NGC 362: Another globular cluster with a split red giant branch

Abstract

We obtained FLAMES GIRAFFE+UVES spectra for both first and second-generation
red giant branch (RGB) stars in the globular cluster (GC) NGC 362 and used them
to derive abundances of 21 atomic species for a sample of 92 stars. The
surveyed elements include proton-capture (O, Na, Mg, Al, Si), alpha-capture
(Ca, Ti), Fe-peak (Sc, V, Mn, Co, Ni, Cu), and neutron-capture elements (Y, Zr,
Ba, La, Ce, Nd, Eu, Dy). The analysis is fully consistent with that presented
for twenty GCs in previous papers of this series. Stars in NGC 362 seem to be
clustered into two discrete groups along the Na-O anti-correlation, with a gap
at [O/Na] 0 dex. Na-rich, second generation stars show a trend to be more
centrally concentrated, although the level of confidence is not very high. When
compared to the classical second-parameter twin NGC 288, with similar
metallicity, but different horizontal branch type and much lower total mass,
the proton-capture processing in stars of NGC 362 seems to be more extreme,
confirming previous analysis. We discovered the presence of a secondary RGB
sequence, redder than the bulk of the RGB: a preliminary estimate shows that
this sequence comprises about 6% of RGB stars. Our spectroscopic data and
literature photometry indicate that this sequence is populated almost
exclusively by giants rich in Ba, and probably rich in all s-process elements,
as found in other clusters. In this regards, NGC 362 joins previously studied
GCs like NGC 1851, NGC 6656 (M 22), and NGC 7089 (M 2).

Full text

arXiv:1307.4085v1 [astro-ph.GA] 15 Jul 2013
Astronomy & Astrophysics
manuscript no. carretta
c
 ESO 2015
June 18, 2015
NGC 362: another globular cluster with a split red giant
branch
⋆,⋆⋆
E. Carretta
1
, A. Bragaglia
1
, R.G. Gratton
2
, S. Lucatello
2
, V. D’Orazi
3,4
, M. Bellazzini
1
, G. Catanzaro
5
, F. Leone
6
Y.
Momany
2,7
, and A. Sollima
1
1
INAF-Osservatorio Astronomico di Bologna, Via Ranzani 1, I-40127 Bologna, Italy
2
INAF-Osservatorio Astronomico di Padova, Vicolo dell’Osservatorio 5, I-35122 Padova, Italy
3
Dept. of Physics and Astronomy, Macquarie University, Sydney, NSW, 2109 Australia
4
Monash Centre for Astrophysics, Monash University, School of Mathematical Sciences, Building 28, Clayton VIC 3800, Melbourne,
Australia
5
INAF-Osservatorio Astrofisico di Catania, Via S.Sofia 78, I-95123 Catania, Italy
6
Dipartimento di Fisica e Astronomia, Universit`a di Catania, Via S.Sofia 78, I-95123 Catania, Italy
7
European Southern Observatory, Alonso de Cordova 3107, Vitacura, Santiago, Chile
ABSTRACT
We obtained FLAMES GIRAFFE+UVES spectra for both first and second-generation red giant branch (RGB) stars in the globular cluster (GC)
NGC 362 and used them to derive abundances of 21 atomic species for a sample of 92 stars. The surveyed elements include proton-capture
(O, Na, Mg, Al, Si), α−capture (Ca, Ti), Fe-peak (Sc, V, Mn, Co, Ni, Cu), and neutron-capture elements (Y, Zr, Ba, La, Ce, Nd, Eu, Dy). The
analysis is fully consistent with that presented for twenty GCs in previous papers of this series. Stars in NGC 362 seem to be clustered into
two discrete groups along the Na-O anti-correlation, with a gap at [O/Na]∼ 0 dex. Na-rich, second generation stars show a trend to be more
centrally concentrated, although the level of confidence is not very high. When compared to the classical second-parameter twin NGC 288,
with similar metallicity, but different horizontal branch type and much lower total mass, the proton-capture processing in stars of NGC 362
seems to be more extreme, confirming previous analysis. We discovered the presence of a secondary RGB sequence, redder than the bulk of
the RGB: a preliminary estimate shows that this sequence comprises about 6% of RGB stars. Our spectroscopic data and literature photometry
indicate that this sequence is populated almost exclusively by giants rich in Ba, and probably rich in all s-process elements, as found in other
clusters. In this regards, NGC 362 joins previously studied GCs like NGC 1851, NGC 6656 (M 22), and NGC 7089 (M 2).
Key words. Stars: abundances – Stars: atmospheres – Stars: Population II – Galaxy: globular clusters – Galaxy: globular clusters: individual:
NGC 362
1. Introduction
In the past few years the paradigm of multiple stellar popu-
lations as basic ingredient of Galactic globular clusters (GCs)
has become well assessed (see the review by Gratton, Sneden
& Carretta 2004, and the recent updates by Martell 2011 and
Gratton, Carretta and Bragaglia 2012, including references also
to the photometric evidence). At least two bursts of star for-
mation must have occurred in GCs, the second generation be-
ing formed from a mix of pristine gas and matter enriched by
nuclear processing in the massive stars of the first generation
(see Gratton et al. 2001). The impact of this chain of events
Send offprint requests to: E. Carretta, eugenio.carretta@oabo.inaf.it
⋆
Based on observations collected at ESO telescopes under pro-
gramme 083.D-0208
⋆⋆
Tables 2, 3, 4, 5, 6, 7 and 8 are only avail-
able in electronic form at the CDS via anonymous
ftp to cdsarc.u-strasbg.fr (130.79.128.5) or via
http://cdsweb.u-strasbg.fr/cgi-bin/qcat?J/A+A/???/???
is not limited to the history of GCs, but it probably had a key
rˆole in the building of a considerable fraction of the Galactic
halo (Carretta et al. 2010a, Vesperini et al. 2010, Martell et al.
2011). Some important links to global cluster parameters (e.g.
total mass, age, location in the Galaxy, Carretta et al. 2010a)
were discovered, and the ouput of extensive surveys performed
with multi-objects facilities like FLAMES (see Carretta et al.
2006, 2009a,b,c) pointed out that the characteristics of the nu-
clear processing occurring in early phases differ from cluster
to cluster. This evidence calls for an in-depth investigation of
a number of GCs with diff erent parameters, like metallicity,
concentration, horizontal branch (HB) morphology. This was
the original motivation of our ongoing FLAMES survey (see
Carretta et al. 2006).
In this paper we present an extensive analysis of chemical
abundances of multiple stellar populationsin the GC NGC 362.

