11institutetext:  Universidade de São Paulo, IAG, Rua do Matão 1226, Cidade
Universitária, São Paulo 05508-900, Brazil 22institutetext:
Astrophysikalisches Institut Potsdam, An der Sternwarte 16, Potsdam, 14482,
Germany 33institutetext: Astrophysics Group, Keele University, Keele,
Staffordshire ST5 5BG, UK 44institutetext: Kavli Institute for the Physics and
Mathematice of the Universe (WPI), University of Tokyo, 5-1-5 Kashiwanoha,
Kashiwa 277-8583, Japan 55institutetext: INAF, Osservatorio Astronomico di
Trieste, Via G.B. Tiepolo 11, I-34143 Trieste, Italy 66institutetext: IFPU,
Istitute for the Fundamental Physics of the Universe, Via Beirut, 2, 34151,
Grignano, Trieste, Italy 77institutetext: E.A. Milne Centre for Astrophysics,
Department of Physics & Mathematics, University of Hull, HU6 7RX, UK
88institutetext: Konkoly Observatory, Research Centre for Astronomy and Earth
Sciences, Hungarian Academy of Sciences, Konkoly Thege Miklos ut 15-17, H-1121
Budapest, Hungary 99institutetext: NuGrid Collaboration, www.nugridstars.org
1010institutetext: Joint Institute for Nuclear Astrophysics - Center for the
Evolution of the Elements, USA 1111institutetext: Università di Padova,
Dipartimento di Astronomia, Vicolo dell’Osservatorio 2, I-35122 Padova, Italy
1212institutetext: INAF-Osservatorio Astronomico di Padova, Vicolo
dell’Osservatorio 5, I-35122 Padova, Italy 1313institutetext: Universidade
Estadual de Santa Cruz, Depto. de Ciências Exatas e Tecnológicas, Rodovia
Jorge Amado km 16, 45662-900, Ilhéus, Brazil 1414institutetext: Universidade
Federal do Rio de Janeiro, Instituto de Física, Av. Athos da Silveira Ramos,
149, 21941-972, Rio de Janeiro, Brazil 1515institutetext: Universidade Federal
do Rio Grande do Sul, Departamento de Astronomia, CP 15051, Porto Alegre
91501-970, Brazil 1616institutetext: South African Astronomical Observatory
(SAAO), P.O. Box 9. Observatory 7935, South Africa

# UVES analysis of red giants in the bulge globular cluster NGC 6522 ††thanks:
Observations collected at the European Southern Observatory, Paranal, Chile
(ESO), under programmes 088.D-0398A, and 097.D-0175 (PI: B. Barbuy),
071.B-0617, 73.B-0074 (PI: A. Renzini.).

B. Barbuy 11 E. Cantelli 11 L. Muniz 11 S.O. Souza 11 C. Chiappini 22 R.
Hirschi 3344 G. Cescutti 5566 M. Pignatari 7788991010 S. Ortolani 11111212 L.
Kerber 1313 F.F.S. Maia 1414 E. Bica 1515 E. Depagne 1616

(Received; accepted )

###### Abstract

Context.  NGC 6522 is a moderately metal-poor bulge globular cluster
([Fe/H]$\sim$$-$1.0), and it is a well-studied representative among a number
of moderately metal-poor blue horizontal branch clusters located in the bulge.
The NGC 6522 abundance pattern can give hints on the earliest chemical
enrichment in the central Galaxy.

Aims. The aim of this study is to derive abundances of the light elements C
and N; alpha elements O, Mg, Si, Ca, and Ti; odd-Z elements Na and Al;
neutron-capture elements Y, Zr, Ba, La, and Nd; and the r-process element Eu.
We verify if there are first- and second-generation stars: we find clear
evidence of Na-Al, Na-N, and Mg-Al correlations, while we cannot identify the
Na-O anti-correlation from our data.

Methods. High-resolution spectra of six red giants in the bulge globular
cluster NGC 6522 were obtained at the 8m VLT UT2-Kueyen telescope with both
the UVES and GIRAFFE spectrographs in FLAMES+UVES configuration. In light of
Gaia data, it turned out that two of them are non-members, but these were also
analysed. Spectroscopic parameters were derived through the excitation and
ionisation equilibrium of Fe I and Fe II lines from UVES spectra. The
abundances were obtained with spectrum synthesis. Comparisons of abundances
derived from UVES and GIRAFFE spectra were carried out.

Results. The present analysis combined with previous UVES results gives a mean
radial velocity of v${}^{\rm hel}_{\rm r}$ = $-$15.62$\pm$7.7 km s-1 and a
metallicity of [Fe/H] = $-$1.05$\pm$0.20 for NGC 6522. Mean abundances of
alpha elements for the present four member stars are enhanced with
[O/Fe]=+0.38, [Mg/Fe]=$\approx$+0.28, [Si/Fe]$\approx$+0.19, and
[Ca/Fe]$\approx$+0.13, together with the iron-peak element
[Ti/Fe]$\approx$+0.13, and the r-process element [Eu/Fe]=+0.40. The neutron-
capture elements Y, Zr, Ba, and La show enhancements in the +0.08 $<$ [Y/Fe]
$<$ +0.90, 0.11 $<$ [Zr/Fe] $<$ +0.50, 0.00 $<$ [Ba/Fe] $<$ +0.63, 0.00 $<$
[La/Fe] $<$ +0.45, and -0.10 $<$ [Nd/Fe] $<$ +0.70 ranges. We also discuss the
spread in heavy-element abundances.

###### Key Words.:

Galaxy: Bulge - Globular Clusters: individual: NGC 6522 - Stars: Abundances,
Atmospheres

††offprints: B. Barbuy

##  1 Introduction

The Galactic bulge formation was probably a result of early mergers and/or
dissipative collapse combined with a buckling bar, as suggested by the
excellent modern data now available in the innermost parts of the Galaxy
(Queiroz et al. 2020a,b, Rojas-Arriagada et al. 2020, Pérez-Villegas et al.
2020, Kunder et al. 2020, Savino et al. 2020 - see also review by Barbuy et
al. 2018a), and chemodynamical models (e.g. Fragkoudi et al. 2020, Debattista
et al. 2020, Baba & Kawata 2020). Within this context, globular clusters are
important tracers of the early formation of the Galactic bulge, assuming that
most of them were formed in situ. In particular, their abundance pattern could
give hints as to the early nucleosynthesis processes in the Galaxy.

The globular cluster NGC 6522 located in the Galactic bulge was identified by
Baade (1946) as having a type II stellar population given its colour-magnitude
diagram (CMD); hence, it falling into the Population II class defined in Baade
(1944). Despite such an early identification of this cluster, it has not been
widely studied since then.

NGC 6522 is an old globular cluster, with a moderate metallicity of
[Fe/H]$\sim$$-$1.0, and a blue horizontal branch. Several other such clusters
are present in the Galactic bulge, such as NGC 6558 (Rich et al. 1998, Barbuy
et al. 2007, 2018b), HP 1 (Barbuy et al. 2006, 2016) , AL 3 (Ortolani et al.
2006), Terzan 9 (Ernandes et al. 2019), and UKS 1 (Fernández-Trincado et al.
2020). Rossi et al. (2015) presented a study gathering these inner globular
clusters, that might represent the earliest stellar populations in the Galaxy.
Based on their derived proper motions, Rossi et al. (2015) concluded that the
inner bulge globular clusters have clearly lower transverse motions and
spatial velocities than halo clusters, and they appear to be trapped in the
bulge bar.

Pérez-Villegas et al. (2020) computed the orbits of 78 inner Galaxy globular
clusters, following the selections given in Bica (2016). It was found that
most of the confirmed bulge-population clusters are confined in the bar region
but are not supporting the bar structure. For each cluster, a set of 1000
initial conditions were generated, following the Monte Carlo technique and
taking into account the observational uncertainties. NGC 6522 has a 99.8%
probability of being a bulge member, and most of the orbits among the
different initial conditions do support the bar shape. A study of the origin
of Galactic globular clusters by Massari et al. (2019) also identified NGC
6522 as having been formed in the main bulge and not, for example, in the
Gaia-Enceladus that merged with the Galaxy about 10 Gyr ago. At the same time,
it was confirmed to be an old cluster by Kerber et al. (2018): 14.1, 14.2 Gyr
for [Fe/H]= $-$1.0, $-$1.15 from BaSTI isochrones (Pietrinferni et al. 2004,
2006) and 12.1, 12.4 Gyr for [Fe/H]=$-$1.0, $-$1.15 from Dartmouth isochrones
(Dotter et al. 2008). This old age indicates that NGC 6522 was formed 4 Gyr
prior relatively to the estimated age of the bar formation of 8$\pm$2 Gyr by
Buck et al. (2018) and $\sim$8 Gyr ago by Bovy et al. (2019). Therefore, the
fact that NGC 6522 follows the bar probably indicates that it was confined
within the bar when the latter formed.

Barbuy et al. (2009) analysed eight member stars based on FLAMES-GIRAFFE
(Pasquini et al. 2002) spectra, included in the survey by Zoccali et al.
(2008). Even with these low-resolution spectra (as compared with the UVES
spectra analysed later) it was possible to detect some enrichment in s-process
elements, which Chiappini et al. (2011a, hereafter C11) interpreted as a
possible signature of an early generation of fast rotating stars (the so-
called spinstars). This was based on the idea that, as the age-metallicity
relation in the bulge was expected to be steeper than in the halo, it would
already be possible to reach metallicities as large as [Fe/H] $\sim-$1 on a
very short timescale. It was then suggested that in the bulge, globular
clusters at [Fe/H] $\sim-$1.0 would already represent tracers of the earliest
chemical enrichment phases. Subsequently, we obtained UVES spectra for four of
the stars previously analysed in Barbuy et al. (2009) and C11. The re-analysis
of these stars based on higher resolution and higher signal-to-noise (S/N)
spectra, obtained with the UVES spectrograph at the Very Large Telescope was
then presented in Barbuy et al. (2014) where some enrichment of s-process
elements in the very old NGC 6522 cluster has been confirmed, although smaller
than what was suggested in the earlier low-resolution spectra. It was then
necessary to expand the sample to better constrain the nature of the stars
that have polluted this very old cluster. With this aim, during our first UVES
observations we also obtained parallel observations with the FLAMES-GIRAFFE
spectrograph, and new candidate members were identified. These, in turn, were
observed with UVES in a new run. In this work, we present results for six
stars in NGC 6522 obtained in 2016 with the FLAMES-UVES spectrograph (Dekker
et al. 2000). Our main aim is to study the abundance signatures of heavy
elements in the cluster.

