NASA Technical Reports Server (NTRS) 20140011041: The Chemistry of Extragalactic Carbon Stars

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by  NASA Technical Reports Server (NTRS)

Mon.Not.R.Astron.Soc.426,2689–2702(2012) doi:10.1111/j.1365-2966.2012.21771.x
The chemistry of extragalactic carbon stars
(cid:2)
Paul M. Woods,1,2 C. Walsh,3 M. A. Cordiner4 and F. Kemper5
1DepartmentofPhysics&Astronomy,UniversityCollegeLondon,GowerStreet,LondonWC1E6BT
2JodrellBankCentreforAstrophysics,AlanTuringBuilding,SchoolofPhysicsandAstronomy,TheUniversityofManchester,OxfordRoad,
ManchesterM139PL
3AstrophysicsResearchCentre,SchoolofMathematics&Physics,Queen’sUniversityBelfast,UniversityRoad,BelfastBT71NN
4AstrochemistryLaboratoryandTheGoddardCenterforAstrobiology,NASAGoddardSpaceFlightCenter,Code691,8800GreenbeltRoad,
GreenbeltMD20771,USA
5AcademiaSinicaInstituteofAstronomy&Astrophysics,POBox23-141,Taipei10617,Taiwan
Accepted2012July23.Received2012July20;inoriginalform2012February17
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ABSTRACT loa
d
e
PromptedbytheongoinginterestinSpitzer InfraredSpectrometerspectraofcarbonstarsin d
theLargeMagellanicCloud,wehaveinvestigatedthecircumstellarchemistryofcarbonstars fro
m
inlow-metallicityenvironments.Consistentwithobservations,ourmodelsshowthatacetylene h
ttp
is particularly abundant in the inner regions of low metallicity carbon-rich asymptotic giant ://m
branchstars–moreabundantthancarbonmonoxide.Asaconsequence,largerhydrocarbons n
ra
have higher abundances at the metallicities of the Magellanic Clouds than in stars with s.o
x
solarmetallicity.Wealsofindthattheoxygenandnitrogenchemistryissuppressedatlower fo
rd
metallicity,asexpected.Finally,wecalculatemolecularlineemissionfromcarbonstarsinthe jo
u
LargeandSmallMagellanicCloudandfindthatseveralmoleculesshouldbereadilydetectable rn
a
ls
withtheAtacamaLargeMillimeterArrayatFullScienceoperations. .o
rg
Key words: Astrochemistry – stars: AGB and post-AGB – stars: carbon – circumstellar at N/
matter–infrared:stars–submillimetre:stars. A
S
A
G
o
d
(Schaefer2008),hasanaveragemetallicityofaround50percentso- da
1 INTRODUCTION rd
lar(Dufour,Shields&Talbot1982;Westerlund1997).TheSmall S
p
TheadventoftheAtacamaLargeMillimeterArray(ALMA)and Magellanic Cloud (SMC), at a slightly larger distance of 66kpc a
c
e
other large (sub-)millimetre telescopes, with their unprecedented (Szewczyk etal.2009),hasanaverage metallicityof20per cent F
spatialresolutionandsensitivity,willallowtheobservationofgiant solar(Dufouretal.1982;Westerlund1997).Theeffectofthislow lig
h
starsinothergalaxiesinsimilardetailtothatachievedforGalactic metallicity regime on dust composition and galactic dust budgets t C
objects.Theseadvancedcapabilitiespromptinvestigationintothe hasbeenstudiedobservationallybymanyauthors(e.g.Zijlstraetal. tr o
n
natureoftheseextragalacticstars,withoneofthemostinteresting 2006; Lagadec et al. 2007; van Loon et al. 2008; Matsuura et al. M
aspects being the study of the effect of sub-solar metallicities on 2009). However, the effects on circumstellar chemistry have not ay
8
circumstellarchemistryanddustcomposition. been,asyet,studiedinanydetail,andouraimhereistopioneerin , 2
RecentstudiesintheinfraredusingtheSpitzerSpaceTelescope thefieldwiththiswork. 01
3
andground-basedinstrumentshavehighlightedthedeepmolecular The chemical modelling of Galactic asymptotic giant branch
absorptionof,primarily,acetyleneinthespectraofevolvedcarbon (AGB)starswas firstattempted 35 years ago. Initialattempts fo-
stars, in the Magellanic Clouds (MCs; e.g. van Loon, Zijlstra & cused on simple physical models and chemistry appropriate for
Groenewegen1999a;Matsuuraetal.2002,2005;vanLoon2006; oxygen-rich(n(O)>n(C))circumstellarenvironments(Goldreich
Sloanetal.2006;Specketal.2006;Zijlstraetal.2006;Lagadec & Scoville 1976; Scalo & Slavsky 1980; Jura & Morris 1981).
etal.2007;Leisenring,Kemper&Sloan2008;vanLoonetal.2008; Physicalmodelshavelargelyremainedsimple(withsomenotable
Woodsetal.2011,etc.).Theseabsorptionfeaturesareingeneral exceptions, e.g. Cordiner & Millar 2009) whereas the chemical
deeperthanthoseseeninGalacticstars,andthisimpliesthatthereis modelling has moved to focus on carbon-rich (n(C) > n(O)) cir-
adifferencebetweenthechemistryofMagellaniccircumstellaren- cumstellar chemistry since it shows a wider variety of molecules
velopes(CSEs)andGalacticcarbonstars,whicharecomparatively (e.g.Huggins&Glassgold1982).Progressinchemicalmodellingis
wellstudied. driveninpartbythedesiretoexplainobservedabundancesofnewly
The MCs are nearby dwarf galaxies with sub-solar metallici- detectedmoleculesinthemostaccessiblecarbonstar,IRC+10216,
ties.TheLargeMagellanicCloud(LMC),atadistanceof∼50kpc forinstance,intheadditionoflongcarbon-chainmolecules(Millar,
Herbst&Bettens2000)oranionspecies(Millaretal.2007).Inthis
case,advancesintechnologyhavedrivenustoinvestigatecarbon-
(cid:2)E-mail:[email protected] richcircumstellarchemistryinpreviouslychallenginglocations.
(cid:3)C 2012TheAuthors
MonthlyNoticesoftheRoyalAstronomicalSociety(cid:3)C 2012RAS
2690 P. M. Woods et al.
Inthispaper,wepresenttheresultsofmodellingthecircumstellar wheredustisformedandthenaccelerated(5<(cid:5)R<(cid:5)100R )dueto
(cid:2)
chemistryaroundcarbon-richAGBstarsatsub-solarmetallicities. non-equilibriumprocesses.TheycompriseCS,SiO,SiS,NH ,SiH
3 4
Wefocusonthreefiducialmodels,atmetallicitiesandwithphysical andCH (e.g.Cherchneff2006).Thesetwosubsetsofmolecules
4
conditionsappropriatefortheGalaxy,theLMCandtheSMC.We areoftencalled‘parentspecies’.
