Comprehensive abundance analysis of red giants in the open clusters NGC 752, NGC 1817, NGC 2360 and NGC 2506
aa r X i v : . [ a s t r o - ph . GA ] S e p Mon. Not. R. Astron. Soc. , 1– ?? (2011) Printed 2 September 2018 (MN L A TEX style file v2.2)
Comprehensive abundance analysis of red giants in theopen clusters NGC 752, NGC 1817, NGC 2360 and NGC2506 ⋆ Arumalla B.S. Reddy ⋆ , Sunetra Giridhar and David L. Lambert Indian Institute of Astrophysics, Bangalore 560034, India W.J. McDonald Observatory, The University of Texas at Austin, Austin, TX 78712 - 0259, USA
ABSTRACT
We have analyzed high-dispersion echelle spectra ( R & , − ± − ± − ± − ± Key words: – Galaxy: abundances – Galaxy: open clusters and associations – stars:abundances: general – open clusters: individual: NGC 752, NGC 1817, NGC 2360,NGC 2506
Open clusters (OCs) are believed to be coeval groups of starsborn from the same proto-cluster cloud which may have beenpart of a larger star-forming region in the Milky Way. Ages ofclusters range from very young where stars are still formingto nearly 10 Gyr (Dias et al. 2002). Since all stars in mostOCs are at the same distance and have the same chemicalcomposition, basic stellar parameters like age, distance, andmetallicity can be determined more accurately than for fieldstars. Thus, OCs provide an excellent opportunity to mapthe structure, kinematics, and chemistry of the Galactic diskwith respect to Galactic coordinates and time.The presence of chemical homogeneity among clustermembers has been shown by the study of OCs, see, for ex-ample, spectroscopic analyses of the Hyades (Paulson et al.2003; De Silva et al. 2006) and Collinder 261 (Carretta etal. 2005; De Silva et al. 2007). This observed homogeneitysignifies that the proto-cloud is well mixed, and hence, theabundance pattern of a cluster bears the signature of chem-ical evolution of the natal cloud. Chemical evolution of theMilky Way has, of course, been well studied using field stars.A large fraction of field stars are from disrupted OCs (Lada ⋆ E-mail: [email protected] (ABSR); [email protected] (SG);[email protected] (DLL) & Lada 2003). The youngest OCs may be intact. The old-est OCs may be totally disrupted. Thus, the field stars donot fully sample the age distribution of OCs and, in partic-ular, the youngest stellar generations are under-representedby field stars.In this paper, we report abundance analyses from high-resolution spectra of red giants in four OCs: NGC 752, NGC1817, NGC 2360, and NGC 2506. These analyses are the firstfor these OCs to report elemental abundances for elementsfrom Na to Eu.The structure of the paper is as follows: In Section 2we describe the data selection, observations and data reduc-tion and Section 3 is devoted to the abundance analysis. InSection 4 we present our results and compare them with theabundances derived from samples of field thin and thick diskstars (i.e. dwarfs and giants). Finally, in Section 5 we givethe conclusions.
Clusters were selected from the
New catalogue of opticallyvisible open clusters and candidates (Dias et al. 2002). Em-phasis was placed on OCs not yet subjected to high reso- (cid:13) A. B. S. Reddy, S. Giridhar and D. L. Lambert
Table 1.
Target clusters and their properties from the literature.Cluster l b Age [Fe/H] phot . R gc (m-M) v E(B-V) [Fe/H] ref (deg.) (deg.) (Gyr) (dex.) (kpc) (mag.) (mag.)NGC 752 137.12 − − − − − − − Table 2.
The observed stars.Cluster Star ID α (2000 . δ (2000 .
0) V B-V RV helio S/N(hh mm s) ( ◦ ′′ ′ ) (mag.) (mag.) (km s − ) at 6000 ˚ANGC 752 77 01 56 21.60 37 36 08.00 9.35 +1.02 +6.3 ± ± ± ± ± ± ± −
15 37 30.49 10.74 +1.04 +30.4 ± −
15 37 32.62 11.03 +1.04 +29.1 ± −
15 34 13.30 11.09 +1.01 +29.2 ± −
15 31 39.80 10.34 +1.16 +29.5 ± −
10 46 37.50 11.95 +1.07 +84.1 ± −
10 48 22.73 13.12 +0.98 +84.9 ± −
10 45 38.50 13.30 +0.91 +84.9 ± lution spectroscopy. Since the main sequence stars in thechosen OCs were faint, we elected to observe the red giantmembers. For each of the target clusters, red giants wereselected using the WEBDA database. Target clusters andtheir properties are shown in the Table 1: column 1 repre-sents the cluster name, columns 2 & 3 the Galactic longi-tude and latitude in degrees, column 4 the age, column 5 thephotometric estimate of the iron abundance, column 6 theGalactocentric distance, column 7 the distance modulus, col-umn 8 the reddening, column 9 the reference to [Fe/H]. Allquantities are from the database entry except for the [Fe/H]abundance and the Galactocentric distance, Rgc, which wecalculate assuming a distance of the Sun from the Galacticcentre of 8.0 ± & IRAF soft-ware of
NOAO within the imred and echelle packages, involv- http://iraf.noao.edu/ Figure 1.
Representative spectra for the four clusters.ing bias subtraction, scattered light correction, flat-fielding, wave-length calibration and continuum fitting. We measured the radialvelocity (RV) of each star on each spectrum. The continuum-fittedc (cid:13) , 1– ?? GC 752, NGC 1817, NGC 2360, NGC 2506 spectrum was corrected for the Doppler shift using the routine dopcor available in IRAF task splot . Our radial velocity mea-surements are in fair agreement with the previous radial velocitymeasurements for the red giants in OCs (Mermilliod et al. 2008).The identification and basic observational data for the starsobserved in each of the clusters are given in Table 2, along with thecomputed radial velocity and S/N of each of the spectra extractedat 6000 ˚A for each of the stars. Spectra of a representative regionare shown in Figure 1 for one star from each of the four clusters.