Furthermore, NGC 6522 was recently shown by Kerber et al. (2018) to have at
least two stellar populations in proportions of 86% as second generation (2G)
and 14% as first generation (1G). For this reason, we inspected possible Na-O
anti-correlations and Na-Al, Na-N, and Mg-Al correlations (Gratton et al. 2012
and references therein) among the present sample stars together with stars
analysed in Barbuy et al. (2014).

Finally, we compare element abundance derivation from UVES and GIRAFFE spectra
for stars observed with both instruments (in their common wavelength region)
to check for further use of the lower resolution spectra. Additionally, in
Table 14 in the appendix we list stars identified to be candidate members of
the cluster, observed with FLAMES-GIRAFFE, selected from their radial
velocities together with Gaia collaboration (2018, 2021) proper motions.

The observations are described in Sect. 2. Photometric stellar parameters’
effective temperature and gravity are derived in Sect. 3. Spectroscopic
parameters are derived in Section 4, and abundance ratios are computed in
Sect. 5\. A discussion is presented in Sect. 6, and conclusions are drawn in
Sect. 7.

##  2 Observations

In Barbuy et al. (2009), eight stars of NGC 6522 observed with the GIRAFFE
spectrograph, within the list of over 600 bulge stars analysed by Zoccali et
al. (2008) (programs 71.B-0617 and 73.B-0074, PI: A. Renzini) were studied.
Four of these were re-observed at higher resolution with UVES and analysed in
Barbuy et al. (2014), in programme 88.D-0398 (PI: B. Barbuy) in 2012. From
GIRAFFE stars observed in the same field within this same programme, we
identified stars with radial velocities and metallicities that could be
cluster members. In program 097.D-0175 (PI: B. Barbuy), we observed five of
these stars, plus star B118, which was previously studied at a moderate
resolution in Barbuy et al. (2009). The log-book of observations is given in
Table 1.

The UVES spectra were obtained using the FLAMES setup centred at 580 nm, with
a coverage ranging from 480 nm to 680 nm. The 2012 GIRAFFE spectra were
obtained using setups HR11 (559.7-584.0 nm) and HR12 (582.1-614.6 nm), and the
2016 GIRAFFE observations used setups HR11, HR13 (612.0-640.5 nm), and HR14A
(630.8-670.1 nm), all with a mean resolving power of R $\approx$ 22,000.

The UVES data were reduced with the UVES pipeline v5.7.0 within the REFLEX
ambient. The extracted spectra were treated, normalised, rest-frame aligned,
and combined using the method described in Section 2.3 of Cantelli (2019).
Cosmic-ray removal was done by sigma rejection. Radial velocities were
measured using the upper wavelength chip of the UVES red arm, through the IRAF
task fxcor, using the Arcturus spectrum as template (Hinkle et al. 2000). The
measurements for each star and each exposure and the heliocentric radial
velocities computed through the IRAF task rvcorrect are reported in Table 2.

For completeness, and thanks to the Gaia collaboration (2018, 2021)
measurements, we were able to proceed to a more robust verification of
membership. In Table 3, we list the proper motions ($\mu_{\alpha}\cos\delta$
and $\mu_{\delta}$) and G magnitude from Gaia Early Data Release 3 (EDR3, Gaia
Collaboration 2020). With the precision improvement of Gaia EDR3 on proper
motions, we recalculated the mean values for NGC 6522. We selected stars
within 15 arcmin of the cluster centre and applied the Gaussian mixture models
(GMM) to separate the cluster stars from field stars. With this method, we
recalculated the cluster proper motion as $\mu_{\alpha}\cos\delta=2.55\pm
0.08$ mas yr-1 and $\mu_{\delta}=$-$6.45\pm 0.07$ mas yr-1. Also, the cluster
and field distributions allowed us to determine the membership probability for
each star. As shown in Table 3, and in the point-vector diagram of Figure 1,
stars 234816 and 244523 turn out to be non-members. This shows the power of
Gaia, given that these two stars have the very compatible metallicities and
radial velocities, coinciding with the other four member stars in the CMD.
This is even more striking given that only 0.5% of field stars show a
metallicity of [Fe/H]$\mathrel{\hbox{\raise 2.15277pt\hbox{$<$}\hbox
to0.0pt{\hss\lower 2.15277pt\hbox{$\sim$}}}}$$-$1.0 in the Galactic bulge
(Barbuy et al. 2018a).

A mean heliocentric radial velocity of v${}^{\rm hel}_{\rm r}$ = $-$16.96 km
s-1 is found for the four UVES sample member stars. A mean value of v${}^{\rm
hel}_{\rm r}$ = $-$14.3$\pm$3.3 km s-1 was obtained from UVES spectra of four
stars analysed in Barbuy et al. (2014). By combining the present data with
these four stars, namely, B-107, B-128, B-122 and B-130, with v${}^{\rm
hel}_{\rm r}$ = $-$7.626, $-$14.651, $-$18.043, and $-$16.808 km s-1,
respectively, we obtain a mean value of v${}^{\rm hel}_{\rm r}$ = -15.62 km
s-1. A range of velocities between v${}^{\rm hel}_{\rm r}$ = $-$7.63 and
$-$22.57 km s-1 gives a dispersion of $\pm$7.7 km s-1. A similar range of
radial velocities was detected in Fernández-Trincado et al. (2019), including
stars with $-$21.97 $<$ v${}^{\rm hel}_{\rm r}$ $<$ $-$6.61 km s-1.

The GIRAFFE data were retrieved from the ESO reduced
data111archive.eso.org/wdb/wdb/adp/phase3–main/form archive. The extracted
spectra belonging to the same setups were then corrected for radial velocity,
normalised, and combined by the median. In the appendix, we give a list of new
candidate member stars observed with GIRAFFE in 2012 and 2016, within a radial
velocity range of $-$14.5$\pm$12 km s-1, that are confirmed members from
proper motions.

Figure 1: Gaia proper motions at the location of NGC 6522. The background is the proper motion density plot within 8 arcmin of the cluster centre. Blue crosses are the member sample stars, and red crosses are the non-member ones. Table 1: Log of the spectroscopic observations of programs 088.D-0398(A) and 097.D-0175(A), carried out in 2011-2012 and 2016, respectively. The seeing and airmasses reported are the mean values along the exposures. Date | UT | Julian | exp | Air- | Seeing |
---|---|---|---|---|---|---
|  | date | (s) | mass |  (′′) |
Program 088.D-0398(A) |
2011-10-08 | 00:45:54 | 2455842.53187 | 2750 | 1.462 |  0.82′′ |
2011-10-08 | 01:34:37 | 2455842.56571 | 2750 | 1.853 |  1.29′′ |
2012-03-06 | 07:38:32 | 2455992.81843 | 2750 | 1.579 |  1.15′′ |
2012-03-06 | 08:28:44 | 2455992.85329 | 2750 | 1.260 |  0.93′′ |
2012-03-07 | 07:47:56 | 2455993.82495 | 2750 | 1.489 |  0.81′′ |
2012-03-07 | 08:39:16 | 2455993.86060 | 2750 | 1.270 |  0.73′′ |
2012-03-25 | 08:31:47 | 2456011.85541 | 2750 | 1.087 |  0.64′′ |
Program 097.D-0175(A) |
2016-05-17 | 07:22:18 | 2457525.80716 | 2400 | 1.007 |  0.40′′ |
2016-05-17 | 08:05:08 | 2457525.83690 | 2400 | 1.033 |  0.47′′ |
2016-05-17 | 08:52:35 | 2457525.86985 | 2400 | 1.099 |  0.47′′ |
2016-07-11 | 02:33:35 | 2457580.60666 | 2400 | 1.028 |  0.96′′ |
2016-07-21 | 03:27:16 | 2457590.64394 | 2400 | 1.016 |  0.51′′ |
2016-07-21 | 04:43:37 | 2457590.69696 | 2400 | 1.112 |  0.54′′ |
2016-07-21 | 06:33:32 | 2457590.75246 | 2400 | 1.373 |  0.54′′ |
2016-07-22 | 04:48:26 | 2457591.70031 | 2400 | 1.131 |  0.48′′ |
2016-07-22 | 05:40:15 | 2457591.73629 | 2400 | 1.288 |  0.45′′ |
2016-07-22 | 06:33:29 | 2457591.77326 | 2400 | 1.574 |  0.63′′ |
Table 2: Observed and heliocentric radial velocities of the UVES sample stars, in each of the 10 UVES exposures, and the mean heliocentric radial velocity. OGLE n∘ |  | 234816 | 244523 | 244819 | 256289 | B118 | 402370
---|---|---|---|---|---|---|---
Date | UT |  Observed and Heliocentric radial velocity (${\rm km.s^{-1}}$)
|  | vobs | vhel | vobs | vhel | vobs | vhel | vobs | vhel | vobs | vhel | vobs | vhel
17-05-2016 | 07:22:18.596 |  $-$31.90 |  $-$15.21 |  $-$29.59 |  $-$12.89 |  $-$39.18 |  $-$22.49 |  $-$31.56 |  $-$14.87 |  $-$36.48 |  $-$19.79 |  $-$27.20 |  $-$10.51
17-05-2016 | 08:05:08.353 |  $-$31.98 |  $-$15.36 |  $-$29.66 |  $-$13.95 | — | — |  $-$31.75 |  $-$15.13 |  $-$36.61 |  $-$19.99 |  $-$27.53 |  $-$10.92
17-05-2016 | 08:52:35.292 |  $-$31.67 |  $-$15.13 |  $-$29.39 |  $-$12.85 |  $-$39.15 |  $-$22.61 |  $-$31.17 |  $-$14.63 |  $-$36.36 |  $-$19.82 |  $-$27.20 |  $-$10.66
11-07-2016 | 02:33:35.361 |  $-$06.85 |  $-$15.93 |  $-$04.61 |  $-$13.68 |  $-$13.99 |  $-$23.07 |  $-$06.42 |  $-$15.49 |  $-$10.98 |  $-$20.05 |  $-$02.02 |  $-$11.09
21-07-2016 | 03:27:16.562 |  $-$01.53 |  $-$15.27 | 0.76 |  $-$12.99 |  $-$08.47 |  $-$22.22 | 01.48 |  $-$12.27 |  $-$05.48 |  $-$19.22 | 02.69 |  $-$11.06
21-07-2016 | 04:43:37.414 |  $-$01.36 |  $-$15.24 | 0.77 |  $-$13.11 |  $-$08.14 |  $-$22.03 |  $-$0.92 |  $-$14.80 |  $-$05.25 |  $-$19.13 | 02.70 |  $-$11.19
21-07-2016 | 06:33:32.375 |  $-$1.33 |  $-$15.33 | 0.55 |  $-$13.46 |  $-$8.21 |  $-$22.21 |  $-$0.81 |  $-$14.82 |  $-$5.11 |  $-$19.11 | 2.84 |  $-$11.18
22-07-2016 | 04:48:26.611 |  $-$01.04 |  $-$15.37 | 01.11 |  $-$13.22 |  $-$08.38 |  $-$22.71 |  $-$0.54 |  $-$14.87 |  $-$05.36 |  $-$19.69 | 03.78 |  $-$10.55
22-07-2016 | 05:40:15.273 |  $-$01.39 |  $-$15.80 | 0.26 |  $-$14.15 |  $-$08.26 |  $-$22.67 |  $-$0.79 |  $-$15.20 |  $-$5.18 |  $-$19.59 | 02.89 |  $-$11.52
22-07-2016 | 06:33:29.572 |  $-$0.88 |  $-$15.36 | 1.11 |  $-$13.37 |  $-$08.66 |  $-$23.14 |  $-$0.44 |  $-$14.92 |  $-$05.12 |  $-$19.60 | 3.48 |  $-$11.00
Mean vhel |  |  |  $-$15.40 |  |  $-$13.37 |  |  $-$22.57 |  |  $-$14.70 |  |  $-$19.60 |  |  $-$10.97
Table 3: Identifications and Gaia proper motions; Gaia G and Johnson B and V magnitudes and membership probability. Stars from Barbuy et al. (2009, 2014) are also included.  OGLE | Name |  pmRA |  pmDEC | $Gmag$ | $V$ | $B$ | $\mathcal{P}_{Member}$
---|---|---|---|---|---|---|---
| Present work |
234816 |  | $-6.1170$ | $-8.3550$ | $15.6950$ | $16.401$ | $17.781$ | 0
244523 |  | $-3.0400$ | $-6.9400$ | $15.3813$ | $15.988$ | $17.283$ | 0
244819 |  | $2.6380$ | $-6.6090$ | $15.6814$ | $16.306$ | $17.614$ | 100
256289 |  | $2.5830$ | $-7.1310$ | $15.3887$ | $15.918$ | $17.192$ | 98
402322 | B118 | $2.4800$ | $-6.5720$ | $15.3817$ | $16.011$ | $17.331$ | 100
402370 |  | $2.4900$ | $-6.6590$ | $15.5926$ | $16.226$ | $17.551$ | 100
Stars from Barbuy et al. (2014) |  |
402361 | B107 | $2.8460$ | $-6.8850$ | $15.3375$ | $15.980$ | $17.292$ | 100
244582 | B122 | $2.4820$ | $-6.0110$ | $15.3401$ | $16.000$ | $17.354$ | 99
402607 | B128 | $2.4810$ | $-6.4190$ | $15.6001$ | $16.260$ | $17.621$ | 100
402531 | B130 | $3.0100$ | $-6.1770$ | $15.6871$ | $16.300$ | $17.597$ | 100
| Other stars from Barbuy et al. (2009) |
412752 | B008 | $2.6170$ | $-6.4500$ | $15.4115$ | $15.990$ | $17.402$ | 100
– | B108 | $3.3340$ | $-6.2990$ | $15.1898$ | $16.290$ | $18.313$ | 99
– | B134 | $0.7350$ | $-5.0020$ | $15.5082$ | $16.040$ | $17.367$ | 0
244829 | F121 | $2.2580$ | $-7.0090$ | $15.7923$ | $16.400$ | $17.849$ | 99