initially assume solar metallicity (Z = 0.02) for Galactic carbon Toobtainaccurateinitialabundancesforthese10parentspecies
stars, and average interstellar metallicities for LMC (Z = 0.008) whenmodellingthechemistryintheouterenvelope(R >(cid:5)100R )
(cid:2)
and SMC (Z = 0.004) carbon stars. In Section 2, we describe atdifferentmetallicities,onemustemploydifferenttechniquesfor
ouradoptedphysicalmodelanddiscussthedifferentphysicaland thetwosubsetsofparents.Forthehighstabilityparentspecies,we
chemicalconsiderations,andalsodiscusssourcesofuncertainties can use a TE model to calculate the relevant data. This approach
in our models in Section 2.6. In Section 3, we describe how we hasbeenusedpreviouslytogoodeffect(e.g.Tsuji1973;Sharp&
calculatethechemicalevolutionoftheCSE.Weshowourresults Huebner1990;Markwick2000).Fortheremainingparentspecies,
in Section 4, preceding a discussion in Section 5 (including our other physical factors must be taken into consideration, such as
calculationsofmolecularlineemissionwhichmaybeobservable depletion through dust formation (SiS, SiO; Bieging & Nguyen-
with ALMA) and finally, in Section 6, we draw our conclusions Quang-Rieu 1989; Scho¨ier et al. 2006), pulsation-driven shocks
regardingthechemistryoflow-metallicitycarbonstars. (SiS,SiO,CS;Willacy&Cherchneff1998)andgas–graininterac-
tion(e.g.hydrogenationleadingtoNH ,SiH andCH ).Modelling
3 4 4 D
thesecomplexitiesindetail,whichindeedthemselveshavebeenthe o
2 CHEMICAL AND PHYSICAL w
CONSIDERATIONS AT LOW METALLICITY subjectofmuchinvestigation,istooadvancedforthisinitialstudy nlo
andthuswemustuseotherargumentsdiscussedbelow. ad
e
Ninutecrlieoorsayrnethmeitxicedprtoodtuhcetssudrrfeadcegeodf(thveiasctaorn,vwehcetiroenm)fartoemriatlhiessatceclelal-r d from
eratedtotheterminalvelocityofthestellarwindwithinaradiusof 2.1.1 TEmolecules(CO,N2,C2H2,HCN) h
ttp
20R(cid:2)(Keady,Hall&Ridgway1988)andpassedintotheCSE.The TheTEmodelusedissimilartothatdetailedbySharp&Huebner ://m
gas,whichismainlymolecularhydrogen,iswellmixedwithdust (1990)andMarkwick(2000),1whichworkbyminimizingtheGibbs n
grains.Weassumeasphericalgeometrywherethegashasa1/R2 freeenergyofthethermodynamicsystem.Weadoptatemperature ras.o
densitydistribution,andatemperatureprofilewhichfollows: of2250Kandapressureof1.033×10−3atm,appropriateforthe xfo
T(r)=max[150(R/R0)−0.79; 10]K (R≥R0 =5×1015cm) photosphere of a carbon-rich AGB star (Ivezic´ & Elitzur 1996; rdjo
Markwick 2000). Results of the TE calculations for the range of u
(e.g. Millar & Herbst 1994; Millar et al. 2000). The CSE is irra- temperatures2500–2000KaredisplayedinFig.1;thosespeciesfor rnals
diatedbytheinterstellarradiationfield(ISRF),butnotbyultravi- whichTEisappropriatearelargelyinvariantacrossthetemperature .o
olet(UV)photonsfromthestaritself,whicharequenchedinthe range.ElementalabundancesforAGBstarsatdifferentmetallicities arg/
stellar atmosphere. Extinction in the CSE is calculated according aretakenfromthenucleosynthesiscalculationsofKarakas(2010) t N
to the approach of Jura & Morris (1981), assuming interstellar- fora3M(cid:6) star.Thismasswaschosenbecauseitwasthelowest AS
A
typegrains,andwetreatCOself-shieldingaccordingtoMamon, mass model for which the star became carbon rich at the three G
Glassgold&Huggins(1988).Wemodelthecarbon-richchemistry metallicitesZ =0.02(Galaxy),Z =0.008(LMC)andZ =0.004 od
3in×th1e0c1i8rccumm).stellarregionbetween∼200and100000R(cid:2) (0.005– (cSleMarCf)r.oTmheorbesearrveatiinocnoanlgervuiidtieenscewitthhatthceasrebomnosdtaerlss,ininththeaMt iitlkiys dard S
p
Inthesubsequentsections,weconsiderthechemicalandfurther Way(MW)canformatmassesaslowas∼1.5M(cid:6)(Wallerstein& ac
e
pkhnyowsicnalabinogurtetdhieencthseomfisoturyr mofodexelt.raOgablsaecrtvicaticoanrablolyn,svtearrys,laitntldesios Kadndarpepss1e9d9(8K).arTahkeasse2m01o1d)e,llainndgwisesuceosnatirneuaecktonouwselethdegerdesaunltdsbseinincge Fligh
wtheenmuacilneloysydnisthcuessisstmheodcehlesmoifstKryaroafkaGsa(l2a0c1ti0c)ctaorbaodnjussttafrosr,luoswinegr tshteelylaarryeiealdcsownseirseteunstesdeatsoifnrpeuatdsiflyoratvhaeilTaEblmeoabduenladnadnctehse.rAesvuelrtaingge t Ctr on
metallicities (see Section 2.1). Physical parameters of LMC and fractionalabundanceswithrespecttoH canbefoundinTable1 M
2 a
SMCcarbonstarsaremorewellconstrainedobservationally,and for CO, N2, C2H2 and HCN. This method (using the results of y 8
wesummarizethoseaspectsinSections2.2–2.5. nucleosynthesis calculations) is preferable to using the elemental , 2
0
abundancesoftheMCsingeneral,sincemuchofthecircumstellar 13
chemistrydependsontheamountofcarbongeneratedbythestar
2.1 Initialchemicalabundances
intheAGBphase(Matsuuraetal.2008).Inthenextsection,we
Ascircumstellarmaterialcoolsduringtheexpansionofthestellar describe our method of determining input abundances for those
wind, it is energetically favourable for atomic elements to form molecules for which TE does not apply. Table 3 summarises the
molecules in the photosphere of the star (R <(cid:5) 5R(cid:2)), where the physicalparametersdiscussedinSection2.2–2.5.
chemistry is in thermal equilibrium (TE) due to the high tem-
peratures and densities. For molecules of high stability, that is to
say, those which contain strong internal bonds, the abundances 2.1.2 Non-TEmolecules(CS,SiO,SiS,NH3,SiH4,CH4)
amounted in this region of the star are carried through into the
TheobserveddistributionofSiOpeakswithinafewstellarradiiin
non-TE outer envelope (R >(cid:5) 100R(cid:2)) without significant change. thearchetypalcarbonstar,IRC+10216(Keady&Ridgway1993;
Incarbon-richenvironments,suchmoleculesincludeCO,N ,HCN
2 Scho¨ier et al. 2006). SiO is seen at high fractional abundances,
andC2H2(e.g.Tsuji1973;Millar2008). n(SiO)/n(H )∼10−6,whereasTEmodelsunderpredicttheabun-
2
A further group of molecules are also observed at high abun-
dancebyaroundafactorof30.Itsformationisthoughttobedueto
dancesattheinneredgeoftheouterenvelopeinGalacticstars.This
subsetisgenerallyformedatlowabundanceintheTEphotosphere
of the star, but attains a higher abundance in the inner envelope 1Availableonrequesttotheauthors.
(cid:3)C 2012TheAuthors,MNRAS426,2689–2702
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The chemistry of extragalactic carbon stars 2691
(discussedbelow).ForIRC+10216,thispremiseagreeswellwith
observations(Bieging&Tafalla1993;Lucasetal.1995;Scho¨ier
etal.2006;Decinetal.2010).
CS is slightly overproduced in TE models compared to obser-
vations (Keady & Ridgway 1993) and the work by Willacy &
Cherchneff (1998) has shown that CS is destroyed by pulsation-
driven shocks by a considerable amount, the extent of which de-
pends on the assumed shock speed. It then slowly reforms in the
outflowingwind.Weassumeslowshockspeeds(10kms−1)andthe
sameshock-destructionefficiencyasthatofWillacy&Cherchneff
(1998),andadoptanabundanceofCSafactorof2lessthanitsTE
abundance.