Line selection
Selection of stellar lines which are free from blends is crucial forderiving accurate elemental abundances. We used Rowland’s pre-liminary table of solar spectrum wavelengths (Moore et al. 1966)and the the Arcturus spectrum (Hinkle et al. 2000) to identifyunblended spectral lines. We employed strict criteria in the se-lection of suitable lines. First, in order to avoid the difficulty indefining the continuum due to heavy line crowding in the bluepart of the spectrum, we selected lines only within the 4300 to7850 ˚A wavelength range. Second, regions containing telluric ab-sorption lines were generally avoided. Third, lines which appearasymmetric were assumed to be blended with unidentified linesand discarded. Fourth, lines with equivalent widths (EWs) below10 m˚A were rejected because they are too sensitive to noise andthe normalization of the continuum, and lines with equivalentwidths greater than 230 m˚A were discarded because they are toosaturated.Application of our criteria resulted in a list of 55 Fe I lineswith lower excitation potentials (LEP) ranging from 0.9 to 5.0eV and EWs of up to 180 m˚Aand 10 Fe II lines with excitationpotentials of 2.8 to 3.9 eV and EWs up to 120 m˚A.A portion of final linelist with solar EWs is given in Table3 with details about the lines including the log gf’s (see below).For each line in the constructed linelist, we provide a gf-valuefrom the literature. In most cases, we located recent experimentaldeterminations or chose values from critical reviews. Referencesto the adopted sources are given in Table 3. EWs were measuredmanually using the task ’ splot ’ contained in
IRAF by fitting aGaussian profile to the observed line.For several absorption features, multiple components of agiven atomic transition contribute to the feature. In such cases,we computed a synthetic spectrum including all components andoccasionally other lines too and matched the synthetic spectrumto the observed spectrum by varying the abundance of the elementin question.As a check on the chosen gf-values, we derived solar abun-dances using the ATLAS9 model for T eff = 5777 K, log g =4.44cgs. A microturbulence of ξ t = 0.95 km s − was found from ironlines. Solar equivalent widths were measured off the solar inte-grated disk spectrum (Kurucz et al. 1984). Abundances are givenin Table 4 along with those from the recent review by Asplundet al. (2009). Our abundances for the majority of elements are ingood agreement with Asplund et al.’s. Small differences in abun-dances are inevitable, for example, the line lists are not necessarilyidentical as to selected lines and/or gf-values and the adopted so-lar models are different; ours is a classical model but Asplund etal. for many elements use a model representing the solar granula-tion. For the purposes of determining the stellar abundances, weadopt our solar abundances when computing [X/H] and [X/Fe],i.e., our analysis is essentially a differential one relative to theSun. Table 3.
Adopted linelist. All the columns are self-explanatory.Atom Wavelength LEP log gf W λ, ⊙ Ref a .( ˚A) (eV) (m˚A)Na I 4668.567 2.100 -1.31 53.4 NIST4982.821 2.100 -0.96 79.0 NIST5688.210 2.100 -0.45 115.4 NIST6154.224 2.100 -1.55 36.3 NIST6160.746 2.100 -1.25 55.9 NISTMg I 5711.090 4.340 -1.72 104.5 NIST6318.708 5.108 -1.90 42.9 NISTAl I 7084.564 4.020 -0.93 21.4 KUR7835.296 4.020 -0.65 41.2 KUR7836.119 4.020 -0.47 55.4 KURSi I 5665.551 4.920 -2.04 39.4 NIST5645.606 4.930 -2.14 34.9 LUCK5701.100 4.930 -2.05 37.7 NIST5753.632 5.610 -1.30 42.5 LUCK6131.570 5.610 -1.70 22.3 LUCK6131.851 5.610 -1.62 23.3 LUCK6145.013 5.610 -1.48 36.5 LUCK6237.319 5.610 -1.14 56.7 LUCK6244.470 5.610 -1.36 45.1 LUCK6243.812 5.613 -1.26 46.5 KUR6142.486 5.620 -1.54 33.0 LUCK6721.840 5.862 -1.06 42.8 LUCK6195.445 5.873 -1.80 14.9 LUCKCa I 6122.221 1.890 -0.32 161.9 LUCK5581.971 2.520 -0.55 93.7 LUCK5590.117 2.520 -0.57 91.7 LUCK6166.434 2.520 -1.14 69.2 LUCK6169.560 2.520 -0.48 108.5 LUCK6455.599 2.520 -1.29 56.2 LUCK6499.649 2.520 -0.82 85.0 LUCK6471.662 2.526 -0.68 90.7 LUCKSc I 5686.832 1.440 0.38 8.3 LUCK5356.090 1.860 0.17 2.3 LUCKSc II 6604.587 1.357 -1.31 35.1 NIST5667.141 1.500 -1.20 30.0 KUR6245.615 1.507 -1.03 35.4 KUR6300.681 1.507 -1.89 8.1 KUR5526.813 1.768 0.06 75.6 KURTi I 5039.960 0.021 -1.13 75.7 NIST5460.497 0.048 -2.75 9.6 LUCK4999.510 0.826 0.31 103.6 LUCK5020.026 0.836 -0.35 72.6 KUR5295.776 1.067 -1.63 13.1 NIST5474.223 1.460 -1.23 10.8 NIST5490.148 1.460 -0.93 21.6 NIST4617.274 1.749 0.39 61.2 NIST5739.980 2.236 -0.60 7.3 NIST5702.656 2.292 -0.57 8.1 NISTTi II 4764.528 1.237 -2.77 37.2 LUCK4708.665 1.240 -2.21 52.9 NIST5005.168 1.566 -2.54 25.5 NIST5381.022 1.566 -1.85 57.3 LUCK5396.244 1.580 -2.92 12.1 LUCK5336.788 1.582 -1.70 69.2 NIST5418.767 1.582 -1.99 48.1 KURV I 6251.823 0.286 -1.34 15.8 NIST6111.647 1.043 -0.71 10.7 NIST5727.653 1.051 -0.87 8.9 NIST6135.366 1.051 -0.75 10.4 NIST5737.062 1.064 -0.74 10.8 NIST5668.365 1.081 -1.03 5.6 NIST5670.848 1.081 -0.42 19.0 NIST5727.044 1.081 -0.01 39.9 NISTc (cid:13) , 1– ?? A. B. S. Reddy, S. Giridhar and D. L. Lambert