2 E. Carretta et al.: Multiple stellar populations in NGC 362
This is a moderately metal-rich object ([Fe/H]= −1.26 dex
1
in
the 2011 update of the Harris 1996, H96 hereinafter, catalogue
that is based on the metallicity scale of Carretta et al. 2009c)
seen projected toward the Small Magellanic Cloud.
NGC 362 is variously classified as a red HB (RHB), a
young halo (Mackey and van den Bergh 2005) or an inner halo
(Carretta et al. 2010a) cluster. It belongsto a group of GCs with
low total orbital energy (small orbit sizes) and unusual orbital
parameters. The orbit has a chaotic behaviour and is confined
close to the Galactic plane (z
max
= 2.1 kpc), with high eccen-
tricity and small inclination off the orbital plane (Dinescu et al.
1999). An extensive variability survey in NGC 362 was made
by Sz´ekely et al. (2007), who find more than 45 RR Lyrae vari-
ables, a metallicity [Fe/H]= −1.16 dex, and a mean period of
RRab stars of 0.585 ± 0.081 days, which places this cluster in
the Oosterhoff type I group.
The pair NGC 362-NGC 288 is considered a classical ex-
ample of second parameter clusters: the colour of the HB of
these two GCs is very different (NGC 362: red; NGC 288:
blue) despite their similar metal-abundance. Catelan et al.
(2001) compared the relative age provided by the HB mor-
phology with that obtained from the main sequence for the pair
NGC 362-NGC 288. They used the bimodal HB of NGC 1851
as a “bridge” (see also Bellazzini et al. 2001). They found that
NGC 362 is about 2 Gyr younger than NGC 288, supporting
the concept that age is a main second parameter (beside metal-
licity) in shaping the distribution of stars along the HB. The
same result was obtained by Dotter et al. (2010) and Gratton et
al. (2010) from the analysis of homogeneous photometric and
spectroscopic databases. However, while Catelan et al. claimed
that the mass dispersion on the HB of NGC 362 is substantially
larger than for NGC 288, the opposite was found by Gratton
et al. According to their analysis, the mass spread required to
account for the HB morphology in NGC 362 is only 0.004 M
⊙
against 0.024 M
⊙
for NGC 288. This is due to the adoption of
an universal mass loss law with a linear dependence on metal-
licity, that well reproduces the median colours of HB stars.
Piotto et al. (2012) found that a few (actually, less then 3%)
subgiant branch stars of NGC 362 are fainter than the other
subgiants of similar colour. This faint sequence might be inter-
preted as a small group of objects having either older ages or
larger total CNO content than the majority of the stars.
Abundance analyses of stars in NGC 362 are rather scanty.
Several studies based on low dispersion spectroscopy deter-
mined the pattern of the CN and CH distribution in NGC 362,
typically found to be bimodal (Smith 1983, 1984; Kayser et al.
2008, Smith and Langland-Shula 2009 and reference therein).
Shetrone and Keane (2000) derived abundances for several el-
ements in a dozen of giants from high dispersion spectra and
compared their chemical pattern to that derived from 13 giants
in NGC 288. At the epoch (before the analysis of O, Na, Mg,
Al in unevolved cluster stars by Gratton et al. 2001) it was still
debated how much evolutionary (mixing) effects were super-
imposed to ab initio abundance variations. However, even from
1
We adopt the usual spectroscopic notation, i.e. [X]= log(X)
star
−
log(X)
⊙
for any abundance quantity X, and log ǫ(X) = log (N
X
/N
H
) +
12.0 for absolute number density abundances.
Fig.1.The V, B−V CMD of NGC 362 (light greycrosses). Stars
selected for the present study are plotted as filled, larger sym-
bols: blue squares are the stars observed with UVES, red circles
are stars with GIRAFFE spectra, and green triangles are stars
observed with GIRAFFE but not analyzed (see text). The faint
blue and red sequences are due to the young main sequence
stars and old red giants of the Small Magellanic Cloud.
their moderate-size samples, Shetrone and Keane (2000) were
able to point out differences in zero-point and slope existing
among the observed Na-O anti-correlation in this pair of GCs.
Recently Worleyand Cottrell (2010)used high dispersion spec-
troscopy to study the pattern of light and heavy neutron-capture
process elements in NGC 362.
In our study we enlarge the sample of stars analysed with
moderate-high resolution spectroscopy to more than 90 red
giant branch (RGB) stars in NGC 362, obtaining homoge-
neous abundances of proton-capture, α−capture, Fe-peak and
neutron-capture elements.
The paper is organized as follows: observations are pre-
sented in Section 2; radial velocities and a brief discussion of
the kinematics of the cluster are in Section 3; Section 4 presents
the derivation of theatmosphericparameters and the abundance
analysis; results of this analysis are given in Section 5; finally,
discussion and conclusions are in Section 6. The Appendix
presents the derivation of errors in the abundance analysis.
2. Observations
Our targets were selected from unpublished Johnson B, V Wide
Field Imager (WFI) photometryobtained at the 2.2mESO-Max
Planck Telescope (La Silla, Chile) and astrometrised by one
of us (Y. Momany), integrated with K band magnitudes from
the Point Source Catalogue of 2MASS (Skrutskie et al. 2006).
The V, B− V colour magnitude diagram (CMD) of NGC 362 is
shown in Fig. 1 with superimposed stars of our spectroscopic

E. Carretta et al.: Multiple stellar populations in NGC 362 3
Table 1. Log of FLAMES observations for NGC 362
Date UT exp. grating seeing airmass
(sec) (”)
Aug. 04, 2009 08:07:18.230 2600 HR11 1.02 1.454
Aug. 04, 2009 08:57:24.456 2600 HR11 0.88 1.442
Aug. 05, 2009 07:27:25.601 2600 HR13 0.65 1.482
Aug. 05, 2009 08:12:33.313 2600 HR13 0.86 1.449
sample, chosen to be near the RGB ridge line and with no
close companion. NGC 362 is seen projected against the Small
Magellanic Cloud, which is responsible for the faint blue and
red sequences visible in that diagram; they are young main se-
quence stars and old red giants, respectively.
The log of the observations is given in Table 1; we obtained
two exposures with the HR11 high resolution grating covering
the Na i 5682-88 Å doublet and two exposures with the grat-
ing HR13 including the [O i] forbidden lines at 6300-63 Å. We
observed a total of 14 (bright) giants with the fibres feeding
the UVES spectrograph (Red Arm, with spectral range from
4800 to 6800 Å and R=45,000; blue squares in Fig. 1) and
136 RGB stars with GIRAFFE (filled circles and triangles). All
stars turned out to be (likely) member of the cluster based on
radial velocities (RV). However, for the present analysis we re-
tained only stars with effective temperature below 5200 K be-
cause for warmer stars (green triangles) the low S/N and noisy
spectra prevented us from measuring accurate enough equiva-
lent widths (EWs)
2
. This limit correspondsto a magnitudelevel
V = 16.47, indicated by a solid line in the CMD.
We used the 1-D, wavelength calibrated spectra as reduced
by the ESO personnel with the dedicated FLAMES pipeline.
Radial velocities (RV) for stars observed with the GIRAFFE
spectrograph were obtained using the IRAF
3
task FXCORR,
with appropriate templates, while those of the stars observed
with UVES were derived with the IRAF task RVIDLINES.
One star (13875) observed with GIRAFFE shows the TiO
band head at 6158 Å in the HR13 spectrum (see Valenti et al.
1998). This star has many peculiarities beside being of spectral
type M: it is a variable (V2: Clement 1997) with a period of
90 days; it was found to be very Li-rich by Smith et al. (1999);
and it presents carbon dust according to Boyer et al. (2009). It
was discarded from the following analysis.
We have 12 stars in common between the UVES and
GIRAFFE datasets; disregarding star V2 our final sample in
NGC 362 consists of 92 RGB stars. Coordinates, magnitudes
and heliocentric RVs are shown in Tab. 2 (the full table is only
available in electronic form at CDS).
2
The hottest star left in our final sample is #5447, with an effective
temperature of 5203 K.
3
IRAF is distributed by the National Optical Astronomical
Observatory, which are operated by the Association of Universities
for Research in Astronomy, under contract with the National Science
Foundation
3. Kinematics
As discussed in Bellazzini et al. (2012; B12 hereinafter) the
samples of RV estimates that are obtained as a natural by
product of this project are not ideal for kinematical analyses.
However, because of the fibre allocation constraints described
in B12, in many cases they are nicely complementary to ex-
isting dataset as they preferentially probe the outskirts of the
clusters. This allowed us to reveal rotation signals in the outer
regions of severalclusters that went unnoticed in previous stud-
ies. In the following - as for the chemical analysis - we use
only stars cooler than 5200 K (red circles in Fig. 1 and below)
since only for these stars the membership is fully ascertained
based on both RV and chemical composition, following B12.
Stars warmer than this values are plotted as green circles (as
in Fig. 1) in the figures presenting kinematic results, for com-
pleteness. We perform the same kind of analysis as in B12, us-
ing the same techniques: we refer to that paper for details and
discussion.
The detailed analysis by Fischer et al. (1993: F93) is based
on 210 RV members of this cluster,all ofthem lying within R =
4.0
′
. On the other hand, our sample cover the range 1.0
′
≤ R ≤
11.0
′
, reaching the tidal radius of the clusters (r
t
= 10.3
′
, H96).
We merged our catalog with the one by F93, keeping our RV
when estimates from both sources are available. From a subset
of 38 stars in common (excluding two stars classified as binary
from RVs at two epochs by F93, our stars 11413 and 18947
4
)
we find a mean differenceof ∆RV=-1.15 km s
−1
with σ = 1.54;
this small zero point offset was corrected before merging the
two samples. The errors on individual RV estimates from the
two datasets are very similar (≃ 1.0 km s
−1
).
In the lower panel of Fig. 2 we show that the velocity dis-
persion profile we obtain fromour global sample ishardly com-
patible with a King (1966) model having the best-fit parameters
reported in H96 and scaled to the central velocity dispersion re-
ported there for this cluster, i.e. σ
0
= 6.4 km s
−1
. In the upper
panel it can be appreciated that even adopting a higher value,
providing an acceptable fit to the observedprofile (σ
0
= 7.5 km
s
−1
), there are obvious members lying outside the ±3− σ enve-
lope of the model, suggesting that some unbound or partially-
bound extra-tidal stars are present, or that King models are not
fully adequate to describe the structure and dynamics of the
system (see McLaughlin & van der Marel 2005, and Correnti
et al. 2011 for discussion and references).
Looking for rotation in the global sample we find no sig-
nificant signal, in agreement with F93. However if we limit to
our sample we find the A
rot
= 2.1 km s
−1
, and PA = 147deg
signal shown in Fig. 3, suggesting that some rotation may be
indeed present in the outer regions. Note that the warmer stars
excluded from the analysis (green points) appear to share the
same pattern. The trend toward V
r
− V
sys
∼ 0.0 at large ab-
solute values of X(PA
0
) is more consistent with an intrinsic
nature of the kinematic pattern, as a velocity gradient induced
4
All the stars classifiedas binaries by F93 have been excluded from
the analysis of the cluster kinematics. The four excluded stars are
H1348, H1419, H2205, H2222, in the nomenclature of F93. H1419
= 18947 and H2205= 11413 are also included in our own sample and
in the following chemical analysis.