Emission from the hydrogenated species NH , CH and SiH
3 4 4
originates in the dust-formation zone and inner envelope around
IRC+10216 (Keady & Ridgway 1993). Although surprising be-
cause of the high temperature of the dust in these regions, these
D
moleculesaremostlikelyformedviahydrogen-additionreactions o
w
onthesurfaceofdustgrains.Theveryshorttime-scalebeforeejec- nlo
tionfromthegrainsurfacewouldindicatethathydrogenisthedom- ad
e
ianreanotfrleoawctiaobnunpdaratnnceerfinorbaodthsoTrbEe(dea.gto.mTssu.jTiy1p9i6c4a)llayntdhensoenm-ToElemcuoldes- d from
els(Willacy&Cherchneff1998,andothers),addingfurtherweight h
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tiongIRraCin+-s1u0r2fa1c6ehhayvderobgeeennatuiosendthteoodrieersiv.OebfrsaecrtviaotnioalnsaboufnwdaarnmceNsHo3f ://mn
0.2–2.0 × 10−6 (Keady & Ridgway 1993; Monnier et al. 2000; ras
.o
Hasegawa et al. 2006). We assume that the abundance of NH at x
3 fo
othermetalliciteswillscalewiththedustsurfacearea.FortheLMC rd
and SMC, this quantity and the associated gas-to-dust ratio and jou
grain-sizedistributionarehighly uncertain, although workiscur- rna
ls
rently ongoing to determine more definite values. Assuming that .o
the grain-size distribution in LMC and SMC carbon stars is sim- arg/
ilar to that in Galactic carbon stars, and that the gas-to-dust ratio t N
scaleswithmetallicity,wewilladopttheupperobservedfractional AS
A
abundanceofNH3forGalacticcarbonstars,andscaleaccordingto G
metallicityfortheLMCandSMCcarbonstars.Weapplyasimilar o
d
d
reasoningforscalingtheinitialfractionalabundancesofCH and a
4 rd
SiH4withmetallicity. Sp
Adopting the initial abundances detailed in Table 1 means that ac
e
onlysmallamountsofnitrogen,oxygenandsulphurareavailable F
for incorporation into other molecules and dust. However, in the ligh
Galacticcase,some5percentoftheelementalcarbonremainsfor t C
incorporation into molecules or carbonaceous dust. At LMC and tr o
n
SMC metallicities, the percentage of free carbon rises to 28 per M
a
centand32percent,respectively,whichisinreasonableagreement y
8
Figure1. Fractionalabundances(logscale)atTEforMW,LMCandSMC with Ferrarotti & Gail (2006) for an evolved star. These authors , 20
metallicities,atapressureofP∼10−3atm.Weutilizevaluesat2250Kin alsofoundthatatmost9percentofelementalsiliconcondensed 13
thesubsequentmodelling. intodust.Ourinitialabundancesmeanthat52–60percentofthe
elementalsiliconisavailablefordustordustseedformation.
circumstellarshocks,andhenceitsabundanceisdependentonthat InTable2,wehavecompiledacomparisonofC/Oratiosandthe
ofatomicoxygen,sinceintheshockedmaterialthereaction carbon excesses (log((cid:3) − (cid:3) )) for the three metallicity regimes.
C O
C/Oratiosincreasedrasticallyinthenucleosynthesiscalculations
Si +OH −→ SiO +H (1)
ofKarakas(2010),largelyduetothedecreaseinelementaloxygen.
dominates(Hartquist,Dalgarno&Oppenheimer1980;Willacy& Thecarbonexcessislargeratlowmetallicity,aneffectwhichhas
Cherchneff1998).Assuch,wechoosetosetthefractionalabun- beenobserved(e.g.Wahlinetal.2006).
danceofSiOto10−3n(O),inlinewiththatobservedinIRC+10216.
SiSisthemainrepositoryforbothsiliconandsulphurinthegas
2.2 Mass-lossrate
phase,andformsatahighabundanceinTEmodels.Observations
indicatethatitisrapidlydepletedasgasflowsthroughtheinterme- Mass-lossratesforAGBstarsaregenerallymeasuredusingobser-
diateenvelope,presumablyduetotheadsorptionofthemoleculeon vationsoftheCOenvelopeandthenassumingaCO-to-H scaling
2
tograinsurfaces(Bieging&Nguyen-Quang-Rieu1989;Boyleetal. ratio(theso-calledX-factor,X ).Alternatively,thenear-andmid-
CO
1994).Thus,weassumethattheinitialabundanceofSiSisequal infraredspectralenergydistributioncanbefittedwitharadiative-
to the elemental abundance of sulphur minus the fraction of CS transfermodeltoobtainadustmass-lossrate,andthenbyassuming
(cid:3)C 2012TheAuthors,MNRAS426,2689–2702
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2692 P. M. Woods et al.
Table1. Initialmodelfractionalabundances(w.r.t.H2).
Species ObservedGalacticinner TEabundance Adoptedinitial Adoptedinitial Adoptedinitial
envelopeabundance (Galactic) abundance(Galactic) abundance(LMC) abundance(SMC)
CO a6×10−4 4.4×10−4 4.4×10−4 1.7×10−4 8.9×10−5
N2 b– 5.4×10−5 5.4×10−5 1.2×10−5 3.9×10−6
C2H2 c8×10−5 3.3×10−5 3.3×10−5 2.8×10−4 3.1×10−4
HCN d2×10−5 2.4×10−5 2.4×10−5 3.3×10−5 2.0×10−5
SiS e3×10−6 3.7×10−6 6.9×10−6 2.1×10−6 1.0×10−6
CS c4×10−6 4.7×10−6 2.4×10−6 1.6×10−6 8.0×10−7
NH3 f2×10−6 3.3×10−11 2.0×10−6 8.0×10−7 4.0×10−7
CH4 c2×10−6 6.1×10−9 2.0×10−6 8.0×10−7 4.0×10−7
SiO g1×10−6 2.0×10−8 1.7×10−6 6.8×10−7 3.6×10−7
SiH4 c2×10−7 8.8×10−15 2.0×10−7 8.0×10−8 4.0×10−8
aKwan&Linke(1982). bNotobservedduetolackofpermanentdipole.2×10−4isusuallyassumedinmodels(e.g.Cordiner
&Millar2009)similartotheelementalabundance. cKeady&Ridgway(1993),Cernicharoetal.(1999). dScho¨ieretal.
D
(2007a),Cernicharoetal.(1999). eScho¨ieretal.(2007b),Decinetal.(2010). fHasegawaetal.(2006),Monnieretal. o
w
(2000),Keady&Ridgway(1993). gScho¨ieretal.(2006),Keady&Ridgway(1993). n
lo
a
d
e
Table 2. C/O ratios, elemental carbon and oxygen wind(e.g.Ho¨fner2007),andtheopacityofthedustgrainswithin d
abundances((cid:3)C,(cid:3)O),andcarbonexcessesinthethree it(Woitke2006). from
metallicityenvironments. Inthehaloofourgalaxy,wherethemetallicityislowerthanthat h
intheplane,asampleof16carbon-richAGBstarshavemass-loss ttp
Environment (cid:3)C (cid:3)O C/O C−O ratesof∼4×10−6M(cid:6)yr−1,towithinafactorof3(Mauron2008). ://m
n
MW 9.05 8.94 1.3 2.5×10−4 Similar, although smaller, rates are estimated by Groenewegen, ras
LMC 9.33 8.53 6.3 1.8×10−3 Oudmaijer & Ludwig (1997) for two of these halo stars. In the .ox
SMC 9.33 8.25 12.0 2.0×10−3 metal-rich–butstillsub-solar–Sagittariusdwarfspheroidalgalaxy, ford
sixcarbonstarsaredetectedbyLagadecetal.(2010),whoestimate jo
u
mass-lossratesoftheorderof10−6M(cid:6)yr−1. rn
a
Table3. Physicalparametersofthefiducialmodels. IntheLMCandSMC,carbonstarshavemass-lossratessimilar ls.o
tothoseintheGalaxy.Tanabe´etal.(1997)foundthatdustycarbon- rg
Parameter MW LMC SMC rich LMC stars have mass-loss rates less than 10−5M(cid:6)yr−1; at N/
Leisenringetal.(2008)foundthatthebrightestMagellaniccarbon A
M˙ (M(cid:6)yr−1) 3×10−5 3×10−5 3×10−5 starswhichmadeuptheir(biased)samplehavemass-lossratesina SA
v1/e(cid:4)xp(kms−1) 12000 21000 5500 npaarrrtiocwulraarnlygehiagrhoumndas1s0-l−o6ssM-r(cid:6)ateyrc−a1rb;ovannsLtaor,oLnIe-tLaMl.C(2010831)3f,owuhnidcha Godda
G(G0) 1 2 4 has a mass-loss rate of nearly 4 × 10−5M(cid:6)yr−1. Since this star rd S
M˙ isthemass-lossrate,(cid:4) thedust-to-gasratio,vexp isinacluster,thebirthmassandmetallicitycouldbedetermined, pac
theenvelopeexpansionvelocityandGtheinterstellar allowingforanestimateofthedust-to-gasratio(seeSection2.3). e F
UVfield. VariouscarbonstarshavebeenobservedintheSMC(e.g.vanLoon lig
h
etal.1999a;Matsuuraetal.2005;vanLoonetal.2008)andrates t C
agas-to-dustratio,thetotal(gas+dust)mass-lossratecanbecalcu- are found to be similar to those in the LMC (see also van Loon tr o
n
lated.TheformermethodhasbeenroutinelyusedforGalacticstars, 2000,2006). M
but extragalactic stars are in general too faint. The availability of Sinceinthisworkweareonlyconsideringcarbonstars,forour ay
high-sensitivitydatafrom,forexample,theInfraredSpaceObser- fiducialmodelsinallthreemetallicityregimeswewillassumeM˙ = 8, 2
vatoryortheSpitzerSpaceTelescope,meansthatthesecondmethod 3 × 10−5M(cid:6)yr−1. This is towards the high end of the observed 01
3
is more practicable for extragalactic sources, and a large number range of mass-loss rates, but typical of stars where the strong in-
of authors have taken this approach – see for example van Loon fraredC2H2 featuresareseenmostclearly,andofthewell-known
etal.(1999b)orGroenewegenetal.(2009)andreferencestherein GalacticcarbonstarIRC+10216.
formass-lossderivationsforstarsintheLMCandSMC.However,
thesetwomethodsprobedifferentpartsoftheCSE(e.g.Kemper
2.3 Gas-to-dustratio
etal.2003),andoftengivedifferingvaluesoftotalmass-lossrate.