Table 3 − continuedAtom Wavelength LEP log gf W λ, ⊙ Ref a .(˚A) (eV) (m˚A)Cr I 4545.958 0.941 -1.37 81.1 SLS5296.696 0.983 -1.36 91.0 SLS5300.747 0.983 -2.00 58.3 SLS5345.802 1.004 -0.95 112.2 SLS5238.959 2.709 -1.30 16.1 SLS5329.139 2.910 -0.06 65.8 KUR5784.967 3.321 -0.38 31.2 NIST5214.129 3.369 -0.74 17.1 NIST5628.640 3.422 -0.74 14.4 SLS5287.174 3.440 -0.87 10.7 SLS5312.853 3.449 -0.55 19.8 SLS5304.178 3.463 -0.67 15.4 SLSCr II 5279.874 4.070 -2.10 19.0 NIST5308.424 4.071 -1.81 25.3 NIST5237.321 4.073 -1.16 52.6 NIST5334.864 4.073 -1.56 33.6 KUR5313.578 4.074 -1.65 33.4 NIST5502.081 4.170 -1.99 18.3 NISTMn I 6013.488 3.072 -0.25 84.9 NIST6021.796 3.075 0.03 90.7 NIST5377.608 3.845 -0.11 48.0 KUR5399.480 3.850 -0.29 37.4 KURFe I 6136.995 2.198 -2.95 65.6 F&W6252.562 2.404 -1.69 120.9 F&W5141.742 2.424 -2.24 84.5 F&W5701.548 2.559 -2.22 83.8 F&W6646.931 2.608 -3.95 9.3 KUR5036.918 3.018 -3.04 24.5 F&W5215.184 3.266 -0.87 119.7 F&W5576.093 3.431 -0.94 106.2 F&W5568.863 3.635 -2.95 10.3 LUCK5054.655 3.640 -1.92 39.8 F&W5636.695 3.640 -2.56 20.7 F&W5760.343 3.642 -2.44 22.9 F&W5539.278 3.643 -2.61 18.7 F&W6411.653 3.654 -0.72 116.9 F&W5466.986 3.655 -2.23 32.5 F&W6336.828 3.687 -0.86 101.9 F&W5379.574 3.695 -1.51 60.5 F&W6003.014 3.882 -1.15 81.6 NIST6187.988 3.943 -1.67 46.0 F&W5293.957 4.143 -1.84 28.7 F&W6165.358 4.143 -1.47 43.9 F&W5608.973 4.209 -2.40 10.4 LUCK5618.631 4.209 -1.28 49.6 F&W5074.753 4.221 -0.23 114.3 F&W5738.230 4.221 -2.34 11.8 LUCK5579.338 4.231 -2.40 10.6 LUCK5016.477 4.256 -1.69 32.7 LUCK5090.783 4.256 -0.44 90.1 NIST5243.777 4.256 -1.12 59.9 F&W5646.682 4.261 -2.50 7.1 LUCK5717.832 4.284 -1.10 60.5 F&W5197.934 4.301 -1.62 35.5 F&W5466.398 4.371 -0.63 77.0 LUCK5295.312 4.415 -1.67 28.5 F&W5560.210 4.435 -1.16 50.5 F&W5577.022 5.033 -1.55 12.3 LUCKFe II 5000.730 2.780 -4.61 11.9 M&B4520.225 2.807 -2.65 81.3 M&B4993.353 2.807 -3.62 39.0 M&B6369.460 2.891 -4.11 19.5 M&B6432.680 2.891 -3.57 40.9 M&B Table 3 − continuedAtom Wavelength LEP log gf W λ, ⊙ Ref a .(˚A) (eV) (m˚A)5256.931 2.892 -4.06 21.1 M&B5425.260 3.199 -3.22 42.6 M&B6084.103 3.199 -3.88 20.2 F&W5234.620 3.221 -2.18 84.8 M&B5414.067 3.221 -3.58 27.7 M&B6149.244 3.889 -2.84 35.8 F&W6247.559 3.892 -2.43 52.1 F&W6456.383 3.903 -2.19 61.0 F&WCo I 6116.994 1.785 -2.49 6.0 NIST5647.232 2.280 -1.56 13.8 NIST5212.680 3.514 -0.14 19.2 KUR6454.990 3.632 -0.25 14.6 NIST5342.701 4.022 0.69 30.8 KURNi I 5578.720 1.676 -2.64 55.3 NIST5748.352 1.676 -3.26 27.6 NIST6191.181 1.677 -2.35 71.0 KUR6177.242 1.826 -3.51 14.4 NIST4998.229 3.606 -0.78 52.9 NIST5462.493 3.847 -0.93 38.6 NIST5468.103 3.847 -1.61 12.4 NIST5589.357 3.898 -1.14 27.6 NIST5593.736 3.898 -0.84 40.2 NIST5638.745 3.898 -1.73 9.5 NIST6111.072 4.088 -0.87 32.7 NIST5625.316 4.089 -0.70 37.7 NIST5682.199 4.105 -0.47 50.4 NIST5760.830 4.105 -0.80 33.9 NISTCu I 5218.201 3.820 0.26 49.3 NISTZn I 4722.160 4.030 -0.34 65.0 KUR6362.342 5.790 0.27 28.1 LUCKY II 5200.409 0.992 -0.57 36.3 HLG4982.133 1.033 -1.29 13.3 HLG5289.817 1.033 -1.85 4.0 HLG4883.688 1.080 0.07 54.3 HLG5402.773 1.840 -0.63 12.3 LUCKZr I 6143.201 0.070 -1.10 2.1 BGHL4739.483 0.650 0.23 6.2 BGHLBa II 5853.678 0.604 -1.02 62.5 LUCK6496.905 0.604 -0.37 97.8 LUCKLa II 5303.538 0.321 -1.35 4.4 LBS6390.486 0.321 -1.41 3.0 LBSCe II 5472.281 1.240 -0.18 2.1 LUCKNd II 5092.800 0.380 -0.61 6.5 DLS5319.813 0.550 -0.14 11.4 DLS5485.539 1.260 -0.12 4.1 DLSSm II 4519.630 0.544 -0.35 6.1 LDSEu II 6645.108 1.379 0.12 4.80 LWD The lines are arranged in the order of their increasing LEP. a References for the adopted gf-values:BGHL − Bi´emont et al. (1981);F&W − Fuhr & Wiese (2006);HLG − Hannaford et al. (1982);KUR − Kurucz (1998);LWD − Lawler et al. (2001);LBS − Lawler et al. (2001);LDS − Lawler et al. (2006);LSC − Lawler et al. (2009);LUCK − Luck (Private communication);M&B − Mel´endez & Barbuy (2009);SLS − Sobeck et al. (2007);NIST − Atomic Spectra Database aa http://physics.nist.gov/PhysRefData/ASD/lines form.htmlc (cid:13) , 1– ?? GC 752, NGC 1817, NGC 2360, NGC 2506 Table 4.