##  3 Photometric stellar parameters

###  3.1 Temperatures

The selected stars, their OGLE and 2MASS designations, coordinates, and
$VIJHK_{\rm s}$ magnitudes are given in Table 4. $V$ and $I$ data were
collected from the Optical Gravitational Lensing Experiment (OGLE) survey, the
OGLE-II release222www.astrouw.edu.pl/$\sim$ogle/photdb, Field Bul-SC45
centered at 18:03:33.0, $-$30:05:00 from Udalski et al. (2002). 2MASS $J$,
$H$, and $K_{\rm s}$ are from Skrutskie et al. (2006)333
$\mathtt{http://ipac.caltech.edu/2mass/releases/allsky/}$;
$\mathtt{https://irsa.iapc.caltech.edu}$, and VVV $J$, $H$, and $K_{\rm s}$
magnitudes are from the Vista Variables in the Via Lactea survey (Saito et al.
2012)444$\mathtt{horus.roe.ac.uk/vsa}$. Table 3 reports the Gaia (2018, 2020)
G magnitudes and deduced B magnitudes by applying the transformation $G_{\rm
mag}-V_{C}=f(B-V)$ from Riello et al. (2021). In Fig. 2, we show the location
in $B$, $V$ of the sample stars, in the CMD of NGC 6522 from data observed in
F435W and F555W with the Hubble Space Telescope, and converted to $B$ and $V$
by Piotto et al. (2002).

Figure 2: NGC 6522 Colour-magnitude diagram $V$ versus $B-V$ by Piotto et al.
(2002) with the location of the observed stars. Present four stars: blue. Four
stars from Barbuy et al. (2014): cyan. Two non-member stars: orange.

Photometric effective temperatures and bolometric corrections were derived
from $V-I$, $V-K$, and $J-K$ using the colour-temperature calibrations of
Alonso et al. (1999, hereafter AAM99) and Casagrande et al. (2010, hereafter
C10). For the transformation of $V-I$ from the Cousins to Johnson system, we
used ($V-I$)C=0.778($V-I$)J (Bessell 1979). $J$,$H$, and $K_{S}$ 2MASS
magnitudes and colours were transformed from the 2MASS system to the
California Institute of Technology (CIT), and from this to TCS (Telescopio
Carlos Sánchez), following Carpenter (2001) and Alonso et al. (1998). The VVV
$JHK_{s}$ colours were transformed to the 2MASS $JHK_{s}$ system using
relations by Soto et al. (2013) and then transformed to CIT as above.

The derived photometric effective temperatures, which are adopted as initial
guesses, resulting from relations from AAM99 and C10 are both listed in Table
5. The differences in effective temperatures, in the $\Delta$Teff(C10-AAM99)
sense, are of +64.7 K, +54.2 K and $\pm$140 K for $V-I$, $V-K$ and $J-K,$
respectively. These temperatures are used only as a guide to start fitting
them from the Fe I and Fe II lines.

###  3.2 Gravities

For a derivation of photometric gravities, we used the classical formula,
adopting Teff,⊙=5770 K, M∗=0.85 M⊙, and Mbol,⊙ = 4.75. For the cluster, we
used a distance modulus of ($m-$M)0 = 14.40 and a reddening of E(B-V)=0.52 and
AV = 1.61, based on Kerber et al. (2018). Bolometric corrections were derived
using AAM99 and C10 assuming BC${}_{\rm V,\odot}=$-$0.07$, M${}_{\rm
V,\odot}=4.81$, and M${}_{\rm I,\odot}=4.10$ from Willmer (2018). The computed
bolometric magnitudes and gravities are given in Table 5.

Table 4: Identifications, coordinates, and magnitudes. $JHK_{s}$ from both 2MASS and VVV surveys.  OGLE |  ---2MASS ID | $\alpha_{2000}$ |  ---$\delta_{2000}$ |  ---$V$ |  ---$I$ |  ---$J$ |  ---$H$ |  ---$K_{\rm s}$ |  ---$J_{\rm VVV}$ |  ---$H_{\rm VVV}$ | $K_{\rm VVV}$ |
---|---|---|---|---|---|---|---|---|---|---|---|---
234816 | 18032652-3006385 | 18:03:26.52 |  $-$30:06:38.1 | 16.401 | 14.604 | 13.198 | 12.488 | 12.302 | 13.1735 | 12.5075 | 12.3052 |
244523 | 18032757-3003455 | 18:03:27.56 |  $-$30:03:45.1 | 15.988 | 14.325 | 12.667 | 11.872 | 12.274 | 13.0323 | 12.4093 | 12.233 |
244819 | 18033354-3002254 | 18:03:33.51 |  $-$30:02:25.2 | 16.306 | 14.672 | 13.020 | 12.284 | 11.421 | 13.3148 | 12.6694 | 12.499 |
256289 | — | 18:03:31.58 |  $-$30:00:51.0 | 15.918 | 14.337 | — | — | — | 13.0892 | 12.4721 | 12.289 |
402322 | 18034225-3003403 | 18:03:42.25 |  $-$30:03:40.0 | 16.011 | 14.313 | 13.056 | 12.305 | 12.142 | 12.9661 | 12.324 |  |
402370 | 18034235-3002088 | 18:03:42.35 |  $-$30:02:08.5 | 16.226 | 14.554 | 13.391 | 12.673 | 12.550 | 13.2643 | 12.628 |  |
Table 5: Photometric stellar parameters derived using the calibrations by Alonso et al. (1999) (AAM99) and Casagrande et al. (2010) (C10) for $V-I$, $V-K$, $J-K$, bolometric corrections, bolometric magnitudes and corresponding gravity log $g$. star |  T($V-I$) |  T($V-K$) |  T($J-K$) |  T($V-K$) |  T($J-K$) |  log(Teff) | ${\rm BC_{V}}$ | ${\rm M_{bol}}$ | log g | Calib
---|---|---|---|---|---|---|---|---|---|---
|  | 2MASS | 2MASS | VVV | VVV | (mean) |  |  |  |
| (K) | (K) | (K) | (K) | (K) |  |  |  |  |
234816 | $4517.1$ | $4504.8$ | $4693.7$ | $4501.0$ | $4579.4$ | $3.655$ | $-0.475$ | $0.86$ | $2.363$ | AAM99
| $4600.0$ | $4568.4$ | $4843.0$ | $4564.5$ | $4712.6$ | $3.668$ | $-0.477$ | $0.86$ | $2.414$ | C10
244523 | $4792.5$ | $4848.0$ | $7391.8$ | $4799.3$ | $4806.5$ | $3.681$ | $-0.344$ | $0.32$ | $2.249$ | AAM99
| $4854.4$ | $4912.6$ | $7372.1$ | $4865.1$ | $4969.3$ | $3.732$ | $-0.351$ | $0.32$ | $2.453$ | C10
244819 | $4858.8$ | $4027.3$ | $3340.5$ | $4749.2$ | $4746.3$ | $3.687$ | $-0.319$ | $0.62$ | $2.391$ | AAM99
| $4917.7$ | $4039.6$ | $3115.6$ | $4815.7$ | $4902.1$ | $3.639$ | $-0.325$ | $0.61$ | $2.199$ | C10
256289 | $4986.7$ | — | — | $4928.7$ | $4805.2$ | $3.698$ | $-0.277$ | $0.19$ | $2.264$ | AAM99
| $5042.1$ | — | — | $4990.7$ | $4967.9$ | $3.699$ | $-0.279$ | $0.18$ | $2.264$ | C10
402322 | $4715.8$ | $4698.8$ | $4638.5$ | $4681.9$ | $4686.9$ | $3.674$ | $-0.376$ | $0.37$ | $2.242$ | AAM99
| $4782.2$ | $4765.6$ | $4780.3$ | $4748.8$ | $4835.3$ | $3.680$ | $-0.383$ | $0.37$ | $2.266$ | C10
402370 | $4772.5$ | $4887.1$ | $4873.3$ | $4786.4$ | $4778.9$ | $3.679$ | $-0.352$ | $0.57$ | $2.341$ | AAM99
| $4835.4$ | $4950.6$ | $5043.1$ | $4852.4$ | $4938.6$ | $3.692$ | $-0.359$ | $0.56$ | $2.393$ | C10
Table 6: Final UVES spectroscopic parameters.  star |  -Teff |  -log g |  -[FeI/H] |  -[FeII/H] |  -[Fe/H] | ${\rm v_{t}}$
---|---|---|---|---|---|---
| (K) |  |  |  |  |  km s-1
234816 | 4650 | 2.25 |  $-$1.03 |  $-$1.08 |  $-$1.05 | 1.65
244523 | 4800 | 2.00 |  $-$1.10 |  $-$1.12 |  $-$1.11 | 2.30
244819 | 4690 | 2.30 |  $-$1.23 |  $-$1.19 |  $-$1.21 | 1.51
256289 | 4770 | 2.10 |  $-$1.12 |  $-$1.11 |  $-$1.11 | 1.25
B118 | 4820 | 2.20 |  $-$1.18 |  $-$1.16 |  $-$1.17 | 2.10
402370 | 4700 | 2.20 |  $-$1.15 |  $-$1.16 |  $-$1.15 | 1.15