Mass-lossratesforAGBstarsintheGalaxyrangefrom10−9to In the early 1990s, Habing, Tignon & Tielens (1994) postulated
10−4M(cid:6)yr−1,withthemedianbeing3×10−7M(cid:6)yr−1(Scho¨ier that the dust-to-gas ratio would depend on metallicity, and this
&Olofsson2001;Olofsson2008b).Atypicalmass-lossratefora was refined by van Loon (2000) who determined observationally
carbonstarwitharichcircumstellarchemistrysuchasIRC+10216 that it had an approximately linear dependence. Dust production
is ≈10−5M(cid:6)yr−1 (e.g. Men’shchikov et al. 2001; Woods et al. at low metallicity is limited by the availability of heavy metals
2003). There seems to be no differentiation in the mass-loss rate (e.g. titanium) to form the condensation seeds for dust formation
based on chemistry (M-, C- and S-stars) in the Galactic sample (vanLoonetal.2008).
of Olofsson (2008b) (see Fig. 2). There is a weak dependence Therehavebeenseveralattemptstodeterminegas-to-dustratios
on metallicity for O-rich stars compared to C-rich stars, with the inthediffuseinterstellarmedium(ISM)oftheMCsfromHIand
difference appearing to be due to the driving mechanism of the infraredmaps.Athigherdensitiesthisbecomesalmostimpossible.
(cid:3)C 2012TheAuthors,MNRAS426,2689–2702
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The chemistry of extragalactic carbon stars 2693
Moreover, it is not clear how the gas-to-dust ratio in the diffuse determined.Itshouldbenotedthatgasvelocitiesanddustveloc-
ISMrelatestothatincircumstellarenvironments.IntheISMofthe ities are not necessarily well coupled, and drift velocities can be
LMC, the dust-to-gas ratio is approximately a quarter that of the substantialatlowmass-lossrates(e.g.Jones2001).
solarvalue,andthishasbeenestablishedforseveraldecades(e.g. OnecanalsolooktotheGalacticHaloasalow-metallicityen-
vanGenderen1970;Koornneef1982;Clayton&Martin1985)and vironment.Expansionvelocitieshavebeendeterminedforseveral
confirmedmorerecentlywiththeHerschelandSpitzertelescopes Halostars:Groenewegenetal.(1997)measuredanexpansionve-
(e.g. Gordon et al. 2003; Meixner et al. 2010). In the SMC, the locity of 3.2kms−1 from a 12CO J = 2–1 line, which is excep-
ISMdust-to-gasratioisaboutatenththatofsolar(vandenBergh tionally low for a carbon star. This value was confirmed by La-
1968;vanGenderen1970;Lequeuxetal.1982;Bouchetetal.1985; gadec et al. (2010), who observed the 12CO J = 3–2 line. They
Botetal.2004;Gordonetal.2009).Polycyclicaromatichydrocar- alsocalculatedalow(dust)mass-lossrateforthisobject.Lagadec
bons behave differently, but are also depleted at low metallicity etal.(2010)observedtwoothercarbonstarsintheHalo,andde-
(Sandstrometal.2012). termined expansion velocities of 6.5 and 8.5kms−1. Three other
IntheCSEsofAGBstars,thedust-to-gasratioishardertode- carbon-richobjectsintheirsamplewerethoughttobeintheGalac-
termine,especiallywithoutagoodmeasureofthegasmass.This tic Halo, but instead reside within the Galaxy’s more metal-rich
ratioalsodependsonchemistry:inoxygen-richenvelopesthede- thick disc. Expansion velocities for these stars were measured as
pendenceisroughlylinear(Marshalletal.2004).Incarbon-richen- 11.5–16.5kms−1.TheoreticalmodelsofAGBoutflows(Mattsson
D
velopes,thereissomedeclinewithmetallicity,butpossiblyslightly et al. 2008; Wachter et al. 2008) show that although low expan- o
w
shallower than linear (van Loon 2000). Estimates derived from sion velocities may seem to be related to lower metallicity, they nlo
molecularbandstrengths(e.g.Sloanetal.2006)canbemisleading are more closely related to a low carbon excess: less dust for- ad
e
(avsasunmLeodonvaertioalu.s20va0l8u)e.sI:nvlaienuLoofofinrmetdaelt.e(r2m0i0n3a)tioansssu,mauethaorgsahs-atvoe- mmaattieorinali.mplies a less efficient acceleration of the circumstellar d from
dustvalueof300–500forthecarbonstarLI-LMC1813;Leisenring Forthepurposesofourfiducialmodels,wehaveadoptedv = h
etal.(2008)assume100fortheGalaxy,200fortheLMCand500 20,10and5kms−1fortheGalaxy,LMCandSMC,respectiveexlpy. ttp://m
fortheSMC.Weadopt100,200and500,respectively,inlinewith n
themetallicityforthefiducialmodels. ras
2.5 UVradiationfield .o
x
fo
2.4 Windexpansionvelocity BLMernCaridsientgael.ne(2ra0l0s8t)rownegreerathbalenttohaetsitnabthlieshMthWa.tTthheeyISfiRnFditnhatthiet rdjou
Tofonhrcemeaettxhipoeanwnzsinioodnneh,vadseleuopnceidnteydrsgoofonnaetaChceSceEdl,uerstaht-tatioot-nigsia,nstharneadtitoec,rlmowsiehniatcolhvthienelodtcuuirtsnyt Swvauirtcihhesaanfgraoavmlearc0at.gi8ceGvva0arliiunaetdiooiffnf∼uins2efiGree0lgd.iWostnreseuntsoget3ht.h5iissGav0lasilonusemeieonnloeicunurtlmhaerorGdeegalilloainxngys.,, arnals.org/
dependsonthemetallicity.Thereisalsoalesserrelianceuponthe wherethestrengthattheinnermolecularringcanbeuptofivetimes t N
luminosityofthestar,suchthatvexp ∝ (cid:4)1/2L1/4,where(cid:4) isthe thatinthesolarneighbourhood(Paladinietal.2007).FortheSMC, AS
A
dust-to-gasratio,andLtheluminosity(Habingetal.1994;Elitzur recentworkbySandstrometal.(2010,fig.10)hasshownthatthe G
&Ivezic´ 2001;Marshalletal.2004).Thus,expansionvelocityis ISRFhasanaveragevalueof∼4G0,althoughitcanbeashighas od
d
expectedtobelowerintheMCsforagivenluminosity. 10–30G inregionsofstarformation. a
0 rd
There is also a link between expansion velocity and the mass- S
p
lossmechanism,sincepulsation-drivenwinds(whichareassociated ac
2.6 Caveats e
withlowmass-lossrates)showlowexpansionvelocities.Starsex- F
periencing a superwind, where the mass-loss driver is radiation Although we are presenting here three fiducial, and exploratory, ligh
pressureondustgrains,generallyshowlargerexpansionvelocities models,notintendedtorepresentanyparticularsource,itisworth t C
(e.g.Wintersetal.2000). noting sources of uncertainties in our models, and how they may tr o
n
Ingeneral,expansionvelocitiesintheLMChavebeenmeasured affectchemistryandpredictedlineintensities.Manyofthephysical M
a
to be fairly small compared to the Galaxy. For all chemistries, parametersabovehavemainlyaneffectontheradiusatwhichUV y
8
Galactic expansion velocities cover a wide range, from low (e.g. photons will dominate the chemistry of the envelope. Decreasing , 2
1.5–22.5kms−1; Olofsson 2008b) to high (e.g. 4.3–35.4kms−1; themass-lossrate,increasingtheexpansionvelocityofthewindor 01
3
Loup et al. 1993) and even higher velocities are expected theo- ISRFstrength,ordecreasingthedust-to-gasratio(i.e.reducingthe
retically(e.g.upto60kms−1;Mattsson,Wahlin&Ho¨fner2010). dust content) will have the effect of allowing greater penetration
ForGalacticcarbonstarsundergoingasuperwindmass-loss,typ- forUVphotons,meaningthattheenvelopechemistrytransitionsto
icalexpansionvelocitiesrangefrom13to22kms−1 (e.g.Woods a photochemistry (rather than an ion–molecule or neutral–neutral
et al. 2003). In the LMC, van Loon et al. (2003) derived v = chemistry) closer to the star. In practice, this means that parent
exp
9.5kms−1 from modelling the carbon star LI-LMC 1813. In a moleculesaredestroyedmoreeffectively,daughterspeciesarepro-
sample of oxygen-rich CSEs in the LMC, Marshall et al. (2004) ducedandthendestroyedinanarrowershell,andthedistribution
measured the expansion velocity of a number of OH masers and ofspeciesbecomesmorecondensed.Sinceabundantparentspecies
obtainedresultsmostlyintheregionof8–17kms−1.Circumstellar arefoundinthemoredensepartsoftheenvelope,wherecolumn
OHmasersarethoughttotracetheterminalvelocityofthewind, densitiesarehigh,anychangestothisthresholdradiuswillhavea
ratherthanH Omasers,forexample,whichtracetheacceleration minimaleffectontheemissionintensityfromparents.