4 E. Carretta et al.: Multiple stellar populations in NGC 362
Fig.2. Upper panel: radial velocity as a function of the distance
from the cluster centre. Colored symbols are the same as in
Fig. 1; filled triangles are stars from F93. The solid lines repre-
sents the ±3σ envelope of the King (1966) model that best fits
the surface brightness profile, according to H96. The profile
is normalized to the central velocity dispersion σ
0
= 7.5 km
s
−1
. Lower panel: velocity dispersion profile as derived from
the global sample (this work + F93, excluding stars lacking
abundance estimates). The solid line is the profile of the same
King (1966) model as above, the dotted line is the same model
normalized at the central velocity dispersion reported in H96,
σ
0
= 6.4 km s
−1
.
by Galactic tides would show larger amplitudes at larger dis-
tances. NGC 362 nicely fits into the correlations recently found
by B12.
Finally, we note that the velocity dispersion changes as a
function of the Na abundance, Na-poor stars having lower dis-
persion than Na-rich ones. This is due to the radial trend shown
by the different concentration of Na-poor and Na-rich stars, the
first preferentially populating more external regions (see be-
low). There is no significant trend of the rotation pattern with
Na abundance, but it is very likely that our sample is too sparse
to investigate such subtle effects.
4. Atmospheric parameters, abundance analysis,
and metallicity
Only a brief summary of the analysis methods will be reported
here, since the atmospheric parameters were derived following
the same procedure adopted for the other GCs targeted by our
FLAMES survey (see e.g. Carretta et al. 2009a,b).
First-pass input effective temperature T
eff
(from V − K
colours) and bolometric corrections were derived from the cal-
Fig.3. Rotation in NGC 362. Upper panel: difference between
the mean velocities on each side of the cluster with respect to
a line passing through the cluster centre with a position an-
gle (PA, measured from north to east, north=0
o
, east=90
o
), as
a function of the adopted PA. Our preferred solution is repre-
sented by the dotted line and is obtained from bona-fide mem-
bers from our sample (i.e. stars with abundance estimates).
Lower panels: the rotation curve of NGC 362. In the left panel
the RV in the system of the cluster is plotted as a function of
distance from the centre projected onto the axis perpendicular
to the best fit rotation axis found in the upper panel. the mean-
ing of the symbols are the same as in Fig. 1 and 2, above. The
right panel shows the comparison of the cumulative RV dis-
tributions of stars having X(PA
0
) > 0.0 (continuous line) and
X(PA
0
) < 0.0 (dotted line).
ibrations of Alonso et al. (1999, 2001). These values were
refined using a relation between T
eff
and K magnitudes
5
.
Gravities were obtained from stellar masses and radii, these
last derived from luminosities and temperatures. The redden-
ing and the distance modulus for NGC 362 were taken from
the Harris (1996) catalogue (2011 update). We adopted a mass
of 0.85 M
⊙
for all stars and M
bol,⊙
= 4.75 as the bolometric
magnitude for the Sun, as in our previous studies.
The abundance analysis mainly rests on equivalent widths
(EWs). We checkedthat EWs measured on the GIRAFFE spec-
tra are on the same system defined by high-resolution UVES
spectra. The values of the microturbulent velocities v
t
were
obtained by eliminating trends between Fe i abundances and
expected line strength (Magain 1984). Finally, models with
the appropriate atmospheric parameters and whose abundances
matched those derived from Fe i lines were interpolated within
5
For a few stars with no K in 2MASS we obtained an estimate of
the K values by interpolating a quadratic relation as a function of V
magnitudes.

E. Carretta et al.: Multiple stellar populations in NGC 362 5
the Kurucz (1993) grid of model atmospheres (with the op-
tion for overshooting on) to derive the final abundances. The
adopted atmospheric parameters and iron abundances are listed
in Tab. 3.
Our procedure for error estimates is detailed in Carretta et
al. (2009a,b); results, with a brief description, are given in the
Appendix for UVES and GIRAFFE observations (see Tab. A.1
and Tab. A.2, respectively).
Oxygen abundances were obtained from the [O i] line at
6300.3 Å (after cleaning this spectral region from telluric lines
as described in Carretta et al. 2007a) and, whenever possible,
from the [O i] line at 6363.8 Å. Sodium abundances, derived
from the EWs of the 5682-88 and 6154-60 Å doublets, were
corrected for departures from the LTE assumption following
Gratton et al. (1999). Abundances of O, Na, Al (from the 6696-
99 Å doublet), and Mg (derived as in Carretta et al. 2009b) for
individual stars are listed in Tab. 4.
Beside Mg (reported among elements involved in proton-
capture processes), we measured the abundances of the
α−elements Si, Ca, and Ti either from UVES or GIRAFFE
spectra (see Tab.5)
6
. We measured Ti abundances from lines of
both neutral and singly ionised species on UVES spectra (with
larger spectral coverage). On average, the abundances obtained
from these two ionization states are in excellent agreementwith
each other (see below), supporting our adopted scale of atmo-
spheric parameters.
We obtained abundances for the Fe-peak elements Sc ii,
V i, Cr i, Co i, and Ni i (stars with GIRAFFE and UVES spec-
tra) and additionally Mn i and Cu i for stars with UVES spec-
tra only. The corrections due to the hyperfine structure (see
Gratton et al. 2003 for references) were applied whenever rele-
vant (e.g. Sc, V, Mn, Co). We obtained the abundances of sev-
eral neutron capture elements (Y ii, Zr ii, Ba ii, Laii, Ce ii, Nd ii,
Eu ii, and Dy ii, mostly from UVES spectra) using a combina-
tion of spectral synthesis and EW measurements. Details on
transitions can be found in Carretta et al. (2011).
Abundances of Fe-peak elements for individual stars are re-
ported in Tab. 6. Results for neutron capture elements are given
in Tab. 7 and Tab. 8 for stars with UVES and GIRAFFE spec-
tra, respectively. The averages of all measured elements with
their r.m.s. scatter are listed in Tab. 9.
The mean metallicity we found for NGC 362 from stars
with UVES spectra is [Fe/H]= −1.168 ± 0.014 ± 0.051 dex
(rms = 0.052 dex, 14 stars), where the first error bar is from
statistics and the second one refers to the systematic effects.
From the large sample of stars with GIRAFFE spectra we de-
rived a value of [Fe/H]= −1.174 ± 0.004 ± 0.062 dex (rms =
0.041 dex, 90 stars). The metal abundance of NGC 362 is very
similar to that of NGC 1851 ([Fe/H]= −1.19 dex, Carretta et al.
2011) and 0.14 dex higher than the metal content of NGC 288
([Fe/H]= −1.31 dex, Carretta et al. 2009c) on our homoge-
neous metallicity scale.
The abundances of iron obtained from singly ionized
species are in nice agreement with those from neutral lines:
[Fe/H]ii= −1. 21 (rms = 0.08, dex 14 stars) from UVES and
6
The case of Si as an element involved in the p-capture reactions is
discussed in Section 5.1
Table 9. Mean abundances from UVES and GIRAFFE
Element UVES GIRAFFE
n avg rms n avg rms
[O/Fe]i 14 +0.14 0.20 64 +0.89 0.18
[Na/Fe]i 14 +0.19 0.19 90 +0.11 0.25
[Mg/Fe]i 14 +0.33 0.04 84 +0.33 0.04
[Al/Fe]i 14 +0.24 0.19
[Si/Fe]i 14 +0.22 0.04 87 +0.26 0.04
[Ca/Fe]i 14 +0.30 0.03 89 +0.34 0.02
[Sc/Fe]ii 14 −0.03 0.05 90 −0.07 0.04
[Ti/Fe]i 14 +0.22 0.04 84 +0.16 0.03
[Ti/Fe]ii 14 +0.21 0.05
[V/Fe]i 14 −0.03 0.03 86 −0.05 0.03
[Cr/Fe]i 14 −0.02 0.07 88 −0.03 0.04
[Mn/Fe]i 14 −0.33 0.04
[Fe/H]i 14 −1.17 0.05 90 −1.17 0.04
[Fe/H]ii 14 −1.21 0.08 69 −1.18 0.06
[Co/Fe]i 14 −0.28 0.06 36 −0.05 0.09
[Ni/Fe]i 14 −0.13 0.03 90 −0.09 0.04
[Cu/Fe]i 14 −0.50 0.12
[Zn/Fe]i 12 +0.21 0.06 43 +0.28 0.08
[Y/Fe]ii 14 +0.07 0.11
[Zr/Fe]ii 14 +0.50 0.12
[Ba/Fe]ii 14 +0.30 0.27 68 +0.18 0.21
[La/Fe]ii 14 +0.33 0.09
[Ce/Fe]ii 14 +0.14 0.12
[Nd/Fe]ii 14 +0.35 0.10
[Eu/Fe]ii 14 +0.70 0.07
[Dy/Fe]ii 14 +0.68 0.13
[Fe/H]ii= −1. 18 (rms = 0.06 dex, 69 stars) from GIRAFFE.
They do not present any trend as a function of the effective
temperature, as shown in Fig. 4.
Our Fe abundanceis inverygoodagreementwith that listed
in the old (1996) version of the Harris catalogue and with that
obtained by Szekely et al. (2007) from analysis of the pulsa-
tional properties of RR Lyrae. It is higher (although in agree-
ment within the uncertainties) than the value listed in the most
recent version of the Harris catalogue, and of the average val-
ues obtained by Shetrone and Keane (2000)and Kraft and Ivans
(2003).
5. Results
5.1. Chemistry of multiple populations in NGC 362
The Na-O anti-correlation in NGC 362 is based on 71 stars
for which we were able to provide O and Na abundances from
UVES and/or GIRAFFE spectra (we have 58 detections and 13
upper limits for O; the total number of stars with [Na/Fe] values
is 92). Results are shown in Fig. 5.
The interquartile range IQRfor the ratio [O/Na] inthis clus-
ter is 0.644 dex. This value nicely fits into the correlation with
the total absolute magnitude of the cluster (M
V
= −8.41, Harris
1996) established in Carretta et al. (2010a). NGC 362 also par-
ticipates to the relation with the maximum temperature along
the HB (logT
eff
= 4.079, from Recio-Blanco et al. 2006), dis-
covered by Carretta et al. (2007b) and repeatedly updated and