Solar abundances derived by employing the solar modelatmosphere from Castelli & Kurucz (2003) compared with thephotospheric abundances from Asplund et al. (2009).Species log ǫ ⊙ log ǫ ⊙ (our study) (Asplund)Na I 6.29 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± Initial estimates of the effective temperature of a red giantwere derived from dereddened B and V photometry using theempirically-calibrated colour-temperature relation by Alonso etal. (1999) based on a large sample of field and globular clustergiants of spectral types from F0 to K5. An error of 3% is expectedin the derived temperatures.Gravities were computed using the known distance to theOCs, temperature and bolometric corrections, and the clusterturn-off mass. We have adopted a turn-off mass of 1.5M ⊙ forNGC 752 (Bartasiute et al. 2007), 2M ⊙ for NGC 1817 (Jacobsonet al. 2009), 1.98M ⊙ for NGC 2360 (Hamdani et al. 2000),and 1.69M ⊙ for NGC 2506 (Carretta et al. 2004). The relationbetween log g and T eff is given by (Allende Prieto et al. 1999)log g ⋆ = log g ⊙ +log( M ⋆ /M ⊙ )+4 log( T eff /T eff ⊙ )+ 0 . V + BC V ) + 2 logπ + 0 .
12 (1)with the corresponding luminosity given by log ( L ∗ /L ⊙ ) = − [0 . V + BC V ) + 2 log π + 0 .
12] (2)where π is the parallax, V is the apparent Johnson V magnitudecorrected for reddening, and BC V is the bolometric correction.We adopt log g ⊙ = 4.44, T eff , ⊙ = 5777 K. Bolometric correctionswere taken from the calibration by Alonso et al. (1999).We suppose that the errors in different quantities involved in equation (1) are independent of each other. Then, by assuming anerror of 10% in the stellar mass, an uncertainty of 3% in T eff , anuncertainty of 5% in photometric V magnitude and the bolometriccorrections, and an error of 10% in parallax, we get an error of ≃ g and the corresponding uncertainty in log L ∗ amounts to 0.08. The spectroscopic abundance analysis was performed with the2010 version of the local thermodynamical equilibrium (LTE)line synthesis program MOOG (Sneden 1973) . Model atmo-spheres were generated by linear interpolation from the ATLAS9model atmosphere grid (Castelli & Kurucz 2003) . A model at-mosphere is characterized by the effective temperature T eff , thesurface gravity log g , microturbulence velocity ξ t , and composi-tion. These models use the classical assumptions of line-blanketedplane-parallel uniform atmospheres in LTE and hydrostatic equi-librium with flux conservation.Spectroscopically, we determine the stellar parameters in theconventional way when LTE is the paramount assumption. Thekey lines are those of Fe I and Fe II for which we take gf-valuesfrom Fuhr & Wiese (2006) and Mel´endez & Barbuy (2009). Themicroturbulence assumed to be isotropic and depth independentis determined from Fe II lines by the requirement that the abun-dance be independent of a line’s EW. A model atmosphere withthe photometrically determined parameters was used initially forthis determination. The effective temperature is found by impos-ing the requirement that the Fe abundance from Fe I lines beindependent of a line’s lower excitation potential. Finally, thesurface gravity is found by requiring that Fe I and Fe II lines givethe same Fe abundance for the derived effective temperature andmicroturbulence.A check on the microturbulence is provided by lines of speciesother than Fe I. For example, for the star NGC 752 ξ t , from 0 to 6 kms − . It is clear that the minimum value of dispersion for all speciesis in the range 1.2-1.6 km s − . Thus, we adopt a microturbulenceof 1.45 km s − with an uncertainty of 0.20 km s − .Several elements other than Fe provide both neutral andionized lines and so offer a check on the condition of ionizationequilibrium of Fe. Consider for example the four giants from NGC752: the abundance differences [X/H] between neutral and ionizedlines of Sc, Ti, V and Cr are on average 0.03, -0.03, -0.01, and-0.05 dex, respectively where ± logg by ∓ of 0.15.The uncertainties in the derived surface temperatures fromspectroscopy are provided by the errors in the slope of the relationbetween the Fe I abundance and LEP of the lines. A perceptiblechange of slope occures for variations of the temperature from50 −
100 K about the adopted model.Therefore, the typical errors considered in this analysis are100 K in T eff , 0.25 cm s − in log g and 0.20 km s − in ξ t .The derived stellar parameters for program stars in each ofthe cluster are given in Table 5: column 1 represents the clustername, column 2 the star ID, columns 3 & 4 the photometric T eff and log g values, columns 5-7 the spectroscopic T eff , log g and ξ t estimates. Finally, the spectroscopic and photometric luminosities( log(L/L ⊙ )) are presented in columns 8 & 9. Photometric andspectroscopic estimates are in excellent agreement. Mean differ-ences in T eff , log g and logL/L ⊙ across the 14 stars are − ± http://kurucz.harvard.edu/grids.htmlc (cid:13) , 1– ?? A. B. S. Reddy, S. Giridhar and D. L. Lambert
Figure 2.
The standard deviation for various species about the mean abundances as a function of microturbulence for the star NGC752 . ± .
10 cgs, and − . ± .
10, respectively. The correspond-ing comparison of the spectroscopic with the photometric [Fe/H]in Table 1 also illustrates fair agreement: ∆[Fe/H] = 0.09 (NGC752), 0.21 (NGC 1817), 0.05 (NGC 2360), and 0.12 (NGC 2506).