2
zone.Wood,Habing&McGregor(1998)suggestedthatstarswith
higherexpansionvelocitieshavehighermetallicities.Morerecently,
3 CHEMICAL MODELLING
vanLoon(2000),Marshalletal.(2004)andScho¨ieretal.(2007a)
havesuggestedthatexpansionvelocitiesincreasewithevolutionon OurchemicalnetworkisbasedonthatofCordiner&Millar(2009),
theAGB.ExpansionvelocitiesforstarsintheSMChavenotbeen augmented with silicon chemistry from the UMIST Database for
(cid:3)C 2012TheAuthors,MNRAS426,2689–2702
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2694 P. M. Woods et al.
Astrochemistry2006(UDfA;Woodalletal.2007).Reactionswith millimetreregime,givenitshighabundanceandsignificantdipole
Si-bearing species are included only for those species already moment. However, it is clear from the calculations that at lower
presentinthenetworkofCordiner&Millar(2009).Theresulting metallicity(evenatjusthalfsolarmetallicityintheLMC),acety-
network contains 481 chemical species, including atoms, neutral lene (C H ) becomes the dominant carbon-bearing molecule for
2 2
molecules,cationsandanions,linkedby6174reactions,whichis R(cid:2) 1016cm(Fig.2).Intermsofcolumndensitythroughtheen-
considerablylargerthanthemostrecentreleaseofUDfA,Rate06. velope,N(C H )(cid:3)N(CO)forLMCmetallicityandisthreetimes
2 2
Speciesrangeinsizefromsingleatoms(H,He,C,Mg,N,O,S,Si) greateratSMCmetallicity.TheCOabundanceislowerinLMCand
tolonghydrocarbonchains(e.g.C ,C H ). SMCcarbonstars,andthisreducestheefficiencyofself-shielding,
23 23 2
Wefollowaparcelofgasasitpassesfromtheinneredgeofthe meaningthattheextentoftheCOenvelopeissmallerincomparison.
CSE outwards at the terminal wind velocity. Since we assume a Thepredominanceofacetyleneatlowermetallicitymeansthat
constantmass-lossprocess,thisprocedureistimeindependent,and the hydrocarbon chemistry is also more developed in such envi-
weobtainasnapshotofthechemicalstructureoftheCSEduring ronments(Fig.3).Acetyleneandtheethynyl(C H)radicalarethe
2
theAGBphase.ChemicalrateequationsaresolvedusingtheGear basis of much of the carbon-chain growth that occurs in carbon-
method for stiff differential equations (Gear 1971), resulting in richcircumstellarenvironments(Millaretal.2000),andsolarger
fractionalabundancesforallspeciesateachradialgridpoint. species such as triacetylene (C H ) experience a large boost in
6 2
production.IntheLMCandSMCmodels,N(C H )=3.5–3.9×
6 2 D
1017cm−2,whereasintheGalacticmodelthecolumndensityis30 o
4 RESULTS w
timessmaller.Thepeakabundance ofC23H2,thelargestpolyyne nlo
RadialabundanceprofilesforvariousspeciesareplottedinFigs2– inthemodel,iscomparableforthetwolowermetallicityregimes, ad
e
6,andwediscussdifferentfamiliesofmoleculesasfollows. blaurtge∼h1y0d0rotcimarebsonsmsaallsloerapfoprliethsetoGthaelaccytiacnompoodlyeyl.nTehcihsaiinnscr(eHaCse3Nin, d from
HC N,etc.)despitethereductioninelementalnitrogenabundance h
4.1 Carbonmonoxide,polyynesandcyanopolyynes 5 ttp
(Fig. 4). This would indicate that as metallicity decreases, nitro- ://m
Carbonmonoxide(CO)inGalacticcarbonstarsisthemostabun- gen is preferentially sequestered in HCN and the cyanopolyynes n
dantmoleculeafterH2,andthemostreadilyobservableinthesub- ratherthaninothernitrogen-bearingspecies.Table4,whichgives ras.o
anoverviewofthenitrogenrepositoriesinthemodelbycomparing x
fo
columndensitiesthroughtheenvelope,showsthatthisisthecase. rd
tTichemaebtuanlldicaintycetoofLHMCCxNm(eotadldlicxi≥ty3b)ustpdeeccireesaisnecsraetasSeMsfCrommeGtaallliacci-- journa
ls
ties,followingthechangesininitialfractionalabundanceofHCN .o
(Table1). arg/
Thedistributionofthecyanopolyynechainsisverydifferentto t N
thatseeninMillaretal.(2000,fig.6);however,itmatchesverywell AS
A
with the smooth density distribution model of Cordiner & Millar G
(2009). The difference arises from the inclusion of an increased od
d
photodissociationrateforHC N,resultinginasubstantialreduction a
3 rd
oftheHC3N/HC5Nratio. Sp
a
c
e
F
4.2 Anions lig
h
t C
Fcaitgeusrtehe2.GFarlaaccttiiocnmaloadbeul,ntdhaendcoestteodfTlinEepianrdeinctastepsecthiees.LTMhCesmoloiddellinaendintdhie- SIRevCe+ra1l0a2n1i6o:nsCh6Hav−ewnoawsibdeeenntifideedtecbtyedMicnCtharethGyaelatcatli.c(c2a0r0b6o)n,astnadr tr on M
confirmedbyCernicharoetal.(2007),whoalsoreportedtheinitial a
dashedlineindicatesthetheSMCmodel. y
8
, 2
0
1
3
Figure3. Fractionalabundancesofthepolyynespecies.Linecharacteristics Figure4. FractionalabundancesofHNCandcyanopolyynespecies.Line
areasinFig.2. characteristicsareasinFig.2.