6 E. Carretta et al.: Multiple stellar populations in NGC 362
Fig.4. Abundance ratios [Fe/H] i (upper panel) and [Fe/H] ii
(lower panel) as a function of T
eff
for all analysed stars. Blue
squares are stars with UVES spectra, filled circles are those
with GIRAFFE spectra observed with both HR11 and HR13
gratings, whereas crosses indicate stars observed with only the
HR11 (black) or the HR13 (green) grating.
Fig.5. The Na-O anti-correlation observed in NGC 362. Blue
squares are stars observed with UVES, while filled circles indi-
cate stars with GIRAFFE spectra. Upper limits in O are shown
as arrows, and star-to-star (internal) error bars are plotted in the
upper-right corner of the plot.
Fig.6. The distribution of the [O/Na] ratios in NGC 362 from
the combined sample UVES+Giraffe.
strengthened during our FLAMES survey. NGC 362 stars seem
to be clustered along the Na-O anti-correlation in two distinct
groups, one with high O and low Na and the other more en-
riched in Na and depleted in O. This sub-division is even more
evidentif we consider only the 14 stars observedat higher reso-
lution with UVES, and it is nicely supported by the distribution
of the [O/Na] ratios shown in Fig. 6, that maximise the signal
along the Na-O anti-correlation. We also note that the same
separation into two groups was already present in the indepen-
dent analysis by Shetrone and Keane (2000) (see their Fig.2).
Using the quantitative criteria introduced by Carretta et al.
(2009a) we can use O and Na abundances to quantify the size
of the primordial (P) population and of the intermediate (I) and
extreme (E) components of the second generation stars. We
found that the fractions of P, I, and E stars for NGC 362 are
22 ± 6%, 75 ± 10%, and 3 ± 2%, respectively. We note that
according to these criteria, the lower Na-high-O group visible
in Fig. 5 includes not only the P component ([Na/Fe]< −0.03
dex, in NGC 362), but also part of the intermediate fraction I.
We warn however that the distribution of stars along the Na-O
anticorrelation from GIRAFFE spectra may be not optimally
suited to study the discrete vs continuous nature of multiple
populations.
The radial distribution of these components is shown in
Fig. 7, where the second generation stars of the I and E com-
ponent are merged together, since only two stars belong to the
E fraction. We find that second generation stars are more cen-
trally concentrated, a result already found in several other clus-
ters (see e.g. Lardo et al. 2011, Nataf et al. 2011, Milone et al.
2012).
The interplay of the various elements involved in proton-
capture processes in NGC 362 is summarized in Fig. 8 for the
stars observed with UVES. For this sub-sample, many light el-
ement abundances (O, Na, Mg, Al, Si) are available from our

E. Carretta et al.: Multiple stellar populations in NGC 362 7
Fig.7. Cumulative distribution of the stars of first (primordial
P component) and second generation (the merged I+E compo-
nents together) as a function of the projected distance from the
cluster centre.
Fig.8. Relations of the [Al/Fe] ratios in NGC 362 (filled blue
squares) from UVES spectra as a function of [Na/Fe] (upper
left panel), [Mg/Fe] (upper right), [O/Fe] (lowerleft panel), and
[Si/Fe] (lower right). The star to star error bars are indicated in
each panel. As a comparison, the open circles are results for
giants in NGC 288 from Carretta et al. (2009b).
data, thanks to the large spectral coverage. As a reference, we
also plot in this figure the results for 10 RGB stars observed
with UVES in NGC 288 by Carretta et al. (2009b). The range
of Al variations is not extreme among giants in NGC 362, yet
it is almost twice the spread in Al observed in NGC 288. The
Al-O anti-correlation is well developed in NGC 362, whose O
Fig.9. [Mg/Fe] ratios (upper panels) as a function of [O/Fe]
(left) and [Na/Fe] (right) ratios from GIRAFFE spectra. The
same relations for [Si/Fe] are plotted in the lower panels. The
internal errors are indicated in each panel.
abundances reach a level significantly lower than in NGC 288,
as found by Shetrone and Keane (2000). All these features can
be explained by the larger mass of NGC 362, which is almost
five times more massive than NGC 288 (see McLaughlin and
van der Marel 2005).
Due to the limited size of the sub-sample of stars with
UVES spectra, it is not possible to evince whether the pollut-
ing matter producing second generation stars was processed at
very high temperature. As discussed by Yong et al. (2005) and
Carretta et al. (2009b), when the H-burning temperature ex-
ceeds 65 MK a leakage from the Mg-Al cycle on
28
Si occurs
and some amount of Si is produced at the expense of Mg, be-
side the main outcome, Al. To verify that this is the case also
for NGC 362 we show in Fig. 9 the run of [Mg/Fe] and [Si/Fe]
ratios as a function of O and Na for the much larger sample of
giants observed with GIRAFFE in this cluster.
Mg is correlated to elements depleted in proton-capture re-
actions and anti-correlated to elements which are enhanced in
this burning, and the opposite is seen to occur for Si. All these
relations are statistically robust, the level of confidence for-
mally exceeds 99% in all four cases. However, the relations
involving Si seem to be driven by a few stars only. While the
stars with UVES spectra give Mg abundances that nicely fall
on the relationsdefined by the GIRAFFE sample involving Mg,
the agreement is not so satisfactory concerning Si. In summary,
it is not clear that the first generation polluters were of the
right mass (inner temperature) range to significantly change
the Si abundance above the level typical of enrichment from
type II supernovae. In NGC 362 we did not observe the typical
Si-Al correlation that in massive or metal-poor GCs (such as
NGC 2808 or NGC 6752, Carretta et al. 2009b) suggest that

