Chemical abundances in bright giants of the globular cluster M62 (NGC 6266)

0.82 MB·

by  David Yong

Mon.Not.R.Astron.Soc.000,1–15(2014) Printed17January2014 (MNLATEXstylefilev2.2)
Chemical abundances in bright giants of the globular
⋆
cluster M62 (NGC 6266)
4
1 David Yong1,2†, Alan Alves Brito1, Gary S. Da Costa1, Javier Alonso-Garc´ıa3,4,
0
Amanda I. Karakas1, Marco Pignatari5, Ian U. Roederer6, Wako Aoki7,
2
Cherie K. Fishlock1, Frank Grundahl8 and John E. Norris1.
n
a 1Research School of Astronomy and Astrophysics, Australian National University,Canberra, ACT 2611, Australia
J 2Stromlo Fellow

5 3Instituto de Astrof´ısica, Facultad de F´ısica, Pontificia Universidad Cato´lica de Chile, Av. Vicun˜a Mackenna 4860, 782-0436 Macul, Santiago, Chile
1 4The Milky Way Millennium Nucleus, Av. Vicun˜a Mackenna 4860, 782-0436 Macul, Santiago, Chile
5Department of Physics, University of Basel, Klingelbergstrasse 82, CH-4056 Basel, Switzerland

] 6Department of Astronomy, Universityof Michigan, 500 Church Street, AnnArbor, MI48109, USA
A 7National Astronomical Observatory, Mitaka, Tokyo 181-8588, Japan
G 8Stellar Astrophysics Centre, Department of Physics and Astronomy, Aarhus University,Ny Munkegade 120, DK-8000 Aarhus C, Denmark
.
h
p
- Accepted 2014January15.Received2014January14;inoriginalform2013November26
o
r
t
s
a ABSTRACT
[ With the exception of Terzan 5, all the Galactic globular clusters that possess signif-

icant metallicity spreads, such as ω Cen and M22, are preferentially the more lumi-
1
nousclusterswithextendedhorizontalbranches.Herewepresentradialvelocitiesand
v
chemicalabundances for sevenbright giants in the globular cluster M62, a previously
4
8 little-studied cluster. With MV = −9.18, M62 is the ninth most luminous Galactic
7 globular cluster and has an extended horizontal branch. Within our sample, we find
3 (i) no evidence for a dispersion in metallicity, [Fe/H], beyond the measurement un-
. certainties, (ii) star-to-star abundance variations for C, O, Na and Al with the usual
1
correlationsbetweentheseelementsasseeninotherglobularclusters,and(iii)aglobal
0
enrichmentfortheelementsZr,BaandLaatthelevel[X/Fe]≃+0.4dex.Forelements
4
1 heavier than La, the abundance ratios are consistent with the scaled-solar r-process
: distribution. Below La, the abundances are anomalous when compared to the scaled-
v
solars-processorr-processdistributions.Fortheseelements,theabundancesignature
i
X in M62 is in agreement with predictions of the s-process from fast-rotating massive
stars,althoughthe high[Rb/Y]ratio wemeasure maybe achallengeto this scenario.
r
a
Key words: Stars:abundances–Galaxy:abundances–globularclusters:individual:
NGC 6266
1 INTRODUCTION a range in metallicity1 from [Fe/H] 2.0 to [Fe/H]
≃ − ≃
0.5; for certain elements, the abundance ratios with re-
A small, but growing, number of Galactic globular clus- −
spect to iron, [X/Fe], exhibit considerable variations with
ters exhibit a star-to-star variation in the relative abun-
metallicity(e.g.,Norris & DaCosta1995;Smith et al.2000;
dances of the heavy elements. ω Centauri is the most
Johnson & Pilachowski 2010; D’Orazi et al. 2011). As the
well-known member of this group with its stars spanning
most massive Galactic globular cluster and in light of its
complex chemical enrichment history, ω Cen is regarded to
bethecoreofadisrupteddwarfgalaxy(e.g.,Freeman1993;
Bekki& Freeman 2003; Bekki& Norris 2006).
M54isthesecondmost massiveGalactic globularclus-
⋆ Based in part on data collected at Subaru Telescope, which
isoperated bytheNationalAstronomicalObservatoryofJapan.
This paper includes data gathered with the 6.5 meter Magellan 1 Hereand throughout the paper, ironisthe canonical measure
Telescopes locatedatLasCampanasObservatory,Chile. of metallicity and we adopt the standard spectroscopic notation
† E-mail:[email protected] [Fe/H]=log10(NFe/NH)⋆−log10(NFe/NH)⊙
2 D. Yong et al.
ter. It displays a key characteristic shared by ω Centauri clusterandisatypicalOosterhofftypeIsystem.Thiscluster
butabsentin themajority ofclusters,namely,ametallicity liesinthevicinityofthebulge,althoughproper-motionmea-
dispersion (Carretta et al. 2010a), albeit of smaller ampli- surementsindicatethatitprobablybelongstothethickdisk
tude than that found in ω Centauri. That M54 is the nu- (Dinescu et al.2003).M62 ishighly reddened,E(B V) =
−
clear star cluster of theSagittarius dwarf spheroidal galaxy 0.47 (Harris 1996),and it is severely affected by differential
lends further support to the hypothesis that massive glob- reddening,∆E(B V) 0.25 (Alonso-Garc´ıa et al. 2012).
− ∼
ular clusters with metallicity variations may be the nuclei The goal of the present paper is to conduct a de-
of disrupted dwarf galaxies (although Bellazzini et al. 2008 tailed chemical abundance analysis of M62 based on high-
argue that the nucleus of Sgr likely formed independently resolution, high signal-to-noise ratio spectra. To ourknowl-
of M54). When considering the ensemble of globular clus- edge this is the first such analysis of this cluster and this
ters with unusual chemical abundance ratios2 (dispersions workthereforerepresentsanessentialsteptowardscomplet-
in metallicity and/or neutron-capture elements), it is clear ing the census and characterisation of the Galaxy’s most
thattheseobjectsarepreferentiallythemoreluminous,and massive globular clusters.
therefore the more massive, clusters.