(cid:3)C 2012TheAuthors,MNRAS426,2689–2702
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The chemistry of extragalactic carbon stars 2695
D
o
w
n
lo
a
Figure5. Fractionalabundancesoftheanionspecies.LinecharacteristicsareasinFig.2. d
e
d
fro
m
eplesakasrewCithHC−6Ha,nadnCdhHen−ce(Fthige.m5o).stWabeucnadnanctoamnpioanresitnhetheosbesemrvoedd- http
columnd6ensitiesand7ratiosfortheseanionswiththosecalculated ://m
n
inourGalacticmodel(alsoshowninTable5).Despitenotspecif- ra
s
ically modelling IRC+10216, the agreement between model and .ox
observationisreasonablygoodforthesmallerhydrocarbons.For ford
columndensitiesofthelargerhydrocarbonanions,C H−andlarger, jo
4 u
theagreementisnotasclose,withthemodeloverproducingbyfac- rn
a
torsof100–1000.Theabundancesofthesehydrocarbonanionsare ls.o
particularly sensitive to the initial abundance of acetylene, as has rg
beendiscussedbyRemijanetal.(2007).Theyshowthatareduction at N/
intheinitialabundanceofacetylenebyafactorof5–10bringsabout A
S
amuchbetteroverallagreement.TheC H−anionhasnotyetbeen A
detected, and we predict a column den2sity of ∼109cm−2, several Go
d
ordersofmagnitudebelowthatofthedetectedanions.Thus,unless da
Figure6. Fractionalabundancesofthesilicon-bearingspecies.Linechar- we have underestimated the radiative electron attachment rate of rd S
acteristicsareasinFig.2. C2H, C2H− is not likely to be detectable. Agu´ndez et al. (2010) pac
have determined an upper limit for this species in IRC+10216, e F
detection of C4H−. C8H− was detected by Remijan et al. (2007) of<0.0014percentoftheC2Hcolumn(Table5).Similarly,C10H− ligh
andKawaguchietal.(2007).Theanionsofthecyanopolyynerad- hasnotbeendetected,buthasasimilarion–neutralratiotoC6H− t C
icals C3N and C5N were discovered by Thaddeus et al. (2008) andC8H−,andacolumndensitywhichisaboutaquarterofthat tr o
andCernicharoetal.(2008),respectively.Mostrecently,Agu´ndez of C8H−, according to the model. Also abundant in the envelope n M
esutmalm.(a2r0y1o0f)odbesteecrvteeddtchoelusmmanlldeesntsmitoielsecaunldareastniimonatteosdoaftnee,uCtNra−l-.toA- dareensniteigeastoivfe8ly×ch1a0r1g3e–d1×car1b0o1n5ccmha−i2n.s, C−4...9, which have column ay 8, 2
anionratiosisshowninTable5.Theanionchemistryisdiscussed Thedistributionofhydrocarbonanions(Fig.5)issimilarforall 01
indetailbyCordiner&Millar(2009). membersofthefamilyfromC H− toC H−.Similardistributions 3
4 8
Ingeneral,theradiativeelectronattachmentrateofthesecarbon are also found for C N−. At the inner peak of its distribu-
1,3,5,7
chainsincreaseswithlength(Herbst&Osamura2008);however, tion,atlog(R)=15.85,CN− isformedviaelectronattachmentto
the abundance of the neutrals from which the anions are created MgNC,eveninthelowestmetallicitycase.Agu´ndezetal.(2010)
foundthatHCN+H−−→CN−+H providesaminorcontribu-
2
Table 4. Main nitrogen repositories at differing
tioninthisregion,butwefindthisreactionineffective.Themuch
metallicities. broader,outerpeakoftheCN− distributionarisesduetothereac-
tionN+C− −→ CN−+C .Ingeneral,thelargermembersof
Species MW LMC SMC 4...7 3...6
(percent) (percent) (percent) thefamilyareallformedbyelectronattachmenttocorresponding
neutrals. The distribution of CN− fits well with the observations
N2 71 37 25 of Agu´ndez et al. (2010), and the column density calculated in
HCN 14 42 58 themodel(3.0×1012cm−2)matchestheobserveddetermination
N 12 12 12 (8.0 × 1012cm−2) very closely. Similarly, the column density of
HC(cid:5)(3...11)N 1 6 3 C N−(1.6×1012cm−2)matchesthatdeterminedforIRC+10216
NH3 1 1 1 by3Thaddeusetal.(2008)exactly.Cernicharoetal.(2008)deriveda
CN <1 <1 <1
ratherlowcolumndensityforC N−of3.4×1012cm−2(lowerthan
5
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2696 P. M. Woods et al.
Table5. Columndensitiesandion–neutralratiosforanionicspeciesinIRC+10216.
Species Observedcolumn Neutral–ion Reference Calculatedcolumn Calculated
density(cm−2) ratio density(cm−2) neutral–ionratio
CN− 8.0×1012 400 Agu´ndezetal.(2010) 3.0×1012 ∼8000
C3N− 1.6×1012 200 Thaddeusetal.(2008) 1.6×1012 ∼240
C5N− 3.4×1012 2 Cernicharoetal.(2008) 1.3×1014 6
C2H− <7×1010 >70000 Agu´ndezetal.(2010) 8.1×108 ∼108
C4H− 7.1×1011 4200 Cernicharoetal.(2007) 7.2×1014 17
C6H− 4.0×1012 16–100 McCarthyetal.(2006);Cernicharoetal.(2007) 8.2×1014 4
C8H− 2.4×1012 3–4 Remijanetal.(2007);Kawaguchietal.(2007) 2.1×1014 5
forCN−),butnotethatthisfigureislikelytobeanunderestimate. Table 6. Main silicon repositories at differing
WecalculateacolumndensityofN(C N−)=1.3×1014cm−2. metallicities.
5
At lower metallicity, small anions are a factor of a few lower
inabundancethanatGalacticmetallicity.Thisissurprisinggiven Species MW LMC SMC D
(percent) (percent) (percent) ow
thatthecorrespondingneutralsaremoreabundant,andthenumber n
offreeelectronsintheMagellanicenvelopesislarger.Thelarger SiS 70 65 64 load
amneiotanlslicCit6yHt−haanndGCal7aHct−ica.repredictedtobemoreabundantatLMC SSiiO 176 216 235 ed fro
Si+ 6 4 6 hm
SiH4 2 2 2 ttp
4.3 Siliconchemistry SiC2 <1 2 2 ://m
n
Silicon chemistry was not included in the model of Cordiner & ra
s
Millar(2009),butisincludedhereduetothepotentiallyimportant .o
x
reaction 4.4 Theeffectofassumptionsaboutphysicalparameters fo
rd
C H +Si−→SiC +H , (2) Inordertofullyunderstandthechemistryatlowmetallicitydistinct jou
2 2 2 2 fromtheeffectofthedifferingphysicalconditions,wecalculated rna
whichcouldbeamajorsinkofC H intheinnerregionsoftheen- ls
2 2 threemodelsadoptingthephysicalconditionsoftheGalacticmodel .o
vweiltohpoethaetrlohwydmroectaalrlbicointys.aHreowtheevmer,ajtohrelmososdmeleschhoawnsistmhastfroeraCct2ioHn2s, (dTifafbelreen3t),maentdalluicsiintygrtehgeimineisti.aHl cehnecme,icaanlyavbaurnidatainocnesinofabtuhnedtahnrceee at Nrg/
even in the highest metallicity case (i.e. where the abundance of or column density is solely due to the chemistry of the CSE, or AS
siliconatomsishighest). A
theinitialabundancesadopted.Forclarity,wewillnamethesethe G
IRSCi+C120w2a1s6re(Cceenrtnliycdheatreocetetdabl.y2t0h1e0H),ewrshcehreelSitpsadciestOribbsuetriovnattorrayceins ‘fixedphysics’models,asopposedtothe‘fullphysics’models,as odd
describedpreviously. a
thedust-formationzoneofthestar,similartoSiOandSiSemission InTable7,wepresentthecolumndensitiesderivedfromthefull rd S
(Fonfr´ıaetal.2008).However,italsohasasignificantabundancein physicsmodels,andtheirratios.Inthefinalcolumn,wehavelisted pac
theouterenvelope(Gensheimer,Likkel&Snyder1995;Lucasetal. theratioofcolumndensitiesderivedfromthefixedphysicsmodels. e F
1pe9a9k5i)n,gasawt∼e5fin×di1n0o16ucrmmoadnedllrienagch(sienegFxi(gS.iC6)2,)w=ith5t×he1d0i−st7r.ibTuhteiosne Lusooakniningdiactahtioownothfehotwwothfienaadlocpotleudmpnhsysinicathlecotanbdlietiocnhsanagffeecgtivthees light C
figurescorrespondwellwiththeobservationsofLucasetal.(1995) chemistryofthemodel. tr o
andthechemicalmodelofCernicharoetal.(2010).Theabundance Formanyofthespecies,metallicityisthedominatingfactor,with n M
ofSiC increasesthroughthereactionbetweenacetyleneandsilicon a
2 columndensitiesfallingwithmetallicitydespiteanincreaseinthe y
batyomresac(teiqounawtioitnh2a)b,uannddanittbCe+coiomness.TdehsetrpoeyaekdaobuutnsiddaenRce=of1S0i1C7cmis overallH2columndensity.Manyofthenitrogen-bearingspeciesare 8, 20
2 includedinthisgroup,whichmakesthemgoodtracersofmetallic- 1
3
similarinallmetallicityregimes,andcolumndensitiesareslightly ity:CN,CH CN,C N−.However,thereareexceptionstothis:C N,
enhancedintheLMCandSMCmodels(1.9×1015:4.2×1015:2.6× C N,C N−3andthe5cyanopolyynesincreaseinrelativecolumnd3en-
1015; MW:LMC:SMC). The chemistry is slightly different in the 5 3
sityatLMCmetallicity,butdecreaseatSMCmetallicity.Species
LMCandSMCmodels,sincetheneutralizationofSiC H+toform
2 containingheavymetalssuchassiliconormagnesiumalsoreflect
SiC dominatesoverreaction(2).SiC H+itselfispartiallyformed
2 2 changes in metallicity, with the exception to this being SiC . For
2
fromSiC ,viathechain
2 SiC ,theeffectoftheloweringmetallicityistoproduceadecrease
2
Si++C H −→ SiC++H inthecolumndensity(seethefinalcolumnofTable7).However,
2 2
C++SiC −→ SiC++C columndensitiesinthefullphysicsmodelsincreaseforLMCand
2 2
SiC++H −→ SiC H++H SMCmetallicities,counteractingtheeffectsduetochemistryalone.