8 E. Carretta et al.: Multiple stellar populations in NGC 362
Fig.10. Abundance ratios of α−elements Mg, Si, Ca, Ti i as
a function of the effective temperature. The average of [α/Fe]
ratios are shown in the last two panels on the right column (in-
cluding and excluding the Mg abundance from the mean, re-
spectively). Error bars indicate internal star-to-star errors.
Si is partly produced by proton-capture reactions, beside the
classical α−capture processes.
5.2. Other elements
The pattern of the α− and Fe-group element abundances is
summarized in Fig. 10 and Fig. 11 where the abundances are
plotted as a function of the effective temperature for individual
stars in NGC 362.
In the last panel (bottom right) of Fig. 11 we also show
the abundance ratios of Ba, which is a neutron-capture ele-
ment whose abundance in the solar system is mostly due to the
s−process. None of these elements present a trend as a function
of temperature, and as shown in Table 9 their abundances are
usually very homogeneous in NGC 362, apart from the light
elements involved in proton-capture reactions discussed in the
previous Section. The case of Ba with its large spread will be
discussed in detail in Section 5.4 below.
Analysis of the UVES spectra allows the measurements of
the abundances of several more n-capture elements, including
Cu, Y, La, Ce, Nd, Eu, and Dy, beside Ba which is measured
also from GIRAFFE spectra. Since stars observed with UVES
are typically very cool and luminous, some of the relevant
lines are strong and then derived abundances are highly sen-
sitive to microturbulence velocity. For some lines we obtained
strong correlation between abundances and the adopted micro-
turbulent velocities. As an example, Fig. 12 shows the correla-
tions obtained for three clean Y ii lines. The lines at 4883.7 and
5087.4 Å show a marked trend, therefore we only used the line
at 5200.4 Å. With similar considerations, we selected only Nd
Fig.11. Abundance ratios of elements of the Fe-peak (Sc ii,
V, Cr, Co, Ni) and the s−process element Ba from GIRAFFE
spectra as a function of the effective temperature. Error bars
indicate internal star to star errors.
lines with EW smaller than 70 mÅ, yielding typically six lines
per stars. In the case of Ba the three available lines, 5853.7,
6141.7, 6469.9 Å are all very strong, and thus all sensitive to
microturbulence. Fig. 13 shows [Zr/Fe], [Ba/Fe], and [La/Fe]
as a function of microturbulent velocity. The trend of Ba with
microturbulence is quite obvious. However, part of the trend
is due to star 11413, which is richer in all n-capture elements
than the other giants in the UVES sample. When this star is ex-
cluded, the trend between Ba abundances and microturbulence
values is scarcely significant, at less than two-sigma.
We note that in order to further flatten the trend of Ba (and
of the strong Nd and Y lines) with microturbulence, the star-
to-star variation of such parameter should be decreased, i.e. the
low microturbulence values increased and the high ones de-
creased. However, this would create a trend of the Fe abun-
dances with microturbulence of the opposite sign. Given that
the measured Fe lines typically form deeper in the stellar atmo-
spheres than the strong lines of heavy elements, this opposite
behaviour with respect to microturbulence hints at an intrin-
sic inadequacy of the model atmosphere for cool and luminous
stars. Note that this effect is likely not as strong in those stars
observed using GIRAFFE (see Fig. 11), as they are typically
warmer and fainter. Therefore, the derived errors for Ba indi-
cated in Table A.1 for the UVES stars should be taken with
caution, as the true errors are likely larger, due to the very high
sensitivity to the adopted microturbulent velocity. Finally, we
note that an even stronger trend of Ba abundances with micro-
turbulent velocity is present in the analysis by Shetrone and
Keane (2000), albeit in that case Ba abundances increase as V
t
increases.

E. Carretta et al.: Multiple stellar populations in NGC 362 9
Fig.12. logǫ(Y ii) as derived from the three available clean
lines in our spectral range as a function of micro-turbulent ve-
locity. Star 11413 on the secondary, red RGB is indicated with
a red open circle.
Fig.13. Abundances of Zr, Ba and La as a function of micro-
turbulent velocity for stars with UVES spectra. The empty cir-
cle indicates star 11413 (see text).
Fig. 14, Fig. 15, and Fig. 16 show the derived abundances
for the heavy elements from UVES spectra as a function of
T
eff
. There is a weak trend of Ba abundances with T
eff
due to
the above discussed correlation between derived Ba abundance
and microturbulent velocity, given that the latter has a trend
with effective temperature. Apart from the case of Ba, the mea-
sured n-capture elements abundances are quite uniform, show-
ing remarkably small scatter, especially when star 11413 is ex-
cluded from the average. This giant has the highest abundance
of Y, Ba, La, Nd, Ce, and Dy among the sample of 14 stars
with UVES spectra, whereas its Eu abundance does not stand
out among the other stars in this sample.
Fig.14. Abundance ratios of elements of the n-capture ele-
ments Zr ii, Y ii, and Cu i, from UVES spectra as a function
of the effective temperature. Error bars indicate internal star-to-
star errors. In each panel, the Ba-rich star 11413 is highlighted.
Fig.15. As in Fig. 14 for the n-capture elements Ba ii, La ii,
and Ce ii.
5.3.
s−
and
r−
process contributions
The abundances of neutron capture elements derived from
UVES spectra can be used to estimate the relative weight of
the s− and r−process contributions to the chemical pattern in
NGC 362, as done in Carretta et al. (2011) for NGC 1851 and
in D’Orazi et al. (2011) for ω Cen. We refer to those papers for
the details on the adopted procedure. Results for Ba, La, Ce,

10 E. Carretta et al.: Multiple stellar populations in NGC 362
Fig.16. As in Fig. 14 for the n-capture elements Eu ii, Nd ii,
and Dy ii.
and Eu are well consistent with each other: the pattern of abun-
dances is reproduced if we assume that most of these elements
are made by the r−process (1/3 of the solar r−process abun-
dances) with a small contribution from the s−process (about
1/25 of the solar s−process abundances).
On the other hand, results for the first-peak elements Y and
Zr are discrepant: we found less Y and too much Zr with re-
spect to the above recipe. The uncertainties in the determina-
tion of Y and Zr do not allow to determine with precision the
pattern of the first peak species. Since we have well defined
abundances only for the second peak, we are unable to calcu-
late the ratio between heavy and light s−process elements in
NGC 362. Hence, nothing more can be said about the origin of
these elements; in turn, the timescales of the enrichment cannot
be properly determined.
5.4. Ba and the red sequence on the RGB
The Ba abundances for the about 70 stars observed with
GIRAFFE are based on a single line at 6141 Å, and the asso-
ciated internal error is rather large (∼ 0.15 dex). However, the
r.m.s. scatter of the mean (0.21 dex) is larger than the internal
errors (see Fig. 11, bottom right panel, and Fig. 15). Additional
evidence that the spread may be real, albeit small, comes from
the Str¨omgren photometry. We used the photometry collected
by Grundahl and coworkers, and presented by Calamida et al.
(2007)
7
to whom we refer for details. There are 69 stars hav-
ing both photometric data and abundances of Ba. We show the
Str¨omgren y, v− y CMD for these stars in Fig. 17, where we in-
dicated with different symbols stars with Ba abundances lower
or higher than the average value of [Ba/Fe]=+0.209 dex. Stars
7
The catalogue was downloaded from the web page
http://www.oa-roma.inaf.it/spress/gclusters.html
Fig.17. Str¨omgren y, v−yCMD for stars of our sample with ac-
curate photometry. The stars are divided according to their Ba
abundances: blue filled circles are stars with [Ba/Fe] lower than
the average value [Ba/Fe]=+0.209 dex, whereas red open cir-
cles indicate stars with Ba content larger than the mean value.
with lower than average Ba abundances define a very narrow
sequence at the blue ridge of the RGB, while a fraction of the
Ba-rich stars populate a sequence slightly to the red of the main
RGB (apart from an outlier with a very blue v − y colour likely
due to errors in the photometry). We will come back to this red
sequence in Section 6.
We note that the above discussed uncertainty in the Ba
abundance due to microturbulence does not affect this result.
In fact, at a given effective temperature T
ef f
, the relative Ba
abundances are robust with respect to microturbulence, there-
fore at a given magnitude the separation between Ba-rich and
Ba-poor stars is reliable.
This behaviour reproduces the splitting of the RGB of
NGC 1851 observed in the v − y colours, as shown in Fig. 18
using data from Carretta et al. (2011). In both cases, the reddest
sequence is populated entirely by Ba-rich stars.
What is the physical explanation for this phenomenon? In
Carretta et al. (2011) we suggested that the redder v − y colour
of this secondary sequence is most likely due to large enhance-
ments in N. Very recently, this suggestion was supported by the
analysis of N abundances based on a number of CN features
measured on GIRAFFE spectra (Carretta et al. 2013, in prepa-
ration), as well as on a limited sample of giants by Villanova
et al. (2010) in NGC 1851: the stars populating the reddest se-
quence all have strong CN bands and higher N abundances.
The low resolution data for bright giants by Smith &
Langland-Shula (2009) may be used to check if the same holds
for NGC 362. These authors provided values of the index
S(3839) which measure the strength of the CN band strength.
These values were converted into the quantity δS(3839) as
defined by Norris (1981) using for the baseline of equation