We compared the stellar spectra to the synthetic spectra to deriveabundances for the lines having intrinsic multiple components andlines affected by blends. Figures 3 and 4 show synthetic spectra fitto the observed one using three different abundances. The dottedline is the stellar spectrum. The red line is the best fit to thestellar spectrum, with the other lines representing different valuesfor [Ba/H] and [Eu/H] abundances, based on χ goodness of fitprovided by MOOG.In this analysis, we have adopted the hfs data of Prochaska& McWilliam (2000) for the synthesis of Mn I line at 6013 ˚A,Allen et al. (2011) for Cu I line at 5218 ˚A, McWilliam (1998)for Ba II line at 5853 ˚A and Mucciarelli et al. (2008) for Eu IIline 6645 ˚A. Isotopic ratios for Cu I, Ba II and Eu II were takenfrom Lodders (2003). Further, we have synthesized the lines CeII line at 5472 ˚A and Sm II line at 4577 ˚A since the blends makeit impossible to measure their EWs. The spectrum synthesis wascarried out by running the MOOG in ’ synth ’ mode.Since all odd species exhibit hfs effects of relatively varyingstrengths, we have performed spectrum synthesis over Sc II lineat 6245 ˚A V I line at 5727 ˚A and Co I line at 5647 ˚A . Here, wehave adopted the hfs data of Prochaska & McWilliam (2000) forSc II and for V I and Co I hfs components were taken from Kurucz linelists . We noticed that the these lines are not severely effectedby hfs effects, causing an abundance difference of 0.0 − The abundance analysis was conducted with the model atmo-spheres having the stellar parameters determined from the spec-tra (Table 2), the line list (Table 12) and the program MOOG.Abundances [X/H] are expressed relative to the solar abundancesderived from the adopted gf-values. Results for the individualstars in each of the OCs are given in Tables 8, 9, 10, and 11.For each abundance based on analysis of EWs, the abun-dance and the standard deviation were calculated from all linesof a given species. The tables give the abundances of [Fe/H] and[X/Fe] for all elements. The quantity [X/Fe] minimizes the sen-sitivity to errors in the model atmosphere arising from uncer-tainties affecting the stellar parameters. Inspection of the Tables8-11 shows that, in general, the compositions [X/Fe] of stars in agiven cluster are generally identical to within the (similar) stan-dard deviations computed for an individual star. Exceptions tendto occur for species represented by few lines, as expected whenthe uncertainty in measuring equivalent widths is a contributor tothe total uncertainty. From the spread in the abundances for thestars of a given cluster we obtain the standard deviation σ in theTables 8-11 in the column headed ‘average’. Errors in the adoptedgf values are unimportant when providing differential abundances http://kurucz.harvard.edu/linelists.htmlc (cid:13) , 1– ?? GC 752, NGC 1817, NGC 2360, NGC 2506 Figure 3.
Synthetic spectra and the observed spectrum of NGC 2360
Figure 4.
Synthetic spectra and the observed spectrum of NGC 752 − and 0.2 km s − withrespect to the chosen model atmosphere are summarized in Table6. We quadratically added the three contributors, by taking thesquare root of the sum of the square of individual errors associatedwith uncertainties in temperature, gravity and microturbulence,to obtain σ . The total error σ tot for each of the element is thequadratic sum of σ and σ . The error bars in the abundancetables correspond to this total error. The four clusters support the widely held impression that thereis an abundance gradient such that the metallicity [Fe/H] at thesolar galactocentric distance decreases outwards (Magrini et al.2009) at about -0.1 dex per kpc: (R gc , [Fe/H]) = (8.3, − .
03) forNGC 752, (9.3, − .
07) for NGC 2360, (9.9, − .
12) for NGC 1817,and (10.5, − .
20) for NGC 2506.Results – Tables 8–11 – for the individual clusters are con-sistent with the assumption that stars within a given cluster havethe same composition.If OCs are the principal supplier of field stars, there shouldbe a very close correspondence between the composition of starsin clusters and the field. Such a correspondence represents a stiffchallenge to the idea that field stars have come from clustersbecause modern studies of field stars show that there is no dis-c (cid:13) , 1–, 1–
20) for NGC 2506.Results – Tables 8–11 – for the individual clusters are con-sistent with the assumption that stars within a given cluster havethe same composition.If OCs are the principal supplier of field stars, there shouldbe a very close correspondence between the composition of starsin clusters and the field. Such a correspondence represents a stiffchallenge to the idea that field stars have come from clustersbecause modern studies of field stars show that there is no dis-c (cid:13) , 1–, 1– ?? A. B. S. Reddy, S. Giridhar and D. L. Lambert
Table 5.
Basic photometric and spectroscopic atmospheric parameters for the stars in each cluster.Cluster Star ID T eff,phot log g phot T eff,spec log g spec ξ spec log( L/L ⊙ ) log( L/L ⊙ )(K) (cm s − ) (K) (cm s − ) (km sec − ) spectroscopy photometryNGC 752 77 4780 2.75 4850 2.65 1.26 1.71 1.54137 4780 2.57 4850 2.50 1.36 1.81 1.72295 4899 2.80 5050 2.85 1.47 1.53 1.53311 4761 2.62 4850 2.60 1.45 1.71 1.66NGC 1817 1027 5177 2.66 5100 2.60 1.39 1.92 1.892038 4968 2.57 5100 2.45 1.44 2.07 1.902059 5059 2.57 4800 2.40 1.38 2.01 1.94NGC 2360 5 4899 2.56 4900 2.70 1.29 1.75 1.896 4899 2.68 5000 2.50 1.34 1.98 1.778 4962 2.74 5050 2.60 1.37 1.90 1.7412 4668 2.27 4650 2.10 1.23 2.26 2.10NGC 2506 2212 4710 1.86 4700 1.75 1.21 2.56 2.453231 4893 2.44 5000 2.50 1.42 1.92 1.944138 5048 2.60 5100 2.60 1.47 1.85 1.84 Table 6.
Sensitivity of abundances to the uncertainties in the modelparameters for the star 5 in NGC 2360 with T eff = 4900 K, log g =2.70 cm s − ,and ξ t = 1.29 km s − . T eff ±
100 K log g ± ξ t ± σ T eff σ logg σ ξ t σ Na I +0 . / − . − . / + 0 . − . / + 0 .
06 0.05Mg I +0 . / − . − . / + 0 . − . / + 0 .
05 0.05Al I +0 . / − . − . / . − . / + 0 .
03 0.03Si I − . / + 0 .
04 +0 . / − . − . / + 0 .
03 0.04Ca I +0 . / − . − . / . − . / + 0 .
11 0.08Sc I +0 . / − .
15 +0 . / − . − . / + 0 .
03 0.08Sc II − . / + 0 .
04 +0 . / − . − . / + 0 .
07 0.08Ti I +0 . / − .
12 +0 . / + 0 . − . / + 0 .
07 0.08Ti II − . / + 0 .
05 +0 . / − . − . / + 0 .
15 0.11V I +0 . / − .
14 +0 . / . − . / + 0 .
07 0.08Cr I +0 . / − .
09 0 . / . − . / + 0 .
08 0.07Cr II − . / + 0 .
09 +0 . / − . − . / + 0 .
06 0.09Mn I − . / + 0 .
09 +0 . / − . − . / + 0 .
13 0.09Fe I +0 . / − .