Lee et al. (2007) considered globular clusters with an
extended horizontal branch (i.e., a horizontal branch with
2 OBSERVATIONS AND ANALYSIS
extremely blue stars) which likely signifies the presence of
He abundance variations (D’Antonaet al. 2002). They ar- 2.1 Target selection and spectroscopic
gued that not only are clusters with an extended horizon- observations
tal branch preferentially more massive, but also when con-
Six targets were selected from Rutledgeet al. (1997) (see
sidered together theclusters havedifferent kinematics from
Table 1). All were high probability members based on the
othergroupingsoftheGalactic globularclusterpopulation.
publishedradialvelocities andwereobservedwith theHigh
Thesampleofclusterswithextendedhorizontalbranchesin-
Dispersion Spectrograph(HDS,Noguchiet al.2002)onthe
cludeω Cen, M54 and M22, all of which exhibitmetallicity
Subaru Telescope on 2013 July 17. The exposure times
variations. Lee et al. (2007) showed that such clusters had
ranged from 16 to 40 minutes per program star. We used
kinematics consistent with accretion and thus likely have a
theStdYbsettingandthe0′.′8slitwhichresultedinawave-
different origin to thebulk of the cluster population.
lengthcoveragefrom 4100˚Ato 6800˚Aataspectralreso-
The identification of additional globular clusters with ∼ ∼
lution ofR=45,000. Therewere approximately fourpixels
peculiarchemicalabundancepatterns(e.g.,star-to-starvari-
per resolution element. Spectra of a telluric standard (HD
ations in metallicity) would advance our understanding of
163955) and a radial velocity standard (HD 182572) were
the range and relative frequency of the abundance spread
also taken duringthis observing run.
phenomenon in the Galactic globular cluster system. Ad-
Two additional targets were selected from Str¨omgren
ditionally, quantifying the number of accreted systems has
photometry obtained by J.A-G and observed using the
important consequences for understanding the formation
MIKEspectrograph(Bernstein et al.2003)attheMagellan-
and evolution of our Galaxy (Freeman & Bland-Hawthorn
II Telescope on 2012 June 26. The exposure times were 1
2002). To this end, the most obvious objects to study are
hour for each star. The 0′.′7 slit was employed providing
the most massive globular clusters with extended horizon-
a spectral resolution of R = 35,000 (blue CCD) and R =
tal branches. M62 (NGC 6266) is the ninth most luminous
27,000(redCCD)withawavelengthcoveragefrom 3800˚A
cluster (Harris 1996) and has an extended blue horizontal ∼
to 9000˚A. There were approximately four pixels per res-
branch.IntheF439W F555WversusF555Wcolourmag- ∼
− olution element. A telluric standard (HD 170642) was also
nitude diagram presented by Piotto et al. (2002), the hori-
observed.
zontal branch extends down in luminosity to the main se-
In Figure 1, we show the locations of our targets in
quenceturn-off.Theextendedhorizontalbranchisalsoread-
JHK colour-magnitude diagrams (CMD). The photometry
ily identified in dereddened V versus B V and V versus
− is from Valenti et al. (2007) and 2MASS (Skrutskieet al.
V Idiagrams(Alonso-Garc´ıa et al.2011,2012).M62ranks
− 2006).
fifth in the number of millisecond pulsars (Cocozza et al.
2008), and there may be a connection between multiple
stellar populations and the numbers of millisecond pulsars
2.2 Data reduction and analysis
(Lanzoni et al.2010).Contreras et al.(2010)findthatM62
could contain the most RR-Lyrae of any Galactic globular TheSubaruspectrawerereducedusingIRAF3.TheMagel-
lan spectra were reduced using a combination of IRAFand
themtoolspackage4.Inbothcases,theapproachwassimilar
2 In addition to ω Cen and M54, the other clusters include tothatdescribedinYonget al.(2006).InFigure2,weplot
M15 (Sneden etal. 2000; Otsukietal. 2006; Sobecketal. 2011; a region of the spectra for the program stars. The telluric
Worleyetal. 2013), NGC 362 (Carrettaetal. 2013), NGC 1851 standards were used to correct regions affected by telluric
(Yong&Grundahl 2008; Villanovaetal. 2010; Carrettaetal.
2011), NGC 2419 (Cohenetal. 2010, 2011; Cohen&Kirby
2012;Mucciarellietal.2012),NGC3201(Simmereretal.2013), 3 IRAFisdistributedbytheNationalOpticalAstronomyObser-
NGC 5824 (Savianeetal. 2012), NGC 6656 (DaCostaetal. vatories,whichareoperatedbytheAssociationofUniversitiesfor
2009; DaCosta&Marino 2011; Marinoetal. 2009, 2011; Research in Astronomy, Inc., under cooperative agreement with
Roedereretal. 2011; Alves-Britoetal. 2012), NGC 7089 theNationalScienceFoundation.
(Lardoetal. 2012, 2013) and Terzan 5 (Ferraroetal. 2009; 4 www.lco.cl/telescopes-information/magellan/instruments/mike/iraf-
Origliaetal.2011,2013) tools/iraf-mtools-package
Chemical abundances in NGC 6266 3
Table 1.Programstars,radialvelocitiesandstellarparameters.
Namea RA2000 Dec.2000 Jb Hb Kb RV σRVc Teff logg ξt [m/H] [Fe/H]
(kms−1) (kms−1) (K) (cms−2) (kms−1) dex dex
(1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11) (12) (13)
SubaruTelescopeObservations20130717
A5 170116.21 −300331.81 9.782 8.847 8.583 −75.0 0.2 3950 0.50 1.79 −1.10 −1.12
A7 170117.50 −300320.60 9.602 8.608 8.389 −89.3 0.5 3925 0.20 1.89 −1.10 −1.15
A14 170117.29 −300358.69 11.476 10.710 10.529 −54.2 0.2 4400 1.40 1.53 −1.10 −1.10
A27 170125.19 −300409.85 11.779 11.124 10.989 −69.9 0.4 4625 1.50 1.76 −1.20 −1.24
A124 170100.28 −300738.28 11.449 10.756 10.536 −68.2 0.3 4450 1.00 1.64 −1.20 −1.19
A195 170112.63 −300340.78 11.988 11.318 11.157 −0.1 0.4 ... ... ... ... ...