E. Carretta et al.: Multiple stellar populations in NGC 362 11
Fig.18. Str¨omgren y, v − y CMD for stars of NGC 1851 from
Carretta et al. (2011). The stars are divided according to their
Ba abundances: blue filled circles and red open circles are stars
with [Ba/Fe] lower and higher, respectively than the average
value [Ba/Fe]=+0.492 dex.
S(3839)= 0.44 + 0.0878 × M
V
based on the data in Smith &
Langland-Shula(2009). Note that the stars in commonbetween
the two samples span a very limited range in magnitude and
colour along the RGB of NGC 362 (−2.23 < M
V
< −1.92).
Hence, any variation of the δS(3839) index with luminosity
and/or temperature has no impact in our discussion.
We found that, beside to the already knowncorrelation with
Na and anti-correlation with O abundances, this index is also
well correlated with the [Ba/Fe] ratio (we notice that the same
effect was found in M 22 by Marino et al. 2011). This is shown
in Fig. 19, where we adopted Ba abundances from both the
few stars of the present study and the more numerous giants in
common with Shetrone and Keane (2000). Among the six stars
in our sample, star 2333, with the highest value of δS(3839),
lies on the red RGB sequence.
Since Ba is correlated with Na and anti-correlated with O
in this subset of stars, we checked that the same also holds in
our much larger sample observed with GIRAFFE. In Figs. 20
and 21 we show the run of Ba abundances as a function of the
content of O, Na, Mg, and Si. The best (anti-)correlations are
those between Ba and the elements that are depleted in proton-
capture reactions, like O and Mg. The linear correlation coef-
ficients are r = −0.25 and −0.27, respectively (with 61 and 62
degrees of freedom, respectively). These anti-correlations, al-
though not striking, are significant to a level of confidence of
about 95%, admittedly not very high, while the complementary
correlations with Na and Si are not statistically significant.
At the moment, the reason why a typical s−process element
like Ba must be related to the abundances of elements forged in
the nuclear processing of likely much more massive stars of the
first stellar generation in the cluster still remains an unsolved
Fig.19. The index S(3839) of the CN bandstrength from Smith
and Langland-Shula (2009) as a function of the [Ba/Fe] ratio
from Shetrone and Keane (2000, red filled point) and from the
present work (blue open circles) for stars of NGC 362.
Fig.20. Upper panel: [Ba/Fe] as a function of [O/ Fe] in giants
of NGC 362 from this study. Lower panel: the same, as a func-
tion of [Na/Fe]. The Pearson linear correlation coefficient and
the Spearman rank correlation coefficient, together with star to
star error bars, are indicated in each panel.
issue, owing to the much different lifetimes of typical polluters
for these elements. More efforts should be focused in the future
in ascertaining the groundof these possiblecorrelations inlarge
samples of stars.

12 E. Carretta et al.: Multiple stellar populations in NGC 362
Fig.21. As in the previousfigure, butfor [Mg/Fe] (upperpanel)
and [Si/Fe] (lower panel).
Fig.22. Interquartile range (IQR) of the [O/Na] ratio in
NGC 362 as a function of the total absolute magnitude of the
cluster (left panel), and of the maximum temperature along
the HB (right panel). NGC 362 is represented by a filled
star symbol. The filled square is for ω Cen from Johnson
and Pilachowski (2010) and the filled triangle indicates M 54
(Carretta et al. 2010c). The other clusters are from Carretta et
al. (2009a, 2011). In the left panel, the location of NGC 288 is
indicated. In each panel the Spearman rank correlation coeffi-
cient (r
s
) and the Pearson’s correlation coefficient are reported.
6. Discussion and conclusions
We studied the chemical composition of 92 red giant branch
stars in the globular cluster NGC 362.
We found that NGC 362 behaves like a perfectly ’normal’
cluster for what concerns the elements involved in high tem-
perature H-burning. It nicely falls in the middle of the relation
between the extension of the Na-O anti-correlationand the total
absolute magnitude (a proxy for the present-day cluster mass)
defined by the other GCs (Carretta et al. 2010a; see left panel
of Fig. 22). The same holds for the correlation of the extension
of Na-O anti-correlation with the extent of the HB distribution
(right panel of Fig. 22).
In NGC 362 there is a hint that stars are clustered into two
discrete groups along the Na-O anti-correlation: this comes
both from the more conspicuous sample of RGB stars with
GIRAFFE spectra and the more limited fraction of stars with
high-resolution FLAMES/UVES spectra.
First-generation stars (Na-poor and O-rich) are apparently
less concentrated than second generation stars, in agreement
with the result found in most GCs both from spectroscopy (e.g.
Carretta et al. 2010d) and photometry (e.g. Lardo et al. 2011,
Milone et al. 2012 and references therein); however, this find-
ing is not supported by a high level of statistical confidence in
NGC 362.
In this context, it is interesting to compare NGC 362 to its
twin second-parameter cluster NGC 288 (Shetrone and Keane
2000, and, in particular, Carretta et al. 2009a,b). The latter GC
is indicated in Fig. 22 (left panel) to show its position slightly
off the main relation. Carretta et al. (2010a) already discussed
evidence (like tidal tails) pointing toward a large loss of stars
for NGC 288 in the past. Another evidence in this sense is given
by the study by Paust et al. (2010), who determined the global
present day mass functionsfor 17 GCs, including NGC288 and
NGC 362. Their comparison shows that NGC 288 has a flatter
mass function (the slopeofthe powerlaw fit is α = −0.83,com-
pared to α = −1.69 of NGC 362), indicating a lower fraction of
low mass stars. Since NGC 288 is a more loose cluster than its
twin NGC 362, this seems to imply that is also more affected by
external processes of evaporation through disk shocking and/or
tidal stripping (see Paust et al. 2010), well explaining its posi-
tion in the plane IQR[O/Na] vs M
V
.
Also relevant in the context of multiple populations in
GCs is the recent addition of NGC 362 to the increasing
number of clusters where a broad and/or even split subgiant
branch (SGB) was detected. The studies by Milone and col-
laborators include: NGC 1851 (Milone et al. 2008), 47 Tuc
(Anderson et al. 2009, Milone et al. 2012), NGC 6388 and
(maybe) NGC 6441 (Bellini et al. 2013), NGC 362, NGC 5286,
NGC 6656, NGC 6715, and NGC 7089 (Piotto et al. 2012).
Our FLAMES survey, devoted to study the abundances of
proton-capture elements (mainly Na and O) in large samples of
RGB stars in a large number of GCs, is blind to the feature of
double SGBs, which is governed by age and/ or CNO content
(see e.g. Cassisi et al. 2008). However, as already noted in pre-
vious section, a further peculiarity of NGC 362 is its secondary
giant branch in CMD using the v − y color (see Fig. 17).
The reality of the sequenceand the membershipof this pop-
ulation to the cluster is well assessed in the present work. Our
spectroscopic analysis shows that all the stars on this secondary
giant branch havemeasured RV and metallicity consistent with
the cluster. This sequence is neither grossly contaminated by
Galactic field interlopers, nor by stars of the Small Magellanic
Cloud, found at fainter magnitudes and well separated from the
stars in NGC 362 in the y, v − y plane. Thus, beside the split
SGB, NGC 362 seems to be part of the group of GCs where
a secondary giant branch is found in the y, v − y plane. The
most relevant cases include NGC 1851, NGC 6656 (M 22),

E. Carretta et al.: Multiple stellar populations in NGC 362 13
and NGC 7089 (M 2: Lardo et al. 2012). With the addition of
NGC 362 the phenomenon becomes common enough to jus-
tify the dubbing of “secondary RGB” for this feature, since it
seems to enclose only a minor fraction of the total red giants,
although the percentage does vary from cluster to cluster.
In the present study we showed that the “secondary RGB”
of NGC 362 is only populated by stars with high Ba abun-
dance, compared to the average of the sample.
8
The conspic-
uous blue RGB includes instead both Ba-rich and Ba-normal
stars. Is this separation only restricted to the Ba abundances?
Star 11413 is the only star on this sequence observed with
UVES. Unfortunately, this star was classified as binary by F93.
This would have no impact on the abundance analysis, unless
the binary system underwent mass transfer from an AGB com-
panion, a rare event. While this occurrence would explain the
kinematics and part of the chemical pattern of this star, surely
it is not the explanation of the secondary sequence: were the
stars on this sequence this kind of binaries, we would expect
a range in variations (also in colour), not a sequence. Anyway,
to be conservative, we cannot infer from this star alone that the
enhancement observed for Ba is extended to all the s−process
elements.
This pattern with the red sequence only populated by stars
with high Ba abundances is reminiscent of the chemical char-
acterization of the red RGB in NGC 1851 that are found to
be preferentially enriched in s−process elements (and in par-
ticular, Ba) in a number of recent studies (Yong and Grundahl
2008, Villanova et al. 2010, Carretta et al. 2011)
9
. The same
pattern, even more marked, is present in M 22, where however
stars on the bluer RGB are all s−poor and stars on the red RGB
are all s−rich, as demonstrated by Marino et al. (2009, 2011).
In that cluster, the separation is more neat.
A more detailed discussion of the properties of this “sec-
ondary RGB” in NGC 362 and a comparison with other pe-
culiar GCs will be presented in a separate paper (Carretta et
al. 2013b). Here we only note that the splitting in the y, v − y
diagram seems to be different from the split RGB foundin clus-
ters such as, e.g., NGC 104 (47 Tuc) or NGC 6752, when us-
ing colours including ultraviolet magnitudes, which are closely
related to the Na-O anti-correlation (see the extensive photo-
metric studies by Milone et al. 2012, 2013 and the theoretical
interpretation by Sbordone et al. 2011). Not all splits in the
different filter combinations have the same meaning (see the
discussion in Carretta et al. 2013b and Milone et al. 2012). No
“secondary RGB” is easily discernable in the y, v − y diagrams
for the two GCs just mentioned (at least, using the photometry
by Calamida et al. 2007) showing that this feature is present
in some, but not all GCs. These GCs also behave differently
for what concerns Ba abundances. To show this, we used the
Johnson UBVI photometry by Momany (unpublished photom-
etry, reduced as in Momany et al. 2003) for 47 Tuc and the
8
It also seems to show a tendency for lower O and higher Na abun-
dances, although we measured both elements only in six stars. This
will be discussed more in a companion paper.
9
The small sample by Yong and Grundahl is not assigned to one of
the RGBs; however, in Fig.1 of Lee et al. (2009) three s-rich stars are
seen to lie on the reddest branch.
Fig.23. Top: the plot shows 47 Tuc, using the same filter
combination adopted in Milone et al. (2012) to show a split
RGB (see their Figs. 21 and 26); the photometry is not the
same. Ba-rich and Ba-poor stars are not segregated. Bottom:
for NGC 6752 we show the same effect on a plot similar to the
one in Milone et al. (2013; their Fig. 15).
Str¨omgren photometry by Grundahl for NGC 6752 (Calamida
et al. 2007), to reproduce the clear splitsshown by Milone et al.
(2012,2013). We took the Ba abundances from D’Orazi et al.
(2010), cross-identified the stars, and plotted the ones in com-
mon in Fig. 23. When we divide the sample in Ba-rich or Ba-
poor stars with respect to the average [Ba/Fe] value in each
cluster (as done in NGC 1851 by Carretta et al. 2011 and in the
present work for NGC 362), no segregation is evident on the
split RGBs in these two clusters, at odds with what is seen for
C,N and Na,O (Milone et al. 2012, 2013).
In the last years we are learning new things on the nature
of GCs, which are more complex -and more interesting- than

14 E. Carretta et al.: Multiple stellar populations in NGC 362
we thought for long time. True progress and understanding re-
quires a multi-lateral approach,making the best of the informa-
tion that photometry, spectroscopy, and theoretical modelling
can give us.
Acknowledgements. VD is an ARC Super Science Fellow. This pub-
lication makes use of data products from the Two Micron All Sky
Survey, which is a joint project of the University of Massachusetts
and the Infrared Processing and Analysis Center/California Institute
of Technology, funded by the National Aeronautics and Space
Administration and the National Science Foundation. This research
has been funded by PRIN INAF 2011 ”Multiple populations in glob-
ular clusters: their role in the Galaxy assembly” (PI E. Carretta),
and PRIN MIUR 2010-2011, project “The Chemical and Dynamical
Evolution of the Milky Way and Local Group Galaxies” (PI F.
Matteucci) . This research has made use of the SIMBAD database,
operated at CDS, Strasbourg, France and of NASA’s Astrophysical
Data System.
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E. Carretta et al.: Multiple stellar populations in NGC 362 15
Appendix A: Error estimates
We refer the reader to the analogous Appendices in Carretta et
al. (2009a,b) for a detailed discussion of the procedure adopted
for error estimates. In the following we only provide the main
tables of sensitivities of abundance ratios to the adopted errors
in the atmospheric parameters and EWs and the final estimates
of internal and systematic errors for all species analysed from
UVES and GIRAFFE spectra of stars in NGC 362.
The sensitivities of derived abundances on the adopted at-
mospheric parameters were obtained by repeating our abun-
dance analysis by changing only one atmospheric parameter
each time for all stars in NGC 362 and taking the average value
of the slope change vs. abundance. This exercise was done sep-
arately for both UVES and GIRAFFE spectra.
We notice that when estimating the contribution to internal
errors due to EWs and v
t
, the values usually adopted (deter-
mined from the scatter of abundances from individual lines)
are overestimated, because regularities in the data are not taken
into account. These regularities are due to uncertainties in the
gf−values, unrecognisedblends with adjacent lines, not appro-
priate positioning of the continuum, etc. They show up in uni-
form deviations of individual lines from average abundances
for each star. By averaging over all stars the residuals of abun-
dances derivedfrom individuallines with respect to the average
value for each star, we estimated that some 36% of the total
variance in the Fe abundances from individual lines is due to
systematic offsets between different lines, which repeat from
star-to-star. For about 30% of the line, these offsets have trends
with temperature significant at about 2 σ level. However, we
found that the additional fraction of variance that can be ex-
plained by these trends is very small, and we can neglect it. We
conclude that when considering star-to-star variations (internal
errors, according to our denomination), the errors in EWs and
v
t
should be multiplied by 0.8.
The amount of the variations in the atmospheric param-
eters is shown in the first line of the headers in Table A.1,
and Table A.2, whereas the resulting response in abundance
changes of all elements (the sensitivities) are shownin columns
from 3 to 6 of these tables.