##  4 Spectroscopic stellar parameters

The equivalent widths (EW) of Fe I and Fe II lines were measured using IRAF.
The EWs, together with wavelength (Å); excitation potential (eV), damping
constant C6 and oscillator strengths from VALD3 (Piskunov et al. 1995,
Ryabchikova et al. 2015), National Institute of Standards & Technology (NIST,
Martin et al. 2002)555http://physics.nist.gov/PhysRefData/ASD/lines-form.html
and Kurúcz
(1993)666http://www.cfa.harvard.edu/amp/ampdata/kurucz23/sekur.html,
777http://kurucz.harvard.edu/atoms.html and adopted values are given in Table
C. For Fe I, we chose NIST values when available, otherwise they are from
VALD3. In most cases, these coincide with the value from Kurucz, as can be
seen in Table C. For Fe II, values from Meléndez & Barbuy (2009) are used. The
solar abundances were adopted from Grevesse & Sauval (1998), including
$\epsilon$(Fe)=7.50 for Fe, except for A(O)=8.76 for oxygen from Steffen et
al. (2015).

The models were interpolated in the MARCS model atmospheres grid (Gustafsson
et al. 2008). We adopted the spherical and mildly CN-cycled set
([C/Fe]$=-0.13$, [N/Fe]$=+0.31$). These models consider [$\alpha$/Fe]=+0.20
for [Fe/H]=-0.50 and [$\alpha$/Fe]=+0.40 for [Fe/H]$\leq-$1.00. The LTE
abundance analysis and the spectrum synthesis calculations were performed
using the code described in Barbuy et al. (2018c). The code is an update of
the Meudon ABON2 code by M. Spite, continously improved along the years, which
adopts local thermodynamic equilibrium (LTE). In Trevisan et al. (2011), the
calculation of lines and in particular the continuum opacity calculation were
cross-checked with the code by the Uppsala group BSYN/EQWI (Edvardsson et al.
1993, and updates until that date). The basic atomic line list is from VALD3
(Ryabchikova et al. 2015). Molecular lines of CN (A2$\Pi$-X2$\Sigma$), C2 Swan
(A3$\Pi$-X3$\Pi$), TiO (A3$\Phi$-X3$\Delta$) $\gamma$, TiO
(B3$\Pi$-X3$\Delta$) $\gamma$’, TiO $\alpha$ C3$\Delta-$X3$\Delta,$ and TiO
$\beta$ c1$\Phi-$a1$\Pi$ systems are taken into account, as described in
Barbuy et al. (2018c).

We initially adopted the photometric effective temperature and gravity, and
then further constrained the temperature by imposing an excitation equilibrium
for Fe I lines and gravities by imposing ionisation equilibrium from lines of
Fe I and Fe II. Microturbulence velocities vt (km s-1) were determined by
cancelling the trend of Fe I abundance versus EW. Fits to the observed spectra
in regions containing the Fe II lines were carried out, as shown in Fig. 3 for
star 256289. The good match of the Fe II lines indicates that these lines
correspond to the equivalent widths measured, that they are not plagued by
blends, and that the stellar parameters are suitable. The final spectroscopic
parameters Teff, log g, [Fe I/H], [Fe II/H], [Fe/H] and vt values are reported
in Table 6. An example of excitation and ionisation equilibrium using Fe I and
Fe II lines is shown in Fig. 4 for star B118.

Figure 3: Fe II lines in star 256289. Figure 4: Excitation and ionisation
equilibrium in star B118.

##  5 Abundance ratios

Abundances ratios were obtained by means of line-by-line spectrum synthesis
calculations compared to the observed spectra. The abundance derivation
details are explained below, and the results are reported in Table 7 for C, N,
and O; and Table 8 for the $\alpha$-elements O, Mg, Si, Ca, Ti, odd-Z elements
Na, Al, and heavy elements Y, Zr, Ba, La, Nd, and Eu.

###  5.1 Carbon, nitrogen, and oxygen

Table 7: Carbon, nitrogen, and oxygen abundances derived from C2 (0,1), CN (5,1), and [OI] lines. line |  ---$\lambda$(Å) |  --234816 |  --244523 |  --244819 |  --256289 |  --B118 |  --402370
---|---|---|---|---|---|---|---
C2(0,1) |  -5635.5 |  -$\leq$+0.0 |  -$\leq$+0.2 |  -$\leq$+0.2 |  -$\leq$+0.2 |  -$\sim$+0.2 |  -$\leq$+0.0
CN(5,1) |  -6332.2 |  -$\leq$+0.8 |  -$\leq$+0.8 |  -$\leq$+1.2 |  -$\leq$+0.3 |  -$\leq$+0.8 |  -$\leq$+0.3
[OI] |  -6300.3 | +0.60 | +0.40 | +0.40 | +0.40 | — | +0.40

Table 7 gives the results for C, N, and O abundances. The carbon abundances
were estimated from the molecular C2(0,1) Swan bandhead at 5635.5 Å. These
bandheads are faint, and in these stars they allow us to give an upper limit
only. The atomic C I 5380.3 Å lines are essentially absent in these stars and
cannot be used. The nitrogen abundance is derived from the CN(5,1) red system
bandhead at 6332.2 Å. For the oxygen-forbidden line at [OI] 6300.311 Å, a
selection among the original spectra where telluric lines did not contaminate
the line was needed, since most of the observations were contaminated. A few
spectra were retrieved showing a clean [OI] 6300.311 Å line, and the oxygen
abundance was derived. Figure 5 shows fits to C2, CN and [OI] lines for star
244819.

Figure 5: C, N, and O lines in star 244819: Panel (a): C2 5635 Å computed with
([C/Fe], [N/Fe], [O/Fe]) = (0.0, 0.8, 0.4), (0.2, 0.8, 0.4) in red, (0.2, 1.2,
0.4) in green, (0.4, 1.2, 0.4) in blue. Panel (b): CN 6332 Å computed with:
(0.2, 0.8, 0.4) in red, (0.2, 1.2, 0.4) in green, (0.4, 0.8, 0.4) and (0.4,
1.2, 0.4) in blue. Panel (c): [OI] 6300 Å computed with (0.2, 0.8, 0.4) in
red, (0.2, 1.2, 0.4) in green, (0.2, 1.2, 0.5) in blue. In all cases, black
dotted lines are the observed spectra. Figure 6: Fits to Mg I lines in star
234816. Figure 7: Fits to Na I lines in star 234816.

###  5.2 Odd-Z and alpha elements

Line-by-line abundances of the odd-Z elements Na and Al and alpha elements Mg,
Si, Ca, and Ti are reported in Table 8. Ti is an iron-peak element, but given
its behaviour following the alpha-elements, it is often considered as an
alpha.

For star 234816, there is a clear overenhancement of the alpha elements as
well as of the r element Eu. Figure 6 shows the MgI lines studied showing
agreement for a high enhancement of [Mg/Fe]=+0.7. The odd-Z elements Na and Al
are also enhanced in this star, with [Na/Fe]=[Al/Fe]=+0.5, as illustrated for
Na lines in Figure 7.

We exhaustively remeasured equivalent widths using other tools than IRAF, and
we redetermined stellar parameters, and even with somewhat different stellar
parameters. Models of (Teff, log g, [Fe/H], vt) = (4530 K, 2.2, $-$1.04, 1.2
km.s-1) and (4440 K , 2.02, $-$0.78, 1.11 km.s-1) were employed in Cantelli
(2019), and the overenhancement in alpha elements persists.

###  5.3 Heavy elements

Line-by-line abundances of Y, Zr, Ba, La, Nd, and Eu are reported in Table 8.
In Table 9 the mean abundances are reported, including results from Barbuy et
al. (2014). Below we describe details on the lines of the heavy elements
studied.

Barium, Lanthanum and Europium: The hyperfine structure (HFS) for the studied
lines of Ba II 5853.675, 6141.713 and 6496.897 Å and Eu II 6645.064 Å were
taken into account. For Ba II 5853.675, we computed the splitting of lines by
employing a code made available by Andrew McWilliam (McWilliam et al. 2013).
For Ba II 6141.713 Å and Ba II 6496.897 Å lines, as well as for La II lines,
the HFS structure was reported in Barbuy et al. (2014), and for Eu II 6645.064
Å the HFS was adopted from Hill et al. (2002). For Ba II 5853.675 Å, the
magnetic dipole A-factor, and the electric quadrupole B-factor were adopted
from Biehl (1976) and Rutten (1978), as given in Table 15. The nuclear spin is
I=1.5 and the isotopic nuclides Ba138 and Ba137 Ba136, Ba135 and Ba134
contribute with 71.7% and 11.23%, 7.85%, 6.59%, 2.42%, respectively (Asplund
et al. 2009). The hyperfine splitting applies only to the odd-Z nuclides Ba135
and Ba137. The line list taking into account hyperfine structure for Ba II
5853.675 Å is given in Table 16.

Figures 8, 9, 10, 11, and 12 show, respectively, the fits to the Y I and Y II,
three Ba II, La II, Nd II lines, and Eu II studied lines in the six sample
stars.