04 +0 . / − . − . / + 0 .
10 0.06Fe II − . / + 0 .
12 +0 . / − . − . / + 0 .
10 0.12Co I +0 . / − .
03 +0 . / − . − . / + 0 .
06 0.05Ni I +0 . / − .
01 +0 . / − . − . / + 0 .
09 0.06Cu I +0 . / − .
01 +0 . / − . − . / + 0 .
06 0.04Zn I − . / + 0 .
07 +0 . / − . − . / + 0 .
15 0.11Y II − . / + 0 .
01 +0 . / − . − . / + 0 .
08 0.08Zr I +0 . / − .
19 0 . / + 0 . − . / + 0 .
01 0.10Ba II − . / + 0 .
03 +0 . / − . − . / + 0 .
23 0.15La II +0 . / .
00 +0 . / − . − . / + 0 .
05 0.07Ce II +0 . / .
00 +0 . / − . − . / + 0 .
02 0.07Nd II 0 . / .
00 +0 . / − . − . / + 0 .
04 0.07Sm II +0 . / .
00 +0 . / − . − . / + 0 .
07 0.07Eu II − . / + 0 .
02 +0 . / − . − . / + 0 .
02 0.07cernible ‘cosmic’ dispersion in relative abundances – [X/Fe] – ata given [Fe/H] (Reddy et al. 2006). The four OCs are very likelyrepresentatives of the Galactic thin disk but at their metallic-ity thin and thick disk stars very likely have the same relativeabundances.Several studies of thin disk dwarfs and giants have been re-ported recently. For almost all elements over the [Fe/H] rangesampled by these four OCs, the field dwarfs and giants show asolar-like mix of elements, i.e., [X/Fe] ≃
0, with very little star-to-star scatter at a given [Fe/H]. Sample papers echoing this as- sertion include Edvardsson et al. (1993), Bensby et al. (2005),Reddy et al. (2003, 2006), Luck & Heiter (2006) for dwarfs, andMishenina et al. (2006), and Luck & Heiter (2007), and Takeda etal. (2008) for giants. These papers invoke, as we have done, clas-sical methods of abundance analysis involving standard modelatmospheres and LTE line formation.Methods of abundance analysis including choices of gf-values, selection of model atmosphere grid and determination ofsolar reference abundances differ among these papers. Yet, theresults suggest that differences of ± .
05 and possibly ± .
10 dexmay arise among similar analyses by different authors of the sameor similar stars. Such differences are attributable to measurementerrors with the cosmic dispersion masked by such errors. One ex-pects applications of the classical method to give slightly differentresults for dwarfs and giants for several reasons, e.g., the effectsof departures from LTE will be different for giants and dwarfs,and the ability of standard atmospheres to represent true stellaratmospheres may differ for dwarfs and giants. Thus, we restrictcomparisons between our results and those by similar methodsfor field giants i.e. systematic errors will be very similar acrossthis comparison.A useful comparison of abundances between our OCs andfield giants may be made using Luck & Heiter’s (2007) large sam-ple of field giants analysed by methods similar to ours, i.e., a dif-ferential analysis with respect to the Sun. Using their Table 4, wecalculated the mean abundances in field giants across the [Fe/H]range of our clusters (0.0 to − .
2) and those values are presentedin column 6 of Table 7. Our cluster abundances in Table 7 matchthe abundances of the field giants to within about ± ± (cid:13) , 1– ?? GC 752, NGC 1817, NGC 2360, NGC 2506 Table 7.
Elemental abundances for the four clusters in this study and the thin disk meanvalues from Luck & Heiter (2007) in the metallicity range of our clusters (i.e. 0.0 to − . ± .
03 +0 . ± .
02 +0 . ± .
03 +0 . ± .
03 0 . ± . − . ± .
03 +0 . ± .
03 +0 . ± .
03 +0 . ± .
04 0 . ± . . ± .
02 +0 . ± .
02 +0 . ± .
02 +0 . ± .
01 0 . ± . . ± .
02 +0 . ± .
02 +0 . ± .
02 +0 . ± .
02 0 . ± . . ± .
05 +0 . ± .
04 +0 . ± .
04 +0 . ± . − . ± . . ± .
04 +0 . ± .
04 0 . ± . − . ± . . ± .
04 0 . ± .
04 +0 . ± .
04 +0 . ± . − . ± .
05 0 . ± . − . ± .
05 +0 . ± .
05 0 . ± . − . ± .
05 +0 . ± . − . ± .
05 +0 . ± . . ± .
05 +0 . ± .
04 +0 . ± .
04 +0 . ± . − . ± . . ± . − . ± . − . ± .
04 0 . ± .
03 +0 . ± . − . ± .
04 +0 . ± . . ± .
05 +0 . ± .
04 0 . ± . − . ± . − . − . − . − . +0 . ± . − . ± . − . ± . − . ± . − . ± . − . ± . − . ± . − . ± . − . ± . − . ± .
02 +0 . ± .
03 +0 . ± . − . ± .
03 +0 . ± . − . ± . − . ± .
03 +0 . ± . − . ± .
03 0 . ± . − . − . − . − . +0 . ± . − . ± .
04 0 . ± .
05 +0 . ± .
08 +0 . ± . . ± .
03 +0 . ± .
05 +0 . ± .
04 +0 . ± .
07 +0 . ± . . ± .
05 +0 . ± .
05 +0 . ± . . + . + . + . +0 . ± . . ± .
03 +0 . ± .
03 +0 . ± .
05 +0 . ± .
04 +0 . ± . . + . + . + . +0 . ± . . ± .
04 +0 . ± .
04 +0 . ± .
04 +0 . ± . − . ± . . + . + . + . [Eu II/Fe] + . + . + . + . . ± . Note : Abundances calculated by synthesis are presented in bold numbers.implausibly considered to be a systematic error arising from twosimilar but identical analyses .Abundances for the heaviest elements are based on eitherstrong lines (e.g., Ba) or on just one to three lines. Thus, thedifferences between OC and field giants may be in part due tosystematic errors. However, D’Orazi et al. (2009) and Maiorcaet al. (2011) concluded that heavy element abundances increasefrom old to young clusters. This study is a part of a project on the determination of chemicalabundances of OCs through high resolution spectroscopy, whosefinal goal is to chemically tag the disk field stars to re-constructthe dispersed stellar aggregates and to derive the abundance gra-dient in the Galactic disk. In this paper, we have presented ananalysis of giant stars in four OCs located between Rgc ∼ McWilliam et al. (2003) reported low [Mn/Fe] at [Fe/H] ≃ α - elements and Eu, characteristics not carried by the giantsin our OCs. the light elements (Na, Al), α -elements (Mg, Si, Ca, Ti), iron-peak elements (Sc, V, Cr, Mn, Fe, Co, Ni), light s-process ele-ments (Y, Zr), heavy s-process elements (Ba, La, Ce, Nd), and ofthe r-process elements (Sm, Eu).(iii) We have derived average [Fe/H] values of − ± − ± − ± − ± ∼ Acknowledgments:
DLL wishes to thank the Robert A. WelchFoundation of Houston, Texas for support throug grant F-634.