MagellanTelescopeObservations20120626
c53 170115.20 −300836.50 10.101 9.246 ... −73.6 0.2 4450 1.40 1.78 −1.10 −1.11
d10 170116.80 −300818.00 10.125 9.145 ... −59.0 0.2 4175 0.90 1.90 −1.10 −1.15
a AXXXnamesarefromAlcaino(1978)whilec53andd10arefromStr¨omgrenphotometrytaken byJ.A-G.
b Photometry fortheAXXXstarsisfrom2MASS(Skrutskieetal.2006).Photometryforc53andd10arefromValentietal.(2007).
c Theuncertainties arethestandarderrorofthemean.
star against the radial velocity standard, HD 182572, for
whichweadoptedavalue5 of 100.35 kms−1.Forthepro-
−
gramstarsobservedwithMagellan,theradialvelocitieswere
determined by measuring the positions of a large number
(several hundred)of atomic absorption features. The radial
velocities are presented in Table 1 and we note that star
A195 is not a radial velocity member.
Equivalent widths (EW) were measured using routines
inIRAFandDAOSPEC(Stetson & Pancino2008).Wenote
that every line measured using DAOSPEC was also mea-
sured, and visually inspected, using IRAF. For all stars,
therewasgood agreementbetweenthetwosetsofmeasure-
ments; for a given star, the average difference in EWs be-
tween IRAF and DAOSPEC was 0.1 m˚A with a dispersion
of 3 m˚A. The line list and EW measurements (average of
IRAFand DAOSPEC)are presented in Table 2.
The stellar parameters were determined using a tradi-
tional spectroscopic approach. One dimensional local ther-
modynamic equilibrium (LTE) model atmospheres were
taken from the Castelli & Kurucz (2003) alpha enhanced,
[α/Fe] = +0.4, grid. To produce particular models we used
the interpolation software tested in Allende Prieto et al.
Figure 1. JHK colour-magnitude diagrams for M62 from (2004). Chemical abundances were generated using the
Valentietal. (2007) (left panels) and 2MASS Skrutskieetal. LTE stellar line analysis program MOOG (Sneden 1973;
(2006) (right panels; taking only objects within a radius of 3.5’ Sobeck et al.2011).Theeffectivetemperature,Teff,wasob-
and“AAA”photometricqualityflags).Theaquastarrepresents tainedwhentherewasnotrendbetweentheabundancefrom
the radial velocity non-member, star A195, and the red circle Fei lines and the lower excitation potential. The surface
represents star A124, the possible non-member based on chemi- gravity,logg,wasestablishedbyforcingtheironabundance
calabundances.Allotherprogramstarsareshownasgreycircles. derived from Fei and Feii lines to be equal. The microtur-
In the left panels, the solid line represents the mean RGB ridge
bulent velocity, ξt, was adjusted until the abundance from
linefromValentietal.(2007). Fei linesexhibitednotrendwith reducedequivalentwidth,
log(EW/λ). Finally, we required the derived metallicity to
be within 0.1 dex of the value adopted in the model at-
absorption. Star A195 had a signal-to-noise ratio S/N = 40 mosphere. The final stellar parameters (see Table 1) were
per pixel near 6000˚A. Excluding this star, the S/N ranged established only when all conditions above were simultane-
fromS/N=65perpixel(starsA14andA27)toS/N=110 ously satisfied. The uncertainties in Teff, logg and ξt are
per pixel(star d10).
For the program stars observed with Subaru, heliocen-
tricradialvelocitiesweremeasuredbycross-correlatingeach 5 obswww.unige.ch/∼udry/std/stdnew.dat
4 D. Yong et al.
Table 2.Linelistandequivalentwidths.
Wavelength Speciesa L.E.P loggf A5 A7 A14 A27 A124 c53 d10 Sourceb
˚A eV m˚A m˚A m˚A m˚A m˚A m˚A m˚A
(1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11) (12)
6300.31 8.0 0.00 −9.75 33.0 48.9 <15.0 67.0 48.8 <15.0 60.2 B
6363.78 8.0 0.02 −10.25 ... 17.0 ... 21.2 22.1 ... 34.4 A
7771.95 8.0 9.15 0.35 ... ... ... ... ... ... 12.8 B
7774.18 8.0 9.15 0.21 ... ... ... ... ... ... 10.1 B
5682.65 11.0 2.10 −0.67 ... ... 114.8 51.2 ... 110.3 98.0 B
a Thedigits tothe leftof the decimalpoint aretheatomic number.The digitto the rightof thedecimal point isthe ionizationstate (0
=neutral,1=singlyionized).
b A=loggf valuestakenfromYongetal.(2005)wherethereferencesincludeDenHartogetal.(2003),Ivansetal.(2001),Kurucz&Bell
(1995),Prochaskaetal.(2000),Ram´ırez&Cohen(2002);B=Gratton etal.(2003);C=OxfordgroupincludingBlackwelletal.(1979a,b,
1980, 1986, 1995); D = Biemontetal. (1991); E1 = Fuhr&Wiese (2009), E2 = Roederer&Lawler (2012), E3 = Bi´emontetal. (2011),
E4=Biemontetal.(1981),E5=Lawleretal.(2001a),E6=Lawleretal.(2009),E7=Lietal.(2007),E8=DenHartogetal.(2003),
E9=Lawleretal.(2006),E10=Lawleretal.(2001b),E11=Roedereretal.(2012);F=Lambert&Luck(1976)
Thistableispublishedinitsentiretyintheelectroniceditionofthepaper.Aportionisshownhereforguidanceregardingitsformand
content.
catedthatitisasuper-solarmetallicityclumpgiant(T =
eff
4750, logg = 3.1, [Fe/H] = +0.17).
Chemical abundances were determined using the final
model atmosphere, measured EW and MOOG. For Cu and
theneutron-captureelements,abundanceswere determined
via spectrum synthesis. Lines affected by hyperfine split-
ting (HFS) and/or isotope shifts (e.g., Scii, Vi, Mni, Coi,
Rbi, Baii, Laii, Prii, Euii and Pbi) were treated appro-
priately using the HFS data from Kurucz& Bell (1995) or
othersourcesasnotedinTable2.ForRbandPb,thewave-
length coverage necessary to measure these elements was
only obtained in the two stars observed with Magellan (see
Figures3and4).Cabundancesweredeterminedfromspec-
trum synthesis of the 4300˚A CH molecular lines using the
CH line list compiled by B. Plez et al. (2009, private com-
munication).Inouranalysis,thedissociation energyforCH
was 3.465 eV. The Asplund et al. (2009) solar abundances
wereadoptedandthechemicalabundancesarepresentedin
Table 3.
The uncertaintiesin thechemical abundanceswere ob-
tained in the following manner. We repeated the analysis
andvariedthestellarparameters,oneatatime,bytheirun-
certainties.Wealsoconsideredtheuncertaintyofthemetal-
licity used to generate the model atmosphere, [M/H], and
adjusted this value by +0.2 dex. The systematic error was
Figure 2. A portion of the spectra for the program stars.
determinedbyaddingthesefourerror termsinquadrature,
The yellow region highlights two Zr lines used in the analysis.
The positions of other atomic lines and the stellar parameters and we assume these values are symmetric for positive and
(Teff/logg/[Fe/H])areincluded. Theredspectrum isstarA124, negative changes. Following Norris et al. (2010), the ran-
apossiblenon-memberbasedonthechemicalabundances. dom error (s.e.logǫ) was obtained by taking max(s.e.logǫ,
0.20/√N ). The second term is what would be expected
lines
for aset of N with adispersion of 0.20 dex (aconserva-
lines
tive value based on the abundance dispersion exhibited by
estimated to be 50 K, 0.20 dex and 0.20 km s−1, respec- Fei lines). By adding the systematic and random errors, in
tively.WedonotpresentresultsforstarA195sinceitisnot quadrature, we obtain the total error which is presented in
aradial velocitymemberandourpreliminary analysis indi- Table 3.
Chemical abundances in NGC 6266 5
Table 3.Chemicalabundances fortheprogramstars.
Name A(X) Nlines s.e.logǫ [X/Fe] TotalError
CfromCH
A5 6.46 synth ... −0.85 0.21
A7 6.31 synth ... −0.97 0.23
A14 6.41 synth ... −0.92 0.21
A27 6.76 synth ... −0.43 0.22
A124 6.76 synth ... −0.48 0.23
c53 6.71 synth ... −0.61 0.22
d10 7.01 synth ... −0.27 0.22
Oi
A5 7.79 2 0.07 0.30 0.18
A7 7.79 2 0.04 0.39 0.20
A14 7.68 1 ... <0.16 ...
A27 8.40 2 0.09 1.05 0.18
A124 8.08 2 0.03 0.68 0.18
c53 7.66 1 ... <0.17 ...
d10 8.31 4 0.04 0.89 0.14
Nai
A5 5.67 2 0.04 0.64 0.18
A7 5.63 2 0.00 0.68 0.17
A14 5.79 3 0.04 0.72 0.13
A27 4.87 3 0.05 −0.04 0.13
A124 4.86 2 0.01 −0.09 0.15
c53 5.59 3 0.06 0.55 0.12 Figure 3. Observed and synthetic spectra forstars c53 (upper)
d10 5.18 4 0.04 0.21 0.12 andd10(lower)nearthe7800˚ARbfeature.Thethickblackline
representsthebestfit,thinredandgreenlinesare∆[Rb/Fe]=±
Thistableispublishedinitsentiretyintheelectroniceditionof
0.3dexfromthelineofbestfitandthedottedlinehasnoRb.The
thepaper.Aportionisshownhereforguidanceregardingits
stellarparameters(Teff/logg/[Fe/H])and[Rb/Fe]abundanceare
formandcontent.
includedineachpanel.
3 RESULTS
thefollowing constraints: (i) 95 6 RV6 45 kms−1 and
3.1 Cluster membership and radial velocities − −
(ii) 1.35 6 [Fe/H] 6 1.00 dex. And finally, we counted
− −
WiththeexceptionofstarA195, allprogram stars havera- the numbers of stars that lay in a particular region in the
dial velocities consistent with cluster membership. Accord- Teff-logg plane, specifically, the area is bounded at the left
ing to the Harris (1996) catalogue, M62 has a heliocentric edgebythelinefrom (Teff,logg)=(4925,2.0) to(Teff,logg)
radial velocity of 70.1 1.4 km s−1 and a central veloc- =(4225,0.0),attherightedgebythelinefrom(Teff,logg)=
ity dispersion of 1−4.3 0±.4 km s−1. For our seven program (4200,2.0) to (Teff,logg) = (3625,0.0) both with 0.0 6 logg
stars,wefindanavera±geheliocentricradialvelocityof 69.9 62.0.Wefound490starsintheBesan¸conmodelthatsatis-
4.3 km s−1 (σ = 11.4 km s−1). − fiedallcriteriaandthereforeestimatethatgivenasampleof
± As we shall discuss in the following subsection, star starsoccupyingsimilarlocationsinJHK colour-magnitude
A124haslow abundanceratios fortheneutron-captureele- diagrams astheprogram stars, theprobability of observing
mentscomparedtotheotherprogram starsandsothepos- afieldstarwithstellarparametersandaradialvelocitycon-
sibilityexiststhatthisobjectmaynotbeaclustermember. sistentwiththeprogram starsisroughly4%.Inoursample
Witharadialvelocityof 68.2 0.3kms−1,excludingthis of seven stars, having one field star with stellar parameters
starwouldnotsubstantia−llycha±ngethemeanradialvelocity and a radial velocity consistent with the cluster sample is
and velocity dispersion we measure for this cluster. not unreasonable given the small number statistics. Accu-
To investigate the likelihood of observing field stars rate proper-motion and parallax measurements from GAIA
in the vicinity of M62 with stellar parameters (T , logg would establish cluster membership, or otherwise, for this
eff
and [Fe/H]) and radial velocities comparable to the pro- object.
gramstars,wemakeuseoftheBesan¸conmodel(Robin et al.