Strontium: We carefully inspected the Sr lines and give the conclusions here.
The line Sr I 6503.989 Å is in the wing of another line and is very shallow;
Sr I 6791.016 Å is also very shallow, and only at higher S/N it could be used
(as we did in Barbuy et al. 2014). We here describe the blends contained in
the Sr I 6550.244 Å line in detail. In Figure 13, we show that: a) despite the
presence of several TiO lines from the $\gamma$, $\gamma$’, $\alpha,$ and
$\beta$ systems, they are faint, given that TiO only gets stronger is very
cool stars. To take the TiO lines into account, we adopted [O/Fe], [Ti/Fe]
from Table 9; b) There are lines of Mn II, Cr I, Ca I, Fe II, Sc II, Ni I, Si
I, and Tm II, but all of these are extremely faint and do not influence the
strength of the blend; c) The Nd II 6550.178 Å line is not very strong, but it
does contribute to the strength of the blend. We used three Nd II lines to
derive [Nd/Fe] for the sample stars, and the fits are presented in Fig. 11.
The resulting mean Nd abundance is then fixed in order to compute the Sr
abundance. In Fig. 13, we show the Nd line for the abundance derived and also
for a +0.1 or +0.2dex increase; d) The main contributors to a blend are C2 and
CN lines. Although the C and N abundances are upper limits, by adopting these
values the blend is strong. The computations were done for [C/Fe]=0.2 and
[N/Fe]=0.8 (that are the upper limits given in Table 5). Since C and N are
anticorrelated we compute also with [C/Fe]=0.1 and [N/Fe]=0.9 and [C/Fe]=0.3
and [N/Fe]=0.7: the strong variation in the C2 and CN line strengths make the
derivation of Sr abundance very uncertain. The identified lines are C2 Swan
system (v’,v”) = (2,5): R3(30) 6550.660, R2(31) 6550.398 and R1(32) 6550.296
and CN red system (v’,v”) = (6,2): Q1(22) 6550.269 Å; e) In conclusion, we
estimate values of [Sr/Fe]=+1.6 for star 244816 and [Sr/Fe]=+0.7 for star B118
but these cannot be considered reliable.