REFERENCES
Allen D. M., Porto de Mello G. F., 2011, A&A, 525, 63Allende Prieto C., Garc´ıa L´opez R. J., Lambert D. L., GustafssonB., 1999, ApJ, 527, 879.Alonso A., Arribas S., Mart´ınez-Roger C., 1999, A&AS, 140, 261Asplund M., Grevesse N., Sauval A. J., Scott, P., 2009, ARA&A,47, 481.Bartaˇsi¯ut˙e S., Deveikis V., Straiˇzys V., Bogdanoviˇcius 2007,BaltA, 16, 199.c (cid:13) , 1– ?? A. B. S. Reddy, S. Giridhar and D. L. Lambert
Bensby, T., Feltzing, S., Lundstr¨om, I., Ilyin, I., 2005, A&A, 433,185.Bi´emont E., Grevesse N., Hannaford P., Lowe R. M., 1981, ApJ,248, 867.Carretta E., Bragaglia A., Gratton R.G., Tosi M., 2004, A&A,422, 951.Carretta E., Bragaglia A., Gratton R. G., Tosi M., 2005, A&A,441, 131.Castelli, F., Kurucz, R. L., 2003, IAU Symposium 210, Modellingof Stellar Atmospheres, Uppsala, Sweden, eds. N.E. Piskunov,W.W. Weiss, and D.F. Gray, 2003, ASP-S210Clari´a J.J., Piatti A.E., Mermilliod J.-C., Palma T., 2008, AN,329, 609.De Silva, G.M., Sneden, C., Paulson, D.B., Asplund, M., Bland-Hawthorn, J., Bessell, M.S., Freeman, K.C., 2006, AJ, 131,455.De Silva G. M., Freeman K. C., Asplund M., Bland-Hawthorn,J., Bessell M. S., Collet, R. 2007, AJ, 133, 1161.Dias W. S., Alessi B. S., Moitinho A., L´epine J. R. D., 2002,A&A, 389, 871.D’Orazi V. et al. 2009, ApJ, 693, 31Edvardsson, B., Andersen, J., Gustafsson, B., Lambert, D.L., Nis-sen, P.E., Tomkin, J., 1993, A&A, 275, 101.F¨uhr J.R., Wiese W.L., 2006, J.Phys. Chem. Ref. Data, 35, 1669.Ghez, A. M. et al. 2008, ApJ, 689, 1044Hamdani S., North P., Mowlavi N., Raboud D., Mermilliod J.-C.,2000, A&A, 360, 509.Hannaford P., Lowe R. M.,Grevesse N., Bi´emont E., Whaling W.,1982,ApJ,261,736.Henderson C. B., Deliyannis C. P., Hughto J., Simmons A., Crox-all K., Sarajedini A., Platais I., 2007, AAS, 211, 5819.Hinkle K., Wallace L., Valenti J., Harmer D., 2000,
Visible andNear Infrared Atlas of the Arcturus Spectrum 3727-9300 ˚A (San Francisco: ASP).Jacobson H. R., Friel E. D., Pilachowski C. A., 2009, AJ, 137,4753.Kurucz R. L., 1998, http://cfaku5.harvard.edu.Kurucz R. L., Furenlid I., Brault J., & Testerman L. 1984, SolarFlux Atlas from 296 to 1300 nm, ed. R. L. Kurucz, I. Furen-lid, J. Brault, & L. Testerman (Sunspot, NM: National SolarObservatory)Lada C.J., Lada E.A., 2003, ARA&A, 41, 57.Lawler J. E., Bonvallet G., Sneden C., 2001, ApJ, 556, 452.Lawler J. E., Den Hartog E. A., Sneden C., Cowan J. J., 2006,ApJS, 162, 227.Lawler J. E., Wickliffe M. E., Den Hartog E. A., Sneden, C., 2001,AJ, 563, 1075.Lawler J. E., Sneden C., Cowan J.J., Evans I.I., Den Hartog E.A.,2009,ApJ,182,51.Lodders K., 2003, ApJ, 591, 1220.Luck R. E., Heiter U., 2006, AJ, 131, 3069.Luck R. E., Heiter U., 2007, AJ, 133, 2464.Magrini L., Sestito P., Randich S., Galli D., 2009, A&A, 494, 95Maiorca E., Randich S., Busso M., Magrini L., Palmerini S., 2011,ApJ, 736, 120McWilliam A., Rich R. M., Smecker-Hane, T. A. 2003, ApJ, 592,21McWilliam A., 1998, ApJ, 115, 1640.Mel´endez J., Barbuy B., 2009, A&A, 497, 611.Mermilliod J.-C., Mayor M., Udry S., 2008, A&A, 485, 303.Mishenina, T.V., Bienaym´e, O., Gornaeva, T.I., Charbonnel, C.,Soubiran, C., Korotin, S.A., Kovtyukh, V.V., 2006, A&A, 456,1109.Moore C.E., Minnaert M. G. J., Houtgast J., 1966,
The SolarSpectrum 2935 ˚A to 8770 ˚A , Second Revision of the Row-land’s Preliminary Table of Solar Wavelengths. National Bu-reau of Standards Monograph 61.Mucciarelli A., Caffau E., Freytag B., Ludwig H.-G., Bonifacio P., 2008, A&A, 484, 841.Paulson, D.B., Sneden, C., Cochran, W.D., 2003, AJ, 125, 3185.Parisi M.C., Clari´a J. J., Piatti A. E., and Geisler D., 2005, MN-RAS, 363, 1247.Prochaska J. X., McWilliam A., 2000, AJ, 537, 57.Reddy B.E., Tomkin, J., Lambert D.L., Allende Prieto C., 2003,MNRAS, 340, 304.Reddy B. E., Lambert D. L., Allende Prieto C., 2006, MNRAS,367, 1329.Sneden C., 1973, PhD Thesis, Univ. of Texas, Austin.Sobeck J.S., Lawler J.E., Sneden C., 2007, ApJ, 667, 1267.Takeda, Y., Sato, B., Murata, D., 2008, PASJ, 60, 781.Tull, R.G., MacQueen, P.J., Sneden, C., Lambert, D.L., 1995,PASP, 107, 251. c (cid:13) , 1– ?? GC 752, NGC 1817, NGC 2360, NGC 2506 Table 8.