2 2 2
SiC H++e−(oranion) −→ SiC +H(+neutral). Another species withsimilar properties isCS,which is produced
2 2
inherently less in lower metallicity environments, but receives a
ThemajorrepositoryforsiliconinCSEsisSiS(Table6),inline boostinproductioninthedenserenvironmentsofthefullphysics
with the high initial abundances of this molecule (Table 1). This models.
is true in all metallicity regimes tested, although the abundance COisnotaverygoodtracerofmetallicityincarbonstars,despite
of SiS drops slightly with metallicity as sulphur is preferentially itbeingagoodtracerofoxygenabundanceunderthefixedphysics
incorporatedintoCS. conditions.Observationally,COemissionisoftenopticallythick,
(cid:3)C 2012TheAuthors,MNRAS426,2689–2702
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The chemistry of extragalactic carbon stars 2697
Table7. Calculatedcolumndensities(cm−2)forspeciesofinterest.
Species NMW NLMC NSMC Ratio Ratiofiphxyedsics
H2 2.88×1022 5.77×1022 1.15×1023 1.00: 2.00: 4.00 1.00: 1.00: 1.00
C2H 7.88×1016 7.33×1016 2.64×1016 1.00: 0.93: 0.33 1.00: 1.71: 1.78
C2H− 8.12×108 7.48×108 1.10×109 1.00: 0.92: 1.36 1.00: 0.94: 1.02
C2H2 5.28×1017 1.06×1019 2.89×1019 1.00:20.08:54.77 1.00: 7.95: 8.76
C2S 2.67×1014 3.43×1014 1.78×1014 1.00: 1.28: 0.67 1.00: 0.99: 0.50
CH3CN 1.65×1013 1.88×1012 2.81×1012 1.00: 0.11: 0.17 1.00: 0.18: 0.07
CN 2.46×1016 1.48×1016 8.84×1015 1.00: 0.60: 0.36 1.00: 0.62: 0.35
CN− 3.01×1012 2.93×1012 1.46×1012 1.00: 0.97: 0.48 1.00: 0.92: 0.47
CO 1.29×1019 9.93×1018 1.04×1019 1.00: 0.77: 0.81 1.00: 0.39: 0.20
CS 5.66×1016 7.67×1016 8.01×1016 1.00: 1.36: 1.42 1.00: 0.65: 0.32
C3H 5.92×1014 1.04×1015 6.29×1014 1.00: 1.75: 1.06 1.00: 2.01: 1.81
C3H2 3.03×1014 7.37×1014 8.19×1014 1.00: 2.43: 2.70 1.00: 2.80: 2.53
C3N 3.92×1014 6.08×1014 1.87×1014 1.00: 1.55: 0.48 1.00: 1.63: 0.95
CC33NS− 11..6043××11001125 11..7477××11001125 75..9837××11001114 11..0000:: 11..0483:: 00..4587 11..0000:: 11..1230:: 00..6602 Down
CC44HH− 71..1281××11001164 23..0004××11001164 61..5171××11001154 11..0000:: 10..6452:: 00..5145 11..0000:: 30..3909:: 30..8963 loade
CC45HN2 72..5819××11001164 41..3854××11001175 36..5258××11001174 11..0000::152..0445::102..2864 11..0000::112..6739::121..8850 d from
CCCC5686NHHH−−−2 1812....32113454××××1111000011114464 4137....18830465××××1111000011113574 9138....17521312××××1111000011112574 1111....00000000::::33023....53245134::::30230....00580795 1111....00000000::::24022....65103267::::29022....92319913 http://mnra
HCN 5.45×1017 1.44×1018 1.94×1018 1.00: 2.64: 3.55 1.00: 1.20: 0.71 s.o
x
HHCC35NN 16..4849××11001156 18..2676××11001166 53..2967××11001156 11..0000:: 14..9578:: 02..8029 11..0000:: 13..7549:: 12..0321 fordjo
HMNgNCC 32..9925××11001145 93..1011××11001145 71..4861××11001145 11..0000:: 30..0787:: 20..5436 11..0000:: 40..7379:: 30..6125 urnals
SiC 1.10×1014 5.22×1013 3.02×1013 1.00: 0.48: 0.27 1.00: 0.43: 0.24 .o
SiC2 1.93×1015 4.17×1015 2.56×1015 1.00: 2.16: 1.33 1.00: 0.92: 0.47 arg/
SiN 1.05×1014 2.48×1013 1.19×1013 1.00: 0.24: 0.11 1.00: 0.15: 0.04 t N
SiO 4.37×1016 3.49×1016 3.70×1016 1.00: 0.80: 0.85 1.00: 0.40: 0.21 A
S
SiS 1.77×1017 1.08×1017 1.03×1017 1.00: 0.61: 0.58 1.00: 0.31: 0.15 A
G
ListofmoleculesobservedbyHeetal.(2008)inIRC+10216,withtheadditionofthoseinboldface. od
d
a
rd
S
p
and thus the observed flux intensity is not representative of the 5 OBSERVABLES a
c
aCbOunidnanthceefoufllCLOM.ICnaanddditSioMnC,inmooudrelmsoadreelcs,omcopluarmabnled,endseistpieisteoaf With the impending completion of ALMA in 2013, it is interest- e Flig
ingtoconsiderthepotentialforALMAtoobservemolecularline h
factor of 2 difference in metallicity. CO is a very good tracer of t C
the molecular envelope of carbon stars, though, and gives us an ehmaviesseisotnimfraotemdelixnteraignatelancstiitcieAsGfoBrnCeSarElys,hfaolcfuasimngilloionnthlieneMsCasss.uWme- tr on
indication(seeFig.2)thatingeneral,theenvelopesofMagellanic M
ing local thermal equilibrium (LTE) (Section 5.1), and selected a
carbonstarswillonlyextend∼70percentofthedistanceofsimilar ay
fewpotentialcandidatelinesforfurtherinvestigationusingamore 8
Galacticcarbonstars. rigorousnon-LTEradiativetransfercode(Section5.2). , 2
ThecolumndensitiesofhydrocarbonssuchasC2H2,C4H,C4H2 013
and C H increase significantly at lower metallicity, as we have
6 2
seen.Table7showsthatmuchofthisincreaseisduetothechanges
5.1 LTEestimates
in chemistry at lower metallicity, with the relative proportion of
carbonincreasingoveroxygen(Table2).Asignificantincreasein FollowingthecalculationsofOlofsson(2008a),onecanwrite
columndensityforthesespeciesisalsoduetothedifferingphys- (cid:2) (cid:3) (cid:2) (cid:3) (cid:2) (cid:3) (cid:2) (cid:3)
M˙ 1.2 15 1.6 X 0.7 1 2
iacraelfcoormndeidtiothnrsoufoguhnndeuintraMl–angeeultlraanlicreaccatriboonns,satanrds.thTehheisgehsepredceines- SCO(2–1) ≈6 10−6 ve 10C−O3 D Jy, (3)
sities of Magellanic envelopes mean that these reactions proceed wheretheCO(J =2–1)linefluxdensityisgivenintermsofthe
more quickly. This trait is not observed to the same degree in mass-lossrate,expansionvelocity,fractionofCOwithrespectto
anions of hydrocarbon species, which are generally more abun- H andthedistance,D,tothestar.Insertingtypicalvaluesforthe
dant than in Galactic carbon stars for larger species (C6H− and LM2 C gives us an estimated flux density of ≈0.04 Jy. Similarly,
larger), but only by a factor of a few. This is somewhat surpris- fortheSMC,wecalculateSCO(2–1) ≈0.09Jy,sinceweexpectthe
ing, since the ionization fraction of the envelope almost doubles densityofCSEsintheSMCtobehigher.Comparablefluxdensities
from Galactic model to LMC model to SMC model. Smaller hy- areexpectedfortheCO(J=3–2)line.Suchstarsshouldbeeasily
drocarbon species (C4H−, C5H−) are more abundant in Galactic detectablewithinanhour’sobservationsforthefullALMAarray