Figure 8: Fits to Y I 6435.004, and Y II 6795.414 Å in the six sample stars. Observed spectra (black dotted lines) are compared with synthetic spectra (red lines) computed for [Y/Fe]=$-$0.3,0,0.3,0.5,0.8,1.0. Figure 9: Fits to Ba II 5853.675, 6141.713 and 6496.897 Å in the six sample stars. Observed spectra (black dotted lines) are compared with synthetic spectra (red lines) computed for [Ba/Fe]=$-$0.3,0,0.3,0.5,0.8,1.0. Figure 10: Fits to La II 6262.287, 6320.376, and 6390.477 Å in the six sample stars. Observed spectra (black dotted lines) are compared with synthetic spectra (red lines) computed for [La/Fe]=$-$0.3,0,0.3,0.5,0.8,1.0. Figure 11: Fits to Nd II 6740.078, 6790.372 and 6549.525 Å in the six sample stars. Observed spectra (black dotted lines) are compared with synthetic spectra (red lines) computed for [Nd/Fe]=$-$0.3,0,0.3,0.5,0.8,1.0. Figure 12: Fits to Eu II 6437.650 and 6645.064 Å in the six sample stars. Observed spectra (black dotted lines) are compared with synthetic spectra (red lines) computed for [Eu/Fe]=$-$0.3,0,0.3,0.5,0.8,1.0. Figure 13:  Sr I 6550.244 Å line in stars 244819 and B118, showing the contribution of blends to the line. Observed spectra (black dotted lines) are compared to synthetic spectra for lines as indicated in the panels. Dashed lines correspond to the contribution of a particular element only. Full lines correspond to all lines included. The values [Sr/Fe]=+1.6$\pm$0.1 for 244819 and [Sr/Fe]=+0.7$\pm$0.1 for B118 are indicated by full green lines. These results are not considered reliable, however, due to blends.  Table 8: Abundances in the six UVES sample stars.  species |  $\lambda$ (Å) |  --$\chi_{ex}$ (eV) | log gf | 234816 | 244523 | 244819 | 256289 | B118 | 402370 |
---|---|---|---|---|---|---|---|---|---|---
NaI | 5682.633 | 2.10 |  $-$0.71 | +0.50 | +0.00 | +0.30 |  $-$0.40 | +0.30 |  $-$0.30 |
NaI | 5688.194 | 2.10 |  $-$1.40 | +0.50 |  $-$0.30 | +0.30 |  $-$0.30 | +0.30 |  $-$0.30 |
NaI | 5688.205 | 2.10 |  $-$0.45 | +0.50 |  $-$0.30 | +0.30 |  $-$0.30 | +0.30 |  $-$0.30 |
NaI | 6154.230 | 2.10 |  $-$1.56 | +0.50 | +0.15 | +0.50 |  $-$0.30 | +0.30 | +0.00: |
NaI | 6160.753 | 2.10 |  $-$1.26 | +0.50 | +0.25 | +0.30 |  $-$0.30 | +0.30 |  $-$0.30 |
AlI | 6696.185 | 4.02 |  $-$1.58 | +0.50 | +0.30 | +0.40 |  $-$0.30 | +0.60 |  $-$0.30 |
AlI | 6696.788 | 4.02 |  $-$1.42 | +0.50 | — | +0.40 | — | — |  $-$0.30 |
AlI | 6696.788 | 4.02 |  $-$2.72 | +0.50 | — | +0.40 | — | — |  $-$0.30 |
AlI | 6698.673 | 3.14 |  $-$1.65 | +0.50 | +0.30 | +0.40 |  $-$0.10 | +0.60 |  $-$0.30 |
MgI | 5528.405 | 4.34 |  $-$0.50 | +0.70 | +0.00 | +0.30 | +0.30 | +0.60 | +0.30 |
MgI | 6318.720 | 5.11 |  $-$2.10 | +0.70 | +0.45 | +0.30 | +0.30 | +0.30 | +0.30 |
MgI | 6319.242 | 5.11 |  $-$2.36 | +0.70 | +0.30 | +0.30 | +0.30 | +0.30 | — |
MgI | 6319.490 | 5.11 |  $-$2.80 | +0.70 | +0.30 | +0.30 | +0.30 | +0.30 | +0.30 |
MgI | 6765.450 | 5.75 |  $-$1.94 | +0.50 | — | — | — | — | +0.25 |
SiI | 5665.555 | 4.92 |  $-$2.04 | +0.50 | — | — | +0.20 | +0.20 | +0.20 |
SiI | 5666.690 | 5.62 |  $-$1.74 | +0.50 | — | — | — | +0.30 | +0.30 |
SiI | 5690.425 | 4.93 |  $-$1.87 | +0.50 | — | — | +0.25 | +0.30 | +0.20 |
SiI | 5948.545 | 5.08 |  $-$1.30 | +0.50 | +0.30 | +0.30 | +0.20 | +0.30 | +0.10 |
SiI | 6142.494 | 5.62 |  $-$1.50 | +0.50 | +0.50 | +0.30 | +0.30 | +0.30 | +0.30 |
SiI | 6145.020 | 5.61 |  $-$1.45 | +0.50 | +0.35 | +0.30 | +0.25 | +0.40 | +0.30 |
SiI | 6155.142 | 5.62 |  $-$0.85 | +0.50 | +0.35 | +0.30 | +0.00 | +0.30 | +0.15 |
SiI | 6237.328 | 5.61 |  $-$1.01 | +0.50 | +0.40 | — | +0.15 | +0.30 | +0.05 |
SiI | 6243.823 | 5.61 |  $-$1.30 | +0.50 | +0.40 | +0.20 | +0.15 | +0.25 | +0.30 |
SiI | 6414.987 | 5.87 |  $-$1.13 | +0.50 | +0.45 | — | +0.25 | +0.30 | +0.30 |
SiI | 6721.844 | 5.86 |  $-$1.17 | +0.50 | +0.45 | — | +0.15 | +0.55 | +0.30 |
CaI | 5601.277 | 2.53 |  $-$0.52 | +0.50 |  $-$0.30 |  $-$0.30 |  $-$0.15 | +0.00 | +0.12 |
CaI | 5867.562 | 2.93 |  $-$1.55 | +0.50 | +0.20 | +0.15 | +0.10 |  $-$0.05 |  $-$0.10 |
CaI | 6102.723 | 1.88 |  $-$0.79 | +0.30 |  $-$0.30 | +0.00 | +0.30 | +0.00 | +0.00 |
CaI | 6122.217 | 1.89 |  $-$0.20 | +0.30 |  $-$0.30 | +0.00 | +0.00 | +0.00 | +0.00 |
CaI | 6156.030 | 2.52 |  $-$2.39 | +0.40 | +0.30 | — | +0.05 | +0.00 | +0.25 |
CaI | 6161.295 | 2.51 |  $-$1.02 | +0.50 | +0.30 | +0.30 | +0.30 | +0.30 | +0.30 |
CaI | 6162.167 | 1.89 |  $-$0.09 | +0.40 |  $-$0.30 | +0.00 | +0.40 | +0.15 | +0.30 |
CaI | 6166.440 | 2.52 |  $-$0.90 | +0.50 | +0.15 | +0.15 | +0.00 | +0.10 | +0.00 |
CaI | 6169.044 | 2.52 |  $-$0.54 | +0.50 | +0.00 | +0.30 | +0.30 | +0.25 | +0.00 |
CaI | 6169.564 | 2.52 |  $-$0.27 | +0.50 |  $-$0.20 | +0.15 | +0.30 | +0.00 | +0.00 |
CaI | 6439.080 | 2.52 | +0.3 | +0.50 |  $-$0.30 | +0.00 | +0.40 |  $-$0.20 | +0.30 |
CaI | 6455.605 | 2.52 |  $-$1.35 | +0.60 | +0.30 | — | +0.30 | +0.40 | +0.15 |
CaI | 6464.679 | 2.52 |  $-$2.10 | +0.00 | — | — | +0.00 | — | — |
CaI | 6493.788 | 2.52 | -2.44 | +0.30 | +0.00 | +0.30 | +0.30 |  $-$0.20 | +0.10 |
CaI | 6499.654 | 2.52 |  $-$0.85 | +0.50 | +0.10 | +0.30 | +0.25 | +0.00 | +0.10 |
CaI | 6572.779 | 0.00 |  $-$4.32 | — | +0.30 | +0.30 | +0.00 | +0.30 | +0.00 |
CaI | 6717.687 | 2.71 |  $-$0.61 | — | +0.30 | +0.30 | … | +0.50 | +0.50 |
TiI | 5689.459 | 2.29 |  $-$0.44 | +0.50 | +0.00 | +0.30 | — | +0.30 | +0.00 |
TiI | 5866.449 | 1.07 |  $-$0.84 | +0.50 | +0.15 | +0.30 | +0.00 | +0.10 | +0.00 |
TiI | 5922.108 | 1.05 |  $-$1.46 | +0.50 | +0.20 | +0.15 | +0.00 | +0.30 | +0.00 |
TiI | 5941.750 | 1.05 |  $-$1.5 | +0.60 | +0.30 | +0.30 | +0.00 | +0.30 | +0.00 |
TiI | 5965.825 | 1.88 |  $-$0.42 | +0.50 | +0.30 | +0.10 | +0.00 | +0.30 | +0.10 |
TiI | 5978.539 | 1.87 |  $-$0.53 | +0.50 | +0.35 | +0.25 | +0.00 | +0.30 | +0.00 |
TiI | 6064.623 | 1.05 |  $-$1.94 | +0.60 | +0.30 | +0.10 | — | — | +0.00 |
TiI | 6091.169 | 2.27 |  $-$0.42 | +0.60 | +0.30 | +0.10 | +0.00 | +0.20 | +0.00 |
TiI | 6126.214 | 1.07 |  $-$1.43 | +0.55 | +0.30 | +0.15 | +0.10 | +0.30 | +0.00 |
TiI | 6258.110 | 1.44 |  $-$0.36 | +0.30 | +0.05 |  $-$0.30 |  $-$0.15 |  $-$0.30 | +0.15 |
TiI | 6261.106 | 1.43 |  $-$0.48 | +0.60 | +0.20 | +0.05 |  $-$0.10 | +0.00 | +0.00 |
TiI | 6266.010 | 1.75 |  $-$2.98 | — | — | — | — | — | — |
TiI | 6303.767 | 1.44 |  $-$1.57 | +0.50 | +0.30 | +0.25 | +0.00 | +0.30 | +0.00 |
TiI | 6312.240 | 1.46 |  $-$1.60 | +0.60 | +0.30 | +0.30 | +0.10 | +0.30 | +0.00 |
TiI | 6336.113 | 1.44 |  $-$1.74 | +0.55 | — | +0.30 | — | +0.20 |  $-$0.15 |
TiI | 6508.150 | 1.43 |  $-$2.05 | — | — | — | — | — | — |
TiI | 6554.238 | 1.44 |  $-$1.22 | +0.35 | +0.30 | +0.15 | — | +0.40 |  $-$0.30 |
TiI | 6556.077 | 1.46 |  $-$1.07 | +0.50 | +0.30 | +0.30 | +0.00 | +0.30 | +0.00 |
TiI | 6599.113 | 0.90 |  $-$2.09 | +0.60 | +0.40 | +0.30 | +0.00 | +0.35 | +0.00 |
TiI | 6743.127 | 0.90 |  $-$1.73 | +0.50 | +0.30 | +0.30 | +0.00 | +0.30 |  $-$0.10 |
TiII | 5336.771 | 1.58 |  $-$1.70 | +0.50 | +0.10 | +0.30 | +0.00 | — | +0.30 |
TiII | 5381.0212 | 1.57 |  $-$2.08 | +0.50 | +0.10 | +0.30 | +0.00 |  $-$0.30 | +0.00 |
TiII | 5418.751 | 1.58 |  $-$2.13 | +0.50 | +0.30 | +0.50 | +0.20 |  $-$0.20 | +0.10 |
TiII | 6491.580 | 2.06 |  $-$2.10 | +0.35 | +0.30 | +0.30 | +0.30 | +0.40 | +0.15 |
TiII | 6559.576 | 2.05 |  $-$2.35 | +0.35 | +0.30 | +0.30 | +0.30 | +0.30 | +0.30 |
TiII | 6606.970 | 2.06 |  $-$2.85 | +0.50 | +0.30 | — | +0.20 | +0.25 | +0.15 |
Y I | 6435.004 | 0.07 |  $-$0.82 | +0.50 | +0.80 | +0.80 |  $-$0.15 | +0.30 | +0.30 |
Y II | 6795.414 | 1.74 |  $-$1.19 | +0.30 | +0.50 | +1.00: | +0.30 | +0.40 | +0.00 |
Zr I | 6127.475 | 0.15 |  $-$1.06 | +0.50 | — | — | +0.15 | +0.50: | — |
Zr I | 6134.585 | 0.00 |  $-$1.426 | +0.65 | — | — | +0.10 | +0.50: | — |
Zr I | 6140.535 | 0.52 |  $-$1.6 | — | — | — | +0.10 | — | — |
Zr I | 6143.252 | 0.07 |  $-$1.1 | +0.65 | — | — | +0.10 | +0.50 | — |
Ba II | 5853.675 | 0.60 |  $-$1.1 | +0.50 | +0.65 | +0.30 | +0.30 | +0.15 | +0.60 |
Ba II | 6141.713 | 0.70 |  $-$0.08 | +0.50 | +0.60 | +0.30 | +0.30 | +0.00 | +0.60 |
Ba II | 6496.897 | 0.60 |  $-$0.32 | +0.60 | +0.65 | +0.50 | +0.50 | +0.30 | +0.80 |
La II | 6262.287 | 0.40 |  $-$1.60 | +0.40 | +1.00 | +0.40 | +0.00 | +0.60 | +0.00 |
La II | 6320.376 | 0.17 |  $-$1.56 | +0.30 | +0.80 | +0.30 | +0.00 | +0.50 | +0.00 |
La II | 6390.477 | 0.32 |  $-$1.41 | +0.40 | +0.90 | +0.40 | +0.00 | +0.60 | +0.10 |
Nd II | 6549.525 | 0.06 |  $-$2.01 | +0.30 | +0.60 | — | +0.00 | +0.30 | +0.00 |
Nd II | 6740.078 | 0.06 |  $-$1.53 | +0.00 | +0.80 | +0.30 | +0.00 | +0.30 | +0.00 |
Nd II | 6790.372 | 0.18 |  $-$1.77 | +0.30 | +0.80 | +0.00 | +0.00 | +0.30 | -0.30 |
Eu II | 6173.029 | 1.32 |  $-$0.86 | +0.80: | — | — | — | — | — |
Eu II | 6437.640 | 1.32 |  $-$0.32 | +0.80 | +0.80 | +0.50 | +0.40 | +0.50 | +0.50: |
Eu II | 6645.064 | 1.38 | +0.12 | +0.80 | +0.50 | +0.50 | +0.45 | +0.60 | +0.50 |
Table 9: Mean abundances of C, N, odd-Z elements Na, Al, $\alpha$-elements O, Mg, Si, Ca, Ti, and heavy elements Y, Zr, Ba, La, and Eu. The four stars analysed in Barbuy et al. (2014) are also included. For the mean values, only the eight member stars were considered. star | [C/Fe] | [N/Fe] | [Na/Fe] | [Al/Fe] | [O/Fe] | [Mg/Fe] | [Si/Fe] | [Ca/Fe] | [TiI/Fe] | [TiII/Fe] | [Y/Fe] | [Zr/Fe] | [Ba/Fe] | [La/Fe] | [Nd/Fe] | [Eu/Fe]
---|---|---|---|---|---|---|---|---|---|---|---|---|---|---|---|---
The two non-member stars |  |
234816 |  $<$0.2 | +0.80 | +0.50 | +0.50 | +0.60 | +0.66 | +0.50 | +0.42 | +0.52 | +0.45 | +0.40 | +0.60 | +0.53 | +0.36 | +0.20 | +0.80
244523 |  $<$0.2 |  $<$0.8 | -0.04 | +0.30 | +0.40 | +0.26 | +0.40 | +0.02 | +0.26 | +0.18 | +0.65 | — | +0.63 | +0.90 | +0.70 | +0.65
Four stars from the present work |  |
244819 |  $<$0.2 |  $<$1.2 | +0.34 | +0.40 | +0.40 | +0.30 | +0.28 | +0.14 | +0.19 | +0.34 | +0.90 | — | +0.36 | +0.36 | +0.15 | +0.50
256289 |  $<$0.2 |  $<$0.3 |  $-$0.32 |  $-$0.20 | +0.40 | +0.30 | +0.19 | +0.18 | +0.00 | +0.17 | +0.08 | +0.11 | +0.36 | +0.00 | +0.00 | +0.43
B118 |  $\mathrel{\hbox{\raise 2.15277pt\hbox{$<$}\hbox to0.0pt{\hss\lower 2.15277pt\hbox{$\sim$}}}}$0.2 |  $<$0.8 | +0.30 | +0.60 | — | +0.40 | +0.32 | +0.10 | +0.23 | +0.09 | +0.35 | +0.50 | +0.15 | +0.57 | +0.30 | +0.55
402370 |  $<$0.0 |  $<$0.3 |  $-$0.24 |  $-$0.30 | +0.40 | +0.29 | +0.23 | +0.13 |  $-$0.02 | +0.16 | +0.15 | — | +0.66 | +0.03 | -0.10 | +0.50
Four stars from Barbuy et al. (2014) |  |
B-107 | +0.00 | — | +0.03 | +0.28 | +0.50 | +0.33 | +0.17 | +0.16 | +0.03 | +0.17 | +0.32 | +0.20 | +0.45 | +0.20 | — | +0.40
B-122 |  $-$0.20 | +0.70 | +0.09 | +0.18 | +0.20 | +0.10 | +0.06 | +0.00 | +0.03 | +0.15 | +0.20 | +0.10 | +0.05 | +0.35 | — | +0.30
B-128 | +0.10 | +0.60 | +0.01 | +0.08 | +0.23 | +0.23 | +0.14 | +0.20 | +0.05 | +0.17 | +0.43 | +0.40 | +0.55 | +0.35 | — | +0.30
B-130 | +0.00 | +0.70 | +0.05 | +0.26 | +0.50 | +0.27 | +0.13 | +0.15 | +0.03 | +0.18 | +0.23 | +0.00 | +0.22 | +0.00 | — | +0.20
Mean |  $\sim$+0.14 | +0.66 | +0.09 | +0.24 | +0.38 | +0.28 | +0.19 | +0.13 | +0.07 | +0.18 | +0.33 | +0.23 | +0.35 | +0.23 | +0.09 | +0.40
Table 10: Heavy element abundance ratios. star | [Ba/La] | [Ba/Eu] | [La/Eu] | [Y/Ba] | [Y/La] |
---|---|---|---|---|---|---
234816 | +0.17 |  $-$0.27 |  $-$0.44 |  $-$0.13 | +0.04 |
244523 |  $-$0.27 |  $-$0.02 | +0.25 | +0.02 |  $-$0.25 |
244819 | +0.00 |  $-$0.14 |  $-$0.14 | +0.54 | +0.54 |
256289 | +0.36 |  $-$0.07 |  $-$0.43 |  $-$0.28 | +0.08 |
B118 |  $-$0.42 |  $-$0.40 | +0.02 | +0.20 |  $-$0.22 |
402370 | +0.63 | +0.16 |  $-$0.47 | -0.51 | +0.12 |
B-107 | +0.25 | +0.05 |  $-$0.20 |  $-$0.13 | +0.12 |
B-122 |  $-$0.30 |  $-$0.25 | +0.05 | +0.15 |  $-$0.15 |
B-128 | +0.20 | +0.25 | +0.05 |  $-$0.12 | +0.08 |
B-130 | +0.22 | +0.02 |  $-$0.20 | +0.01 | +0.23 |