2003).First,weconsiderallstarswithinaonesquaredegree
3.2 Chemical abundances
fieldcenteredonM62.Secondly,werestrictedthesampleto
lie near the RGB in the K versus J K CMD; specifi- For the program stars, we measure a mean metallic-
−
cally,theleftedgewasdefinedasthelinefrom(J K,K)= ity of [Fe/H] = 1.15 0.02 (σ = 0.05). Our metal-
− − ±
(0.55,12.0)to(J K,K)=(1.1,8.0)andtherightedgeasthe licity is consistent with previous measurements based
−
linefrom(J K,K)=(1.05,12.0) to(J K,K)=(1.4,8.0). primarily on low-resolution spectroscopic measurements:
− −
Wefind13281suchstarsfromtheBesan¸conmodel.Thirdly, [Fe/H] = 1.28 0.15 (Zinn & West 1984); [Fe/H] =
− ±
of these 13281 stars we counted the number that satisfied 1.02 0.04 (Carretta & Gratton 1997); [Fe/H] = 1.19
− ± −
6 D. Yong et al.
Figure 5. Abundance ratios for combinations of the light ele-
ments (C, O, Na, Mgand Al). The reddata point is star A124,
which has low neutron-capture element abundance ratios com-
pared to the remaining program stars. The dashed line is the
linear fit to the data (slope and error are included). The aver-
Figure 4. Observed and synthetic spectra for stars c53 (up- ageerror(<σ[X/Fe]>)anddispersion(σ)inthex-directionand
per)andd10(lower)nearthe4058˚APbfeature.Thethickblack y-directionareincluded.
line represents the best fit (d10) or abundance limit (c53), thin
red and green lines are ∆[Pb/Fe] = ± 0.5 dex from the line of
best fit and the dotted line has no Pb. The stellar parameters
captureelementsfromSrtoEuforA124,wefindanaverage
(Teff/logg/[Fe/H]) and [Pb/Fe] abundance are included in each
panel. difference of 0.29 0.02 (σ = 0.07). This difference is ap-
±
proximately twice the average error and suggests that this
starischemicallydifferentfromtheothersixprogramstars,
(Kraft & Ivans 2003) when using MARCS model atmo- at least for the neutron-capture elements. Thus, our quan-
spheres (Gustafsson et al. 1975); [Fe/H] = 1.18 0.07 titativeanalysis confirmsthevisualimpression from Figure
− ±
(Carretta et al. 2009). The dispersion in [Fe/H] based on 2 in which star A124 has weaker Zr lines relative to other
our program stars can be entirely attributed to the mea- programstarsofcomparablestellarparameters.Asidefrom
surement uncertainty. The mean metallicity and dispersion starA124,wefindnocompellingevidenceforanabundance
areessentiallyunchangedifweexcludethe“neutron-capture dispersionfortheneutron-captureelements.Inallcases,the
poor” star A124 with [Fe/H] = 1.19. observed abundance dispersion is consistent with the mea-
−
For the light elements, we find evidence for a star-to- surementuncertainty.Wefindnoevidenceforastatistically
star variation for C, O, Na and Al and perhaps Mg (see significant correlation between any of the heavy elements
Figure 5). For O, Na and Al, the observed abundance dis- andNa; such acorrelation hasbeen identifiedin theglobu-
persion exceeds the average measurement error by a factor lar cluster NGC 6752 (Yonget al. 2013).
of 2 indicating a genuine abundance spread for these el- In Figure 7, we compare the measured abundance dis-
∼
ements. These abundance dispersions are confirmed by the persion (including star A124) with the average total error.
statistical significance of the correlations between the fol- As already noted, the light elements O, Na and Al exhibit
lowing pairs of elements, > 3-σ for [O/Fe] vs. [Na/Fe] and significantlylargerdispersionsthanexpectedgiventhemea-
[Na/Fe] vs. [Al/Fe]. There∼is a suggestion that Mg and Al surementuncertainty.Forallotherelements,withthepossi-
are anticorrelated, a characteristic found in other, but not bleexception of C, theobserved abundancedispersions can
all,globularclusters.Thereisnoevidenceforanabundance beattributedtothemeasurementuncertainty.Notethatthe
variationforSi,norforcorrelationsbetweenSiandallother abundancesfor Feiiwereforced tomatch thoseof Fei,and
elements as seen in NGC 6752 (Yonget al. 2013). this species is not included in Figure 7. Had we excluded
In Figure 6 we plot Ni and a selection of neutron- star A124 when generating this figure, the σ(observed) val-
capture elements versus Na. Star A124, marked in red this ues would have decreased such that a larger proportion of
and other figures, exhibits systematically lower abundance thedatawouldliebelowthe1:1relation.Assuch,ourmea-
ratiosfortheneutron-captureelementscomparedtothere- surement uncertainties might beoverestimated.
maining program stars. For a given element, we can com- In Figure 8, we plot [X/Fe] vs. [Fe/H] for the pro-
pare the [X/Fe] ratio for star A124 to the average [X/Fe] gramstarsandfieldstarsfromFulbright(2000).ForNaand
ratio for the other stars. When doing this for the neutron- Al, one can clearly see that the abundance dispersions for
Chemical abundances in NGC 6266 7
[Na/Fe]and[Al/Fe]inM62farexceedthatofthefieldstars
ofcomparablemetallicityandthatthelowerenvelopeofthe
cluster values are in good agreement with field stars. Such
resultshavebeenwellestablishedforotherclusters(e.g.,see
reviews by Kraft 1994; Gratton et al. 2004; Gratton et al.
2012). We note that star A124 has Na and Al abundances
in agreement with field stars. Further, the neutron-capture
elementabundancesforthisstar,whilelowerthantheclus-
ter average, appear typical for field halo stars at the same
metallicity, which is consistent with it being a field inter-
loper.
ForSi,Ca,Ti,CrandNi,theabundanceratiosforM62
areingood agreement with field starsofthesamemetallic-
ity.In solar system material, elements mainly produced via
the s-process include Y and Zr, and we refer to these as
s-process elements. Similarly, elements predominantly pro-
ducedviather-processinsolarsystemmaterialarereferred
to as r-process elements and include Eu. For the s-process
elements Y and Zr, the M62 values lie slightly above field
halo stars of comparable metallicity. For the r-process ele-
mentEu,theM62valueslieneartheupperenvelopedefined
by the field stars. (We caution that zero-point offsets may
reduce, or amplify, the abundance differences between M62
and thefield stars.)
InFigure8,wealso includethecomparable metallicity
globularclustersM4([Fe/H]= 1.08;Ivanset al.1999)and
−
M5([Fe/H]= 1.21; Ivanset al.2001)usingthedatafrom
−
Yonget al. (2008a,b). For Na and Al, while the abundance
dispersions found in M4 and M5 are considerably smaller
than in M62, we note that our sample sizes are small. Ex-
amination of the larger samples in Ivanset al. (1999, 2001)
and Marino et al. (2008) also indicate that the abundance
dispersionsforNaandAlinM4andM5areprobablysmaller
Figure 6.SameasFigure5butforNiandtheneutron-capture thanthoseofM62.FortheelementsSitoNi,theabundance
elementsversusNa. ratios [X/Fe] in M4 and M5 match those of M62 and field
stars in their metallicity range. For the s-process elements,
M4 is known to havelarger [X/Fe] ratios than M5.
We conclude this section by noting some physical pa-
rameters for M4, M5 and M62 which may provide in-
sight into any chemical abundance differences between
these clusters. DeAngeli et al. (2005) measured relative
ages for these three clusters and found M5 to be 10-15%
younger6 than M4 and M62 with the latter two being co-
eval.Mar´ın-Franch et al.(2009)alsofoundM5tobeslightly
younger than M4. On the other hand, VandenBerg et al.
(2013) measured ages for M4 andM5 of 11.50 Gyr,but did
not study M62. The absolute luminosities for M4, M5 and
M62are 7.19, 8.81and 9.18,respectively,andthecen-
tralveloc−itydisp−ersionsare−4.0,5.5and14.3kms−1,respec-
tively (Harris 1996). M4 and M62 have orbits that are re-
strictedtotheinnerdiskandbulge,apocentricradiiofRa=
Figure 7. Measured abundance dispersion, σ(observed, includ- 5.9kpc(Dinescu et al.1999)andRa =3.1kpc(Allen et al.
ingstarA124)versusaveragemeasurementerror,σ(model).The 2006),respectivelywhileM5hasanorbitwhichmaybecon-
dotted red line is the 1:1 relation. Feii is not included (see text sistent with an accretion event,Ra = 35.4 kpc.
for details). For clarity, some values have been shifted by up to
0.01dex.
6 DeAngelietal.(2005)regardtheirtypicaluncertainties tobe
5%to8%.
8 D. Yong et al.
Figure 8. Abundance ratios [X/Fe] vs. [Fe/H] for the program stars (the red circle is star A124 with low neutron-capture elements
ratios). For Ti,weadopt the average ofTiiandTiii. A representative errorbar is includedineach panel. Thesmallgreysymbols are
field halo stars taken from Fulbright (2000). The shaded aqua and pink regions represent the data (mean value ± standard deviation)
forM4andM5,respectively, takenfromYongetal.(2008a,b).
4 DISCUSSION IQR[Mg/Al] = 0.89. With MV = 9.18 (Harris 1996), the
−
IQRvaluesforM62areconsistentwiththoseofotherlumi-
ThefocusofthisdiscussionistoexaminethenatureofM62
nousclusters.
in light of its chemical abundances. In Table 4, we present
theaverage cluster abundanceratios with and without star
A124,whichhaslowerneutron-captureelementabundances. 4.2 The neutron-capture elements
In Figure 9 we plot [X/Fe] versus atomic number for M62,
M4 and M5. Regarding the s-process elements, Zr, Ba and
4.1 The light elements
La have ratios of [X/Fe] +0.4. Such values differ from
≃
It is now well established that every globular cluster ex- the majority of field halo stars of similar metallicity and
hibitsstar-to-star abundancevariations for C, N,O,F,Na, from the comparable metallicity globular cluster M5. On
Mg and Al (e.g., see reviews by Kraft 1994; Gratton et al. the other hand, these abundance ratios in M62 are compa-
2012). These abundance variations are believed to have rable (but not identical) to M4, a cluster believed to have
been produced via hydrogen-burning at high tempera- formedfromgasenrichedins-processmaterial.Itisintrigu-
ture in asymptotic giant branch (AGB) stars, fast-rotating ingthatbothM4andM62areinnerdiskglobularclustersof
massive stars and/or massive binaries (e.g., Fenneret al. similar metallicity. Enhancementsin thes-processelements
2004; Ventura& D’Antona 2005; Decressin et al. 2007; are also seen in the globular clusters ω Cen and M22 (e.g.,
deMink et al. 2009; Marcolini et al. 2009). DaCosta & Marino 2011, and references therein). For ele-
Carretta (2006) and Carretta et al. (2010b) found that ments other than Zr, Ba and La, the abundance pattern in
the abundance dispersions for the light elements corre- M62 more closely resembles that of M5 rather than M4.
late with various physical parameters. In particular, they At the metallicity of these clusters, published neutron-
notedthattheinterquartilerange(IQR)for[O/Fe],[Na/Fe], capture element abundances in inner disk and bulge stars
[O/Na] and [Mg/Al] exhibits a correlation with the abso- are limited. The Fulbright (2000) comparison field stars
lute luminosity, i.e., total mass. Within our modest sam- in Figure 8 are solar neighbourhood objects. Since the pi-
ple, assuming for the purposes of this exercise that the oneering work by McWilliam & Rich (1994) on chemical
[O/Fe]limitsaredetections,wefindIQR[O/Na]=1.11and abundances in the bulge, the handful of papers presenting
Chemical abundances in NGC 6266 9
Figure 9. Abundance ratios [X/Fe] vs. atomic number for M62 (large grey circles), M4 (small aqua squares) and M5 (small pink
triangles).ForM62,weexcludestarA124.ForM4andM5,thedataarefromYongetal.(2008b).ForSr,weapplyanon-LTEcorrection
of0.29dex(seeSection4.2fordetails)andtheM4andM5Srvaluesarenotplottedsincetheyarederivedfromadifferentlinethanin
thepresentwork.
neutron-capture element abundances at or below [Fe/H] = values of [Rb/Fe] = +0.09 and +0.13 for stars c53 and
1.0intheinnerdiskorbulgeincludeJohnson et al.(2012) d10, respectively. As can be seen in Figure 9, those values
−
and Bensby et al. (2013). Although their sample sizes are fall closer to M5 ([Rb/Fe] = 0.00) than to M4 ([Rb/Fe] =
modest, some stars are enriched with material produced by +0.39; Yong et al. 2008b). D’Orazi et al. (2013) measured
s-process nucleosynthesis. Among the bulge and inner disk [Rb/Fe] in three globular clusters (NGC 6752, NGC 1904
globularclusterswith[Fe/H]< 1.0,NGC6522exhibitsen- and47Tuc)andfoundconstantvaluesineachclusteratthe
−
hancements in the s-process∼elements (Barbuy et al. 2009; [Rb/Fe] 0.0level.The[Rb/Zr]ratiosare 0.04and 0.05
≃ − −
Chiappini et al. 2011) while HP-1(Barbuy et al. 2006) and forstarsc53andd10,respectively,whentheZrabundances
NGC 6558 (Barbuy et al. 2007) donot. are shifted on the Smith et al. (2000) scale as described in
TheabundancesofRbandPbofferimportantdiagnos- Yonget al.(2006).The[Rb/Zr]ratiosinM62liebelowthose
ticsregarding thenatureofthes-process.Duetoabranch- ofM4andM5,[Rb/Zr]=+0.17and+0.08,respectively,but
ing point at 85Kr along the s-process path, the Rb/Zr ra- abovetheωCenvalues(Smith et al.2000).Insummary,the
tio is sensitive to the neutron density at the site of the s- low [Rb/Fe] and [Rb/Zr] ratios in M62 offer no support for
process. At high neutron density, the [Rb/Zr] ratio is ex- aweak s-processcontribution(i.e.,thes-processinmassive
pectedtobeafactorof 10largerthanatlowneutronden- starsfromthe22Ne(α,n)25Mgreaction)fromwhichthehigh
∼
sity (Beer & Macklin 1989; Kappeler et al. 1989), although neutrondensitywouldresultinconsiderablyhigher[Rb/Fe]
theyieldsdependondetailsinthemodelsandnuclearreac- and[Rb/Zr]ratiosthanobserved.Wewillrevisitthe[Rb/Zr]
tionrates(Pignatari et al.2008,2010;van Raai et al.2012). ratiointhecontextoffast-rotatingmassivestarslaterinthis
Indeed, large Rb abundances have been observed in mas- subsection.
sive AGB stars (Garc´ıa-Hern´andez et al. 2006, 2009). Pb TheSrabundanceinM62,<[Sr/Fe]>= 0.39,appears
−
and Bi are the last stable nuclei along the s-process path. tobeunusuallylowcomparedtotheneighbouringelements,
If the total neutron exposure, or neutron supply per seed <[Rb/Fe]>=+0.11and<[Y/Fe]>=+0.16.Sinceouranal-
nuclei,issufficientlyhigh,large overabundancesof theseel- ysis only considered one line, Sri 4607.33˚A, we regard the
ements are predicted (e.g., Gallino et al. 1998; Busso et al. low Sr abundance with some caution. Indeed,other studies
1999, 2001; Goriely & Mowlavi 2000), and have been ob- thatconsiderthe4607.33˚ASrilineinglobularclustergiants
served (e.g., Van Eck et al. 2001; Ivanset al. 2005). find [Sr/Fe] ratios that are lower than for the neighbouring
As already noted, measurements of Rb and Pb were elements [Rb/Fe], [Y/Fe] and [Zr/Fe] (e.g., Roederer et al.
only possible for the two program stars observed at the 2011, Yong et al. 2014, in preparation). Bergemann et al.
Magellan Telescope.InFigure3weshowobservedandsyn- (2012)investigatednon-LTEcorrectionsforthe4607.33˚ASr
thetic spectra near the 7800˚A Rbi feature. We measure line,and for T /logg/[Fe/H]= 4400/2.2/ 1.2, thecorrec-
eff
−
10 D. Yong et al.
Table 4. Mean abundances and abundance dispersions for the pro-
gramstars(excludingabundance limits).
Species Z [X/Fe] σ[X/Fe]a [X/Fe] σ[X/Fe]a
Allstars ExcludingA124
Oi 8 0.52 0.35 0.49 0.38
Nai 11 0.38 0.35 0.46 0.31
Mgi 12 0.37 0.13 0.34 0.13
Ali 13 0.67 0.40 0.74 0.39
Sii 14 0.52 0.04 0.53 0.04
Cai 20 0.35 0.06 0.35 0.07
Scii 21 0.20 0.09 0.23 0.06
Tii 22 0.34 0.08 0.36 0.06
Tiii 22 0.42 0.09 0.43 0.10
Vi 23 0.22 0.15 0.27 0.11
Cri 24 0.09 0.17 0.12 0.15 Figure 10. Abundance ratios [Eu/X] for X = Rb to Pb. The
Mni 25 −0.30 0.12 −0.30 0.13 dashedredlinerepresents the scaled-solarr-processdistribution
Feib 26 −1.15 0.05 −1.15 0.05 (Bisterzoetal.2011).(ForSr,weapplyanon-LTEcorrectionof
Feiib 26 −1.16 0.05 −1.15 0.05 0.29dex,seeSection4.2fordetails.)
Coi 27 0.08 0.08 0.09 0.08
Nii 28 0.12 0.05 0.14 0.04
Cui 29 −0.22 0.18 −0.19 0.17 deviate from the r-process signature is Ba: while the solar
RZnbiic 3307 00..1111 00..1053 00..1101 00..1063 or-fpro0ce.2ssdveaxl.uIenisth[Eeusa/mBae]rw∼ay,0.a9lldeelexm, Men6t2s leixghhitbeirtsthaanraBtiao
Sri 38 −0.39 0.17 −0.33 0.08 ∼
depart from thesolar system r-process signature.
Yii 39 0.16 0.18 0.20 0.16
The neutron-capture element abundance ratios of the
Zri 40 0.41 0.20 0.48 0.09
Baii 56 0.32 0.11 0.36 0.03 M62 stars are not consistent with s-process nucleosynthesis
Laii 57 0.25 0.15 0.30 0.07 occurring in AGB stars of low metallicity. Low-mass AGB
Ceii 58 −0.01 0.13 0.04 0.07 stars of M < 3M⊙ are the main site of the s-process in
Prii 59 0.18 0.15 0.23 0.09 the Galaxy a∼nd typically produce [hs/ls]7 0.4, [Pb/hs] >
Ndii 60 0.23 0.13 0.27 0.11 1.0dexand[Rb/Zr]60at[Fe/H] 1.2(≈e.g.,Busso et al.
Smii 62 0.32 0.20 0.37 0.14 2001; Cristallo et al. 2011; Fishlock≈e−t al. 2014, in prepara-
Euii 63 0.52 0.14 0.58 0.04 tion). This is a direct consequence of the operation of the
Pbic 82 −0.20 ... −0.20 ... 13C(α,n)16Oreactionwhichisthedominantneutronsource
for M < 3M⊙ (e.g., Busso et al. 1999; Lugaro et al. 2012;
Bisterzo∼et al. 2012). For intermediate-mass AGB stars the
a Standarddeviation 22Ne(α,n)25Mg reaction is the neutron source and operates
b [Fei/H]or[Feii/H] during convective thermal pulses. The abundance pattern
c Measurementsonlyforstarsc53andd10.
produced in those stars should be dominated by lighter
s-process elements (Y, Rb, Sr, Zr) (van Raai et al. 2012;
Karakas et al.2012).At[Fe/H] 1.2,typicalabundances
tion is 0.29 dex. The magnitude and sign of this correction ≈−
in stars M >3M⊙ are [hs/ls] 0.8, [Pb/hs] 0.2, and
wouldhelptoreconciletheSrabundanceswith thoseofRb ≈− ≈−
[Rb/Zr] >0∼.0 dex (Fishlock et al. 2014, in preparation).
andY.InFigure9weadjusttheSrabundanceaccordingly.
Fast∼-rotatingmassivestars(whichwehereafterreferto
In Figure 4, we plot observed and synthetic spectra
interchangeablyas“spinstars”)canproducealargeamount
near the 4058˚A Pbi feature. We measure values of [Pb/Fe]
of s-process products at low metallicity, where the domi-
<−0.09and[Pb/Fe]=−0.20forstarsc53andd10,respec- nantneutronsourceis22Ne(α,n)25MgintheconvectiveHe-
tively. For comparison, the average values in M4 and M5
burning core and in the subsequent convective C-burning
are [Pb/Fe] = +0.30 and 0.35, respectively (Yonget al.
− shell (Pignatari et al. 2008; Frischknechtet al. 2012). In
2008b). Thus, the M62 values for [Pb/Fe] lie closer to M5 thesestars, 22Neis primary,i.e., it is not directly produced
thantoM4(seeFigure9).Roedereret al.(2010)arguethat
fromtheinitialCNOabundancesasinnon-rotatingmassive
the Pb/Eu ratio offers a good diagnostic of the presence of
stars(seeforinstancetheweaks-processinmassivestarsat
s-process enriched material. For M62, the Pb to Eu ratios
solarmetallicity,e.g.,The et al.2007;Pignatari et al.2010).
are at the same level as the (more metal-poor) “r-process The faster the star rotates, the more primary 22Ne can be
standard stars” in Roederer et al. (2010).
madeavailableintheHecore.Themostabundantelements
From Table 4, all element ratios [Eu/X] (for X = La
from thes-process in spin stars areSr,Y and Zrwhile pro-
andheavierelements) areconsistent, within 0.2 dex,with
∼ duction for heavier elements up to Ba strongly depends on
thescaled-solarr-processdistribution(Bisterzo et al.2011),
therotationspeedandonnuclearuncertainties,e.g.,onthe
see Figure10. (Notethat while werefer to s-processand r- α-capturerateson22Ne(Pignatari et al.2008)andontheα-
process elements according to their origin in solar system
material,thosedefinitionsmaybeincorrectinM62,e.g.,La
andCeinthisclusterarelikelyproducedprimarilythrough 7 hsreferstotheheavy s-processelements (Ba, La,etc.)andls
the r-process.) Regarding the lighter elements, the first to referstothelights-processelements(Y,Zr,etc.).