CSEs. (5σ=6mJyat2kms−1resolution).Starswithpropertiessimilarto
(cid:3)C 2012TheAuthors,MNRAS426,2689–2702
MonthlyNoticesoftheRoyalAstronomicalSociety(cid:3)C 2012RAS
2698 P. M. Woods et al.
Table8. PredictedlinestrengthsforcarbonstarsintheLMC,usingequation(4).
Band3(89–119GHz) Band6(211–275GHz) Band7(275–370GHz)
Molecule Frequency Peakflux Molecule Frequency Peakflux Molecule Frequency Peakflux
(GHz) (mJy) (GHz) (mJy) (GHz) (mJy)
C4H2 89.315 0.1 SiO 217.105 0.7 CS 293.912 0.3
C4H2 89.687 0.1 13CO 220.399 1.0 SiO 303.927 0.3
SiC2 93.064 0.1 CN 226.632 0.2 13CO 330.588 1.7
C4H2 97.834 0.1 CN 226.660 0.6 CN 340.008 0.1
CS 97.981 0.1 CN 226.664 0.2 CN 340.020 0.1
C4H2 98.245 0.1 CN 226.679 0.2 CN 340.032 1.1
C4H2 98.655 0.1 CN 226.874 0.6 CN 340.035 0.4
C4H2 107.175 0.1 CN 226.875 0.9 CN 340.035 0.7
13CO 110.201 0.1 CN 226.876 0.4 CN 340.248 1.0
CN 113.491 0.1 CN 226.887 0.1 CN 340.248 1.4
CO 115.271 2.5 CN 226.892 0.1 CN 340.249 0.8
SiC2 115.382 0.2 CO 230.538 ≈40 CS 342.883 0.1 D
C4H2 116.105 0.1 CS 244.936 0.5 CO 345.796 ≈40 ow
n
SiO 260.518 0.5 SiO 347.331 0.1 lo
a
C2H 262.004 1.5 C2H 349.338 1.3 de
CC22HH 226622..000665 11..11 CC22HH 334499..333999 11..00 d from
HCc-2CCHN3H2 222666255...087685769 6060...771 CcH-2CCHN3H2 333465984...425090145 5008...849 http://m
n
+ ra
Band9(602–720GHz):strongestlinesincludeC2H3,HCNandCO,allwithintensitiesoftheorderof1mJy. s.o
x
fo
ourLMCstellarcharacteristicsshouldbedetectableoutto≈150kpc Carloradiativetransfercode,RATRAN(Hogerheijde&vanderTak rdjo
within 1 h, whilst stars similar to our SMC stellar characteristics 2000),toestimatethelineprofilesandlinestrengthsofmolecular urn
could potentially be seen out to distances of ≈200kpc within an lineemissionfromevolvedcarbonstarsintheLMCandSMC.Inour als
hour(Olofsson2008a). simpleestimates(seeequations4and5),weassumeLTEandanex- .org
For species other than CO, the line flux density can also be citationtemperatureof10Kforallrotationaltransitions.Inreality, a/
estimated(Olofsson2008a).ForLMCcarbonstars, theenvelopesofAGBstarshavearadiallydependenttemperature t N
A
S ≈1.62×10−9guAulfXReeQ−E(lT/kXT)X Jy, (4) aapcnecdraudtruearnceysivtoyafrsicterasulccfrtuuolrameti(o∼sne1se0me0.0gad.KJeuiwrnaitt&hheMthineosnrereirsa1ses9nuv8me1l)powtpioehnitcsoh.∼Tlimh1e0itKtsetmhine- SA God
d
using our standard parameters, and where gu, Aul and El are the theouterenvelope.Similarly,thedensitystructureexhibitsanR−2 ard
quantum mechanical degeneracy of the upper energy level of the behavioursothatintheouterenvelopethedensityismuchlower Sp
terragnysiotifotnh,etlhoewEerinesnteeirngycloeevfefilc,ireenstpefoctrivtehley.tQra(nTsiXti)oinstahnedptahretiteionn- itnhdanictahteescariltliceanledrgeynslietyveolfsrlootwateironthaalntrathnesiutipopnesr(lAeuvle/(cid:5)l,iuC)u.i,Wwehuesreedi ace Flig
function,dependentontheexcitationtemperatureofthemolecule, molecular data from the Leiden Atomic and Molecular Database h
TWXo,owdshiecthalw.2e0a0s3su).mReetiostbhee1ra0dKiusfoorfatlhleroemtatiitotinnaglrtergainosnitifoonrsea(ccfh. (cLieAntMsDanAd2)cowllhisiciohnhalasrattaebsufloartemdaennyermgyolleecvuelless,wEiinthstteriannsAiticoonesffiin- t Ctr on
molecule, taken from the molecular distributions in the chemical the(sub-)mmregionoftheelectromagneticspectrum(Scho¨ieretal. M
a
model.SimilarlyfortheSMC, 2005). y 8
S ≈3.72×10−9guAulfXReeQ−E(lT/kT)X Jy. (5) calInmooudrerlalidniga,tiwveettrraenastftehrecsatelcllualrateinovnesl,oapseinasthsephpehryisciaclalylasnydmcmheemtriic- , 2013
X withaconstantoutwardstellarwindvelocityof10kms−1 forthe
Estimatedfluxdensitiesformolecularlinesstrongerthan0.1mJy LMCand5kms−1fortheSMC.Sinceourlineprofilesarebroad-
aregiveninTable8.Thepeaklineintensitiescalculatedfromequa- enedtoawidth(cid:3)20and(cid:3)10kms−1,respectively,weuseaspectral
tions(4)and(5)generallyunderestimatelinestrengthsbyafactor resolution of 1kms−1 for both sources. We assume a distance to
of ≈5–50 in comparison to those from the non-LTE model (Sec-
sourceof50kpcfortheLMCand66kpcfortheSMC.Wegenerated
tion 5.2). The exception to this is lines of HCN, which are over-
lineprofilesandlinestrengthsforrotationaltransitionsofCO,CS,
estimated by ≈9–16 times in comparison to the NLTE estimates.
CN,HCN,SiOandSiSwhichfallintotheexpectedspectralrange
OurassumptionofTX=10KforHCNlinesisprobablyasignificant ofALMA‘FullScience’operations(≈30to≈950GHz).Theseare
underestimateforaspecieswhichisonlyabundantintheinnerre-
themoleculeswhichourcalculationssuggestedmaybesufficiently
gionsoftheCSE.AdoptingTX=75KreducesSHCN(3–2) ≈27mJy, abundanttopossessemissionstrongenoughtobeobservableand
approximatelyhalvingtheestimateforT =10K.
X forwhichcollisionalratesareavailable.UsingtheALMAsensitiv-
itycalculator,3 wedeterminedthoselinetransitionswhichmaybe
5.2 Non-LTEestimates
Inadditiontooursimpleestimatesfortheexpectedlinefluxdensity 2http://www.strw.leidenuniv.nl/moldata/
describedabove,wehaveusedtheone-dimensionalnon-LTEMonte 3almascience.eso.org/call-for-proposals/sensivity-calculator
(cid:3)C 2012TheAuthors,MNRAS426,2689–2702
MonthlyNoticesoftheRoyalAstronomicalSociety(cid:3)C 2012RAS