16 E. Carretta et al.: Multiple stellar populations in NGC 362
Table A.1. Sensitivities of abundance ratios to variations in the atmospheric parameters and to errors in the equivalent widths,
and errors in abundances for stars in NGC 362 observed with UVES
Element Average T
eff
logg [A/H] v
t
EWs Total Total
n. lines (K) (dex) (dex) kms
−1
(dex) Internal Systematic
Variation 50 0.20 0.10 0.10
Internal 5 0.04 0.05 0.04 0.087
Systematic 59 0.06 0.05 0.01
[Fe/H]i 81 +0.041 +0.016 +0.003 −0.033 +0.010 0.017 0.051
[Fe/H]ii 9 −0.051 +0.104 +0.034 −0.012 +0.029 0.041 0.072
[O/Fe]i 2 −0.029 +0.067 +0.032 +0.030 +0.062 0.066 0.068
[Na/Fe]i 3 +0.008 −0.048 −0.032 +0.019 +0.050 0.054 0.053
[Mg/Fe]i 3 −0.010 −0.015 −0.006 +0.018 +0.050 0.051 0.017
[Al/Fe]i 2 +0.003 −0.024 −0.010 +0.028 +0.062 0.063 0.052
[Si/Fe]i 9 −0.053 +0.023 +0.010 +0.025 +0.029 0.032 0.064
[Ca/Fe]i 16 +0.021 −0.030 −0.018 −0.016 +0.022 0.025 0.027
[Sc/Fe]ii 7 +0.042 −0.023 −0.003 −0.016 +0.033 0.034 0.052
[Ti/Fe]i 9 +0.054 −0.019 −0.018 −0.001 +0.029 0.031 0.065
[Ti/Fe]ii 1 +0.030 −0.019 −0.005 +0.004 +0.087 0.087 0.038
[V/Fe]i 9 +0.068 −0.015 −0.013 +0.001 +0.029 0.031 0.081
[Cr/Fe]i 2 +0.037 −0.020 −0.028 −0.010 +0.062 0.064 0.048
[Cr/Fe]ii 16 +0.014 −0.025 −0.020 +0.008 +0.022 0.025 0.026
[Mn/Fe]i 3 +0.021 −0.013 −0.009 −0.018 +0.050 0.051 0.027
[Co/Fe]i 5 −0.008 +0.002 +0.003 +0.024 +0.039 0.040 0.019
[Ni/Fe]i 29 −0.014 +0.016 +0.008 +0.014 +0.016 0.018 0.019
[Cu/Fe]i 1 +0.010 +0.019 −0.003 −0.017 +0.100 0.103 0.091
[Y/Fe]ii 2 +0.040 −0.037 −0.006 +0.008 +0.100 0.087 0.050
[Zr/Fe]ii 2 +0.041 −0.017 −0.005 +0.002 +0.040 0.087 0.067
[Ba/Fe]ii 3 +0.071 −0.030 −0.002 −0.088 +0.010 0.062 0.097
[La/Fe]ii 1 +0.051 −0.071 +0.000 −0.001 +0.100 0.087 0.073
[Ce/Fe]ii 2 +0.041 −0.017 −0.006 +0.002 +0.050 0.087 0.071
[Nd/Fe]ii 7 +0.065 −0.021 −0.001 +0.008 +0.042 0.087 0.065
[Eu/Fe]ii 2 +0.051 −0.037 −0.001 +0.001 +0.062 0.062 0.055
[Dy/Fe]ii 1 +0.031 −0.024 −0.004 +0.002 +0.050 0.062 0.097
Table A.2. Sensitivities of abundance ratios to variations in the atmospheric parameters and to errors in the equivalent widths,
and errors in abundances for stars in NGC 362 observed with GIRAFFE
Element Average T
eff
logg [A/H] v
t
EWs Total Total
n. lines (K) (dex) (dex) kms
−1
(dex) Internal Systematic
Variation 50 0.20 0.10 0.10
Internal 5 0.04 0.04 0.12 0.126
Systematic 59 0.06 0.06 0.01
[Fe/H]i 29 +0.052 +0.000 −0.005 −0.028 +0.023 0.041 0.062
[Fe/H]ii 2 −0.033 +0.091 +0.023 −0.011 +0.089 0.093 0.048
[O/Fe]i 1 −0.041 +0.081 +0.034 +0.031 +0.126 0.133 0.059
[Na/Fe]i 2 −0.011 −0.031 −0.012 +0.018 +0.089 0.092 0.031
[Mg/Fe]i 2 −0.018 −0.007 −0.002 +0.016 +0.089 0.091 0.022
[Si/Fe]i 6 −0.046 +0.025 +0.010 +0.027 +0.051 0.061 0.055
[Ca/Fe]i 5 +0.007 −0.024 −0.006 −0.010 +0.056 0.058 0.011
[Sc/Fe]ii 4 −0.053 +0.080 +0.030 +0.011 +0.063 0.068 0.067
[Ti/Fe]i 3 +0.023 −0.008 −0.008 +0.009 +0.073 0.074 0.027
[V/Fe]i 4 +0.037 −0.006 −0.008 +0.011 +0.063 0.065 0.044
[Cr/Fe]i 3 −0.029 −0.045 −0.043 −0.021 +0.073 0.080 0.037
[Co/Fe]i 1 +0.005 +0.011 +0.007 +0.022 +0.126 0.129 0.018
[Ni/Fe]i 6 −0.010 +0.015 +0.006 +0.020 +0.051 0.057 0.014
[Ba/Fe]ii 1 −0.038 +0.057 +0.037 −0.057 +0.126 0.145 0.055

E. Carretta et al.: Multiple stellar populations in NGC 362 17
Table 2. List and relevant information for target stars in NGC 362 The complete Table is available electronically only at CDS.
ID RA Dec B V K RV(Hel) Notes
986 1 03 20.037 -70 49 55.53 15.227 14.139 11.443 228.83 HR11,HR13
995 1 03 23.159 -70 49 54.98 15.446 14.548 12.013 222.77 HR13

18 E. Carretta et al.: Multiple stellar populations in NGC 362
Table 3. Adopted atmospheric parameters and derived iron abundances. The complete table is available electronically only at
CDS.
Star T
eff
log g [A/H] v
t
nr [Fe/H]i rms nr [Fe/Hii rms
(K) (dex) (dex) (km s
−1
) (dex) (dex)
986 4533 1.59 -1.18 2.00 52 −1.176 0.161 3 −1.141 0.164
995 4661 1.92 -1.26 1.09 28 −1.259 0.111 3 −1.154 0.028

E. Carretta et al.: Multiple stellar populations in NGC 362 19
Table 4. Abundances of proton-capture elements in stars of NGC 362. n is the number of lines used in the analysis. Upper limits
(limO,Al=0) and detections (=1) for O and Al are flagged.
star n [O/Fe] rms n [Na/Fe] rms n [Mg/Fe] rms n [Al/Fe] rms limO limAl
986 1 +0.16 4 +0.21 0.12 3 + 0.30 0.16 1
995 1 +0.17 2 −0.14 0.04 1

20 E. Carretta et al.: Multiple stellar populations in NGC 362
Table 5. Abundances of α-elements in stars of NGC 362. n is the number of lines used in the analysis.
star n [Si/Fe] rms n [Ca/Fe] rms n [Ti/Fe] i rms n [Ti/Fe] ii rms
986 10 + 0.25 0.15 5 +0.33 0.16 5 +0.20 0.11
995 1 +0.26 6 +0.32 0.23

E. Carretta et al.: Multiple stellar populations in NGC 362 21
Table 6. Abundances of Fe-peak elements in stars of NGC 362. n is the number of lines used in the analysis.
star n [Sc/Fe] ii rms n [V/Fe] rms n [Cr/Fe] i rms n [Mn/Fe] rms n [Co/Fe] rms n [Ni/Fe] rms n [Cu/Fe] rms
986 5 −0.10 0.10 6 −0.05 0.08 5 −0.03 0.14 1 −0.12 12 −0.07 0.11
995 2 −0.10 0.04 3 −0.07 0.10 1 −0.11 3 −0.16 0.09

22 E. Carretta et al.: Multiple stellar populations in NGC 362
Table 7. Abundances of n−capture elements in stars of NGC 362 with UVES spectra; n is the number of lines used in the analysis.
star n [Y/Fe] ii rms n [Zr/Fe] ii rms n [La/Fe] ii rms n [Ce/Fe] ii rms n [Nd/Fe] ii rms n [Eu/Fe] ii rms n [Dy/Fe] ii rms
1037 1 +0.00 0.14 2 +0.46 0.06 1 +0.30 0.16 2 +0.13 0.12 9 +0.28 0.16 1 +0.66 0.12 1 +0.75 0.17
12017 1 −0.13 0.14 2 +0.40 0.11 1 +0.17 0.16 2 −0.06 0.12 9 +0.20 0.16 1 +0.57 0.12 1 +0.58 0.17
- For species with only one transition, the scatter is estimated by summing in quadrature the best fit variation obtained by changing temperature, gravity, metal abundance and
microturbulent velocity of amounts corresponding to the errors in these parameters. To this, we added in quadrature a fit error of 0.05 dex.

E. Carretta et al.: Multiple stellar populations in NGC 362 23
Table 8. Abundances of Ba ii in stars of NGC 362. n is the number of lines used in the analysis.
star n [Ba/Fe]ii rms
986 1 +0.400
995 1 +0.447