###  5.4 Errors

Uncertainties in spectroscopic parameters are given in Table 12 for star NGC
6522: 402370. For each stellar parameter, we adopted the usual uncertainties
as for similar samples (Barbuy et al. 2014, 2016, 2018b): $\pm$100 K in
effective temperature, $\pm$0.2 on gravity, and $\pm$0.2 km s-1 on the
microturbulence velocity. Errors were computed by employing models with these
modified parameters, with changes of $\Delta$T${\rm eff}$=+100 K, $\Delta$log
g =+0.2, $\Delta$vt = 0.2 km s-1, and recomputing lines of different elements.
The error given is the abundance difference needed to reach the adopted
abundances. Uncertainties due to non-LTE effects are negligible for these
stellar parameters as discussed in Ernandes et al. (2018). The same error
analysis and estimations can be applied to other stars in our sample. A more
careful discussion is required for Ba. The heavy element abundances for star
B118 reported in Table 11 show that the abundance ratios are confirmed from
one work to another, except for Ba. This is due to the use of strong lines
that fall in the saturated part of the curve of growth, where the abundance is
a function of the square root of the number of atoms; the bottom of the lines
reaches a maximum, and the increase of abundance causes an increase in the
line wings. Therefore, abundance derivations from strong lines are in general
avoided, since they are too sensitive to stellar parameters and spectral
resolution. The La lines are, on the other hand, faint and they are at least
not affected by the same problem.

Finally, it is important to note that the main uncertainties in stellar
parameters are due to uncertainties in the effective temperature, as can be
seen in Table 5. The second most important source of error are the EWs, given
the limited S/N of the spectra, which can be estimated using the formula from
Cayrel (1988): $\sigma_{EW}$ = 1.5 $\sqrt{FWHM.\delta_{x}}/(S/N)$, where
$\delta_{x}$ is the pixel size. The difference in the mean metallicities
between the present work and Barbuy et al. (2014) are probably due to a
difference in the measurements of EWs, and in particular in the placement of
continua.

Table 11: Heavy element results for star B118 from Barbuy et al. (2009, B09), Chiappini et al. (2011, C11), Ness et al. (2014, N14) and present work. work | Teff | log g | [Fe/H] | vrmt | [Y/Fe] | [Zr/Fe] | [Sr/Fe] | [Ba/Fe] | [La/Fe] | [Eu/Fe] |
---|---|---|---|---|---|---|---|---|---|---|---
B09 | 4700 | 2.6 |  $-$0.84 | 1.30 | — | — | — | +1.00 | +0.50 | +0.50 |
C11 | same | same | same | same | +0.50 | — | +1.50 | same | same | same |
N14 | 5000 | 2.25 |  $-$1.04 | 2.45 | +0.30 | — | — | +0.30 | +0.55 | +0.40 |
this | 4820 | 2.20 |  $-$1.17 | 2.10 | +0.25 | +0.50 | +0.70 | +0.15 | +0.57 | +0.55 |

Comparison between results from UVES and GIRAFFE spectra: For the sample
stars, we have the UVES spectra ranging from 4800-5800 Å, with a gap at
5777-5824 Å, and the GIRAFFE spectra in the setups H11 (5597-5840) and H12
(5821-6146) only. Therefore, since most lines used for the stellar parameter
analysis are in the UVES red arm, and most lines for deriving abundances are
also in the UVES red arm, we cannot compare the stellar parameter derivation.
Moreover, we only compare the main lines in common between the two sets of
spectra, which are located at $\lambda$$<$6142 Å.

By comparing the abundances for a list of lines in common between UVES and
GIRAFFE, we give another indicator of uncertainty. In Table 13, we compare the
abundances derived from UVES spectra to those derived from GIRAFFE spectra.
The results show an excellent agreement. In order to be clear, the stellar
parameters are much better derived from UVES spectra, in particular because of
the measurement of the Fe II lines; whereas, given a set of stellar
parameters, the abundances from the same lines at the different resolutions
are both reliable.

Table 12: Abundance uncertainties for star N6522:402370, for uncertainties of $\Delta$Teff = 100 K, $\Delta$log g = 0.2, $\Delta$vt = 0.2 km s-1 and corresponding total error. The errors are to be added to reach the reported abundances.  Element | $\Delta$T | $\Delta$log $g$ |  -$\Delta$vt |  -($\sum$x2)1/2
---|---|---|---|---
| 100 K | 0.2 dex | 0.2 kms-1 |
(1) | (2) | (3) | (4) | (5)
[FeI/H] |  $-$0.10 | +0.01 | +0.05 |  +0.11
[FeII/H] | +0.10 |  $-$0.07 | +0.04 |  +0.13
[C/Fe] | +0.02 | +0.02 |  +0.00 |  +0.03
[O/Fe] | +0.00 | +0.05 | +0.00 |  +0.05
[NaI/Fe] | +0.05 | +0.00 | +0.00 |  +0.05
[AlI/Fe] | +0.06 | +0.00 | +0.00 |  +0.06
[MgI/Fe] | +0.00 | +0.01 | +0.00 |  +0.01
[SiI/Fe] | +0.03 | +0.00 | +0.00 |  +0.03
[CaI/Fe] | +0.08 | +0.00 | +0.01 |  +0.08
[TiI/Fe] | +0.12 | +0.01 | +0.00 |  +0.12
[TiII/Fe] |  $-$0.05 | +0.07 | +0.00 |  +0.09
[YI/Fe] | +0.15 | +0.04 | +0.00 |  +0.10
[YII/Fe] | +0.20 | +0.15 |  +0.00 |  +0.25
[ZrI/Fe] | +0.20 |  $-$0.01 |  +0.00 |  +0.20
[BaII/Fe] | +0.10 | +0.15 |  $-$0.15 |  +0.23
[LaII/Fe] | +0.05 | +0.15 |  +0.00 |  +0.16
[LaII/Fe] | +0.12 | +0.05 |  +0.00 |  +0.13
[EuII/Fe] |  $-$0.05 | +0.05 | +0.00 |  +0.07
Table 13: Comparison of abundances from UVES and GIRAFFE.  species |  $\lambda$ (Å) | 234816 | 234816 | 244523 | 244523 | 244819 | 244819 | 256289 | 256289 | B118 | B118 | 402370 | 402370 |
---|---|---|---|---|---|---|---|---|---|---|---|---|---|---
|  | UVES | GIRAFFE | UVES | GIRAFFE | UVES | GIRAFFE | UVES | GIRAFFE | UVES | GIRAFFE | UVES | GIRAFFE |
Na I | 5682.633 | +0.50 | +0.50 | +0.00 | +0.00 | +0.30 | +0.30 | -0.40 |  $-$0.30 | +0.30 | +0.40 |  $-$0.30 | +0.00 |
Na I | 5688.194 | +0.50 | +0.50 |  $-$0.30 | +0.00 | +0.30 | +0.30 |  $-$0.30 |  $-$0.30 | +0.30 | +0.50 |  $-$0.30 | +0.00 |
Na I | 5688.205 | +0.50 | +0.50 |  $-$0.30 | +0.00 | +0.30 | +0.30 |  $-$0.30 |  $-$0.30 | +0.30 | +0.50 |  $-$0.30 | +0.00 |
Si I | 5665.555 | +0.50 | +0.50 | — | +0.35 | — | +0.35 | +0.20 | +0.00 | +0.20 | +0.35 | +0.20 | +0.00 |
Si I | 5666.690 | +0.50 | — | — | — | — | — | — | +0.00 | +0.30 | +0.00 | +0.30 | +0.30 |
Si I | 5690.425 | +0.50 | +0.50 | — | +0.50 | — | +0.35 | +0.25 | +0.20 | +0.30 | +0.30 | +0.20 | — |
Si I | 5948.545 | +0.50 | +0.00 | +0.30 | +0.15 | +0.30 | +0.30 | +0.20 | +0.00 | +0.30 | +0.15 | +0.10 | +0.00 |
Si I | 6142.494 | +0.50 | +0.50 | +0.50 | — | +0.30 | — | +0.30 | — | +0.30 | +0.15 | +0.30 | +0.30 |
Ca I | 5601.277 | +0.50 | +0.50 |  $-$0.30 | +0.10 |  $-$0.30 | +0.30 |  $-$0.15 | +0.15 | +0.00 | +0.30 | +0.12 | +0.00 |
Ca I | 5867.562 | +0.50 | +0.50 | +0.20 | +0.30 | +0.15 | +0.30 | +0.10 | +0.10 |  $-$0.05 | +0.30 |  $-$0.10 | +0.00 |
Ca I | 6102.723 | +0.30 | +0.30 |  $-$0.30 | +0.00 | +0.00 | +030 | +0.30 | +0.30 | +0.00 | +0.15 | +0.00 | +0.15 |
Ca I | 6122.217 | +0.30 | +0.30 |  $-$0.30 | +0.00 | +0.00 | +0.30 | +0.00 | +0.15 | +0.00 | +0.00 | +0.00 | +0.00 |
Ti I | 5689.459 | +0.50 | +0.50 | +0.00 | +0.30 | +0.30 | +0.30 | — | — | +0.30 | +0.30 | +0.00 | — |
Ti I | 5866.449 | +0.50 | +0.50 | +0.15 | +0.20 | +0.30 | +0.30 | +0.00 | +0.00 | +0.10 | +0.15 | +0.00 | +0.00 |
Ti I | 5922.108 | +0.50 | +0.50 | +0.20 | +0.45 | +0.15 | +0.30 | +0.00 | +0.00 | +0.30 | +0.30 | +0.00 | +0.00 |
Ti I | 5941.750 | +0.60 | +0.50 | +0.30 | +0.50 | +0.30 | +0.30 | +0.00 | +0.20 | +0.30 | — | +0.00 | +0.00 |
Ti I | 5965.825 | +0.50 | +0.50 | +0.30 | +0.30 | +0.10 | +0.30 | +0.00 | +0.00 | +0.30 | +0.30 | +0.10 | +0.00 |
Ti I | 5978.539 | +0.50 | +0.50 | +0.35 | +0.30 | +0.25 | +0.30 | +0.00 | +0.00 | +0.30 | +0.30 | +0.00 | +0.00 |
Ti I | 6064.623 | +0.60 | +0.50 | +0.30 | +0.30 | +0.10 | +0.30 | — | +0.00 | — | +0.15 | +0.00 | +0.00 |
Ti I | 6091.169 | +0.60 | +0.50 | +0.30 | +0.30 | +0.10 | +0.30 | +0.00 | — | +0.20 | +0.30 | +0.00 | +0.00 |
Ti I | 6126.214 | +0.55 | +0.50 | +0.30 | +0.30 | +0.15 | +0.30 | +0.10 | +0.00 | +0.30 | +0.30 | +0.00 | +0.00 |
Zr I | 6127.475 | +0.50 | +0.50 | — | +0.30 | — | — | +0.15 | +0.30 | +0.50: | +0.50: | — | — |
Zr I | 6134.585 | +0.65 | +0.50 | — | +0.50 | — | — | +0.10 | — | +0.50: | — | — | — |
Zr I | 6140.535 | — | — | — | — | — | — | +0.10 | — | — | — | — | — |
Zr I | 6143.252 | +0.65 | +0.50 | — | — | — | — | +0.10 | — | +0.50 | +0.30 | — | — |
Ba II | 5853.675 | +0.50 | +0.50 | +0.50 | +0.60 | +0.30 | +0.30 | +0.30 | +0.30 | +0.00 | +0.20 | +0.60 | +0.30 |
Ba II | 6141.713 | +0.50 | +0.40 | +0.40 | — | +0.30 | — | +0.30 | — |  $-$0.30 |  $-$0.30 | +0.60 | +0.15 |