Elemental abundances for stars in the OC NGC 752.Species star no. 77 star no. 137 star no. 295 star no. 311 Average[Na I/Fe] +0 . ± . . ± . . ± . . ± . . ± . − . ± . . ± . − . ± . − . ± . − . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . − . ± . . ± . . ± . . ± . . ± . − . ± . − . ± . − . ± . − . ± . − . ± . − . ± . − . ± . . ± . − . ± . − . ± . . ± . − . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . − . ± . − . ± . − . ± . . ± . − . ± . . ± . − . ± . . ± . . ± . . ± . − . − . − . − . − . [Fe I/H ] − . ± . − . ± . − . ± . − . ± . − . ± . − . ± . . ± . − . ± . − . ± . − . ± . − . ± . − . ± . . ± . . ± . − . ± . . ± . − . ± . − . ± . . ± . − . ± . − . − . − . − . − . [Zn I/Fe] − . ± . − . ± . − . ± . − . ± . . ± . − . ± . . ± . . ± . . ± . . ± . − . ± . . ± . . ± . . ± . . + . + . + . + . [La II/Fe] +0 . ± . . ± . . ± . . ± . . ± . . + . + . + . + . [Nd II/Fe] +0 . ± . . ± . . ± . . ± . . ± . . + . + . + . [Eu II/Fe] + . + . + . + . + . : The abundances calculated by synthesis are presented in bold numbers. The remaining elemental abun-dances were calculated using line equivalent widths. Numbers in the parentheses indicate the number of linesused in calculating the abundance of that element. In this analysis we have adopted the hfs data of Prochaska &McWilliam (2000) for Mn I, Mucciarelli et al. (2008) for Eu II line, McWilliam (1998) for Ba II line, and Allenet al. (2011) for Cu I lines.c (cid:13) , 1– ?? A. B. S. Reddy, S. Giridhar and D. L. Lambert
Table 9.
Elemental abundances for stars in the OC NGC 1817.Species star no. 1027 star no. 2038 star no. 2059 Average[Na I/Fe] +0 . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . − . ± . . ± . . ± . . ± . . ± . − . ± . . ± . . ± . − . ± . . ± . . ± . . ± . − . ± . − . ± . . ± . − . ± . − . ± . . ± . . ± . − . ± . . ± . . ± . . ± . − . − . − . − . [Fe I/H ] − . ± . − . ± . − . ± . − . ± . − . ± . − . ± . − . ± . − . ± . . ± . . ± . − . ± . . ± . . ± . − . ± . − . ± . − . ± . − . − . − . − . [Zn I/Fe] − . ± . . ± . − . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . − . ± . . ± . . + . + . + . [La II/Fe] +0 . ± . . ± . . ± . . ± . . + . + . + . [Nd II/Fe] +0 . ± . . ± . . ± . . ± . . + . + . [Eu II/Fe] + . + . + . : Same as in table 8. Table 10.
Elemental abundances for stars in the OC NGC 2506.Species star no. 2212 star no. 3231 star no. 4138 Average[Na I/Fe] +0 . ± . . ± . . ± . . ± . . ± . . ± . − . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . − . ± . − . ± . . ± . − . ± . . ± . . ± . . ± . − . ± . − . ± . . ± . . ± . − . ± . . ± . . ± . . ± . − . ± . . ± . . ± . . ± . − . ± . − . ± . . ± . − . ± . − . ± . − . ± . − . ± . − . ± . − . − . − . − . [Fe I/H ] − . ± . − . ± . − . ± . − . ± . − . ± . − . ± . − . ± . − . ± . − . ± . − . ± . . ± . − . ± . − . ± . − . ± . − . ± . − . ± . − . − . − . − . [Zn I/Fe] − . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . + . + . + . [La II/Fe] +0 . ± . . ± . . ± . . ± . . + . + . [Nd II/Fe] +0 . ± . . ± . . + . + . [Eu II/Fe] + . + . + . + . : Same as in table 8. c (cid:13) , 1– ?? GC 752, NGC 1817, NGC 2360, NGC 2506 Table 11.
Elemental abundances for stars in the OC NGC 2360.Species star no. 5 star no. 6 star no. 8 star no. 12 Average[Na I/Fe] +0 . ± . . ± . . ± . . ± . . ± . − . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . − . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . − . ± . . ± . . ± . . ± . . ± . − . ± . . ± . . ± . . ± . . ± . − . ± . − . ± . . ± . . ± . . ± . − . ± . . ± . . ± . . ± . . ± . − . ± . . ± . − . ± . − . ± . − . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . − . ± . . ± . . ± . − . − . − . − . [Fe I/H ] − . ± . − . ± . − . ± . − . ± . − . ± . − . ± . − . ± . − . ± . − . ± . − . ± . . ± . . ± . . ± . . ± . . ± . − . ± . . ± . . ± . . ± . . ± . − . − . − . − . − . [Zn I/Fe] +0 . ± . . ± . . ± . . ± . . ± . − . ± . . ± . . ± . . ± . . ± . . ± . − . ± . . ± . . + . + . + . [La II/Fe] +0 . ± . . ± . . ± . . ± . .
18 +0 .
19 +0 . . ± . . ± . . ± . . ± . . + . + . +0 . . + . + . + . : Same as in table 8. Table 12.
The linelist for all program stars from each of the OCs presented in this paper.EW(m˚A) for NGC 752 EW(m˚A) for NGC 1817 EW(m˚A) for NGC 2360 EW(m˚A) for NGC 2506 λ (˚A) Species a LEP b log gf Notes : a The integer part of the ’Species’ indicates the atomic number, and the decimal component indicates the ionization state(0 = neutral, 1 = singly ionized). b All the lines are arranged in the order of their increasing Lower Excitation Potential (LEP).Only a portion of this table is shown here for guidance regarding its form and content. A machine-readable version of thefull table is available.c (cid:13) , 1–, 1–