams,49
1998,A&AS,133,29 Vasyunin,A.I.,Semenov,D.,Henning,T.,etal.2008,ApJ,inpress
Herbst,E.&Leung, C.M.1986, MNRAS,222,689 Vig,S.,Ghosh,S.K.,Ojha,D.K.,&Verma,R.P.2007,A&A,463,
Hill,T.,Burton,M.G.,Minier,V.,etal.2005, MNRAS,363,405 175
Jansen,D.J.,vanDishoeck,E.F.,Black,J.H.,Spaans,M.,&Sosin, Woodall,J.,Agu´ndez,M.,Markwick-Kemper,A.J.,&Millar,T.J.
C.1995,A&A,302,223 2007, A&A,466,1197
4 Beuther et al.
Table 1
Source parameters
Name Type R.A. Dec. d L Tmb R Tmbdv ∆v(C2H) Ref.
J2000 J2000 kpc log(L⊙) K Kkms−1 kms−1
IRAS07029 IRDC 07:05:11.1 -12:19:02 1.0 b 0.21 0.52 2.4±0.3 (1)
IRAS08477 IRDC 08:49:32.9 -44:10:47 1.8 b 0.23 0.76 3.1±0.5 (2,3)
IRAS09014 IRDC 09:03:09.8 -47:48:28 1.3 b – – – (2,3)
IRAS13039 IRDC 13:07:07.0 -61:24:47 2.4 b – – – (2,3)
IRAS14000 IRDC 14:03:36.6 -61:18:28 5.6 b 0.38 1.13 2.8±0.5 (2,3)
IRAS08211 HMPO 08:22:52.3 -42:07:57 1.7 3.5 0.38 0.85 2.1±0.2 (2,3)
IRAS08470 HMPO 08:48:47.9 -42:54:22 2.2 4.2 1.85 6.81 3.5±0.4 (3)
IRAS08563 HMPO 08:58:12.5 -42:37:34 1.7 3.2 1.08 3.30 2.9±0.1 (2,3)
IRAS09131 HMPO 09:14:55.5 -47:36:13 1.7 3.4 0.23 0.50 2.1±0.3 (2,3)
IRAS09209 HMPO 09:22:34.6 -51:56:26 6.4 4.1 – – – (2,3)
IRAS09578 HMPO 09:59:31.0 -57:03:45 1.7 3.9 0.21 0.64 2.8±0.4 (3)
IRAS10123 HMPO 10:14:08.8 -57:42:12 0.9/3.0c 3.4/4.4c 0.17a 0.46 2.5±0.4 (2,3)
IRAS10184 HMPO 10:20:14.7 -58:03:38 5.4 5.5 0.41 1.68 3.9±0.3 (3)
IRAS10276 HMPO 10:29:30.1 -57:26:40 5.9 4.9 0.24a 0.67 2.6±0.5 (3)
IRAS10295 HMPO 10:31:28.3 -58:02:07 5.0 5.8 0.69 3.45 4.7±0.3 (3)
IRAS10320 HMPO 10:33:56.4 -59:43:53 9.1 5.4 0.85 3.48 3.8±0.6 (3)
G294.97 UCHii 11:39:09.0 -63:28:38 1.3/5.8c 3.9/5.3c 0.31 0.78 2.3±0.4 (4)
G305.20 UCHii 13:11:12.3 -62:44:57 3.0/6.8c 5.1/6.1c 0.44 5.21 11.1±0.5 (4)
G305.37 UCHii 13:12:36.3 -62:33:39 3.0/6.8c d 1.24 6.28 4.8±0.2 (4)
G305.561 UCHii 13:14:25.8 -62:44:32 4.0 5.1 0.94 4.71 4.7±0.6 (4,5)
IRAS14416 UCHii 14:45:22.0 -59:49:39 2.8 5.1 1.24 6.19 4.7±0.6 (6)
Ref.: (1)Klein et al.(2005),(2)Fontaniet al.(2005),(3)Beltr´an et al. (2006),(4)Hill et al.(2005),(5)Fau´ndez et al.(2004),(6)
Vig et al. (2007).
aForthesesourceswelisttheparametersoftheC2Hlineblendat349.338GHz,forallothersourcesitistheC2Hlineat349.399GHz.
b Since theIRDCs are perdefault not detected at short wavelengths, they are no IRASsource and we cannot derivea luminosity.
c Near and far distances and corresponding luminosities.
d No IRAScounterpart, henceno luminosity estimate.
C H in massive star formation 5
2
Fig. 1.—SamplespectraobtainedwithAPEXinadouble-sidebandmode. Thespectracover1GHzofdataaround349.5GHzand1GHz
around338.5GHz. Allspectrallinesarelabeled,linesinthelowersidebandaremarkedwitha“(L)”.Thetop-panelshowsanIRDCexample,
themiddle-panelatypical HMPO/hotcoreandthebottom panel aUCHiiregion. TheC2Hline-widthsareindicatedineachpanel.
Fig. 2.—Thegrey-scaleshows theintegrated emission(from32to35kms−1)ofthe lineblendof C2H(45,5−34,4)andC2H(45,4−34,3)
around349.338GHzobtainedwiththeSubmillimeterArrayusingonlythecompactconfigurationdata(Beuther etal.2005b). Theresulting
synthesized beam isshownatthebottom leftcorner (2.9′′×1.4′′). TheC2Hemissionispresented ingrey-scalewiththincontours (dashed
contours negative features), and the 860µm continuum peak is shown in thick contours. The C2H emission starts at the 2σ level and
continuesin1σstepswiththe1σlevelof450mJybeam−1. The860µmemissioniscontouredfrom10to90%(step10%)ofthepeakemission
of1.4Jybeam−1.
6 Beuther et al.
-6.0
102 years C H
2
-8.0
)
H
/
X
-10.0
(
0
1
g
o
L -12.0 q=-0.4, T=44K
q=0, T=10K
q=0, T=50K
-14.0
3.0 3.5 4.0 4.5 5.0
Log (r, AU)
10
-6.0
5 104 years C H
2
-8.0
)
H
/
X
-10.0
(
0
1
g
o
L -12.0 q=-0.4, T=44K
q=0, T=10K
q=0, T=50K
-14.0
3.0 3.5 4.0 4.5 5.0
Log (r, AU)
10
Fig. 3.—Topleft: RadialprofilesoftheC2Habundanceforthethreecloudmodelsatearlytimes,t.100yr. Bottomleft: Thesamebut
forthelatertimestepof5×104yr. Theparameter q denotes theexponent oftheassumedtemperaturedistribution(seealsomaintext).
This figure "f2.jpg" is available in "jpg"(cid:10) format from:
http://arXiv.org/ps/0801.4493v1