##  6 Discussion

The inspection on abundances of heavy elements in NGC 6522 was triggered by
the variation in Ba abundances reported in Barbuy et al. (2009). In C11, we
tentatively tried to connect these abundances with the s-process
nucleosynthesis calculations in spinstars first presented by Pignatari et al.
(2008) and computed by Frischknecht et al. (2012). We did so for a spinstar of
40 M⊙, a metallicity of [Fe/H] = $-$3.8, and a rotational velocity of Vrot =
500 km.s-1. The lower resolution data (R$\sim$22,000) from GIRAFFE spectra
were compatible with the s-process yields of spinstars boosted by up to four
orders of magnitude with respect to a non-rotating star of the same mass and
metallicity (see their Fig. 2).

A next step was presented in Barbuy et al. (2014), where the analysis of
higher resolution data from UVES (R$\sim$45,000), was studied in terms of an
extended grid of spinstar models from Frischknecht et al. (2016). This paper
reported enhancement of the heavy elements Sr, Y, La, and Ba measurable in
stellar spectra. The new results were shown to be compatible with expectations
from massive spinstars, but allowing for other mechanisms to be invoked.

In a third step using the observations of 2012, besides the four stars
observed with UVES and reported in Barbuy et al. (2014), we also identified
possible new cluster member stars from the GIRAFFE spectra. In 2016, we
obtained new UVES observations of six such newly identified member stars,
which we analyse here.

###  6.1 Analysis of the present results

We derived a mean metallicity of [Fe/H]=$-$1.16$\pm$0.05, somewhat lower than
the B09, B14, and Fernández-Trincado et al. (2019) values of [Fe/H]= -1.0,
-0.95 and -1.04, respectively, and closer to the value of [Fe/H] = -1.15 from
Ness et al. (2014). By gathering the metallicities of the present four member
stars and the other four stars from Barbuy et al. (2014), we obtain
[Fe/H]=$-$1.05$\pm$0.20.

The mean abundances for the six sample stars, as well as the four stars
studied in Barbuy et al. (2014), are reported in Table 9. In the mean we see a
normal expected enhancement of the bona fide alpha elements O and Mg, and a
mild enhancement of Si and Ca (and Ti, noting that Ti behaves as an alpha, but
it is an iron-peak element). There has been evidence that the alpha elements O
and Mg formed during the hydrostatic phases of massive stars nucleosynthesis
are more enhanced than the other alpha elements, Si, Ca, and Ti, which are
formed predominantly during explosive nucleosynthesis (Woosley & Weaver 1995,
McWilliam 2016).

The enhancements of N, Na, and Al vary from star to star, indicating that some
stars are probably second generation ones. We note that the enhancement in N
is also due to a large scatter in its early enrichment history (e.g. Cescutti
& Chiappini 2010) as well as stellar evolution effects. In Fig. 14, we do not
find a clear anti-correlation between the [O/Fe] and [Na/Fe] ratios in our
stellar sample, such as, for instance, in NGC 6121, which is a well-populated
cluster in this diagram (Carretta et al. 2009). In Fig. 14,we also show the
correlated abundance signatures of [Al/Fe] versus [Na/Fe], [N/Fe] versus
[Na/Fe], and [Mg/Fe] versus [Al/Fe]. These diagrams confirm the presence of at
least two stellar populations in NGC 6522, as found by Kerber et al. (2018)
from photometry.

For the heavy elements, we present the plot of abundances including the new
results. In Fig. 15, we compare the present results on heavy element
abundances of Sr, Y, Zr, La, Ba, and Eu relative to Fe, together with those
from Barbuy et al. (2014) for another four stars in NGC 6522. Literature
abundances from field bulge red giants are from Johnson et al. (2012), for Zr,
La, and Eu in Plaut’s field, Siqueira-Mello et al. (2016), van der Swaelmen et
al. (2016), and metal-poor giants from Howes et al. (2016) and Lamb et al.
(2017). Also included are the abundances Bensby et al. (2017 and references
therein) for 39 microlensed bulge dwarfs and subgiants that are older than 9.5
Gyr, selected among their 90 stars. Finally, abundances in the bulge globular
cluster HP 1 (Barbuy et al. 2016), NGC 6558 (Barbuy et al. 2018b), and M62
(Yong et al. 2014) are also shown.

From Fig. 15, the most striking feature is the abundance variation of Sr, Y,
Zr, and to lesser extent Ba and La at the metallicity of [Fe/H] $\sim$ $-$1.0,
where the bulge globular clusters are found. For Sr, we report literature data
only, given the unreliability of Sr derivation in the present sample due to
blends with CN and C2 lines. We note that the spread is clearly larger in halo
metal-poor stars with [Fe/H]$\mathrel{\hbox{\raise 2.15277pt\hbox{$<$}\hbox
to0.0pt{\hss\lower 2.15277pt\hbox{$\sim$}}}}$-2.5 (Cescutti et al. (2013),
Hansen et al. (2014). For Eu, the behaviour of [Eu/Fe] versus [Fe/H] is more
well-defined, indicating a spread at low metallicities and a declining
abundance ratio with increasing metallicities.

From Table 9 and Fig. 15, we find that a) Y tends to be enhanced, showing
strong star-to-star variations; b) Ba tends to be enhanced, showing star-to-
star variations. and c) Eu is enhanced similarly to the alpha elements O and
Mg.

As mentioned, Y and Ba variations are compatible with a large number of
nucleosynthesis processes, with the member stars 244819 and B118 showing
[Y/Ba] excesses of +0.54 and +0.35, respectively. The observation of more
heavy elements would be necessary to differentiate between the potential
astrophysical sources. The variation in [Y/Ba] data compared with a chemical
evolution model for the Galactic bulge is shown in Figure 16.

We note that in C11 there was no model, but instead a calculation for only one
mass and that was showing the impact of rotation already. The argument was
that even if that calculation was done for a very metal-poor metallicity,
because in the bulge the metallicity grows quickly, we would see its effect in
the very old bulge stars at [Fe/H] $-$1.0 as well.

In Fig. 16, we present the result of stochastic models, as presented in
Cescutti et al. (2018). This can be summarised as follows. The nucleosynthesis
adopted for the s process from rotating massive stars comes from Frischknecht
et al. (2016). In this set of yields, the s process for massive stars is
computed for a rotation velocity of vini/vcrit = 0.4 and is composed of a grid
of four stellar masses (15, 20, 25, and 40 M⊙) and three metallicities (solar
metallicity, 10-3, 10-5) (Cescutti & Chiappini 2014, Cescutti et al. 2013).
The model considers the enrichment produced by r-process events as originated
from magneto-rotationally driven supernovae (MRD SNe; see Winteler et al 2012,
Nishimura et al. 2017); MRD SNe are assumed to be 10% of all the SNe II. The
model also takes into account the s-process production from 1.5 to 3 M⊙ stars
and SNIa enrichment, as in Cescutti et al. (2006). In summary, this model
considers the fact that the enrichment in heavy elements takes place both in
spinstars and in MRD supernovae.

A spread in abundances of these elements is observed in metal-poor halo stars
(e.g. François et al. 2003; Cescutti & Chiappini 2014; Rizzutti et al. 2021)
and is expected from spinstar models (Frischknecht et al. 2012, 2016; Choplin
et al. 2018, Limongi $\&$ Chieffi 2018) and from the contribution of neutrino-
driven winds in CCSNe (e.g. Roberts et al. 2010). The observed heavy element
abundance ratios tend to show a higher spread of abundance ratios at the
metallicity of NGC 6522 relative to the models.

However, the models presented in Fig. 16 were optimised for old field stars of
the Galactic bulge, adopting the same nucleosynthesis that worked well for the
Galactic halo (Cescutti et al. 2013). Although this model is not specifically
made for a globular cluster, it is still useful. For instance, it shows the
extension of the dispersion that the enrichment due to rotating massive stars
can produce on these abundance ratios. The goal is to show that the predicted
scatter is indeed compatible with the dispersion observed in NGC 6522.
However, this scatter in field stars seems to appear at lower metallicities
(see Barbuy et al. 2018 where our model is compared with field bulge stars).
It is then plausible that other physical mechanisms are at play in the cluster
evolution (involving dynamical effects, and mass loss through winds). The
model for the field bulge stars would just give an idea of the mean cluster
abundances but not its scatter. A detailed description of the models presented
in Fig. 16, with a focus on the expected differences in the abundance ratio
scatter in the bulge and halo, will be presented in Cescutti et al. (2021, in
preparation).

In order to better interpret the heavy-element abundances of star B118, in
Fig. 17 we show the abundance pattern of stars 244819 and B118 in terms of
A(X)-A(Eu) (where A(X)=log(NX/NH+12), in a diagram idealised by Honda et al.
(2007) and Roederer et al. (2010, their Fig. 11). In this figure, A(X)-A(Eu)
versus Z of 244819 and B118 are compared with data from the typical r-element
star CS 31082-001 (Hill et al. 2002), and the typical LEPP star HD 122563
(Honda et al. 2007, Montes et al. 2007), the identified spinstar-enriched star
ROA 276 in $\omega$ Centauri (Yong et al. 2017), and the reference dwarf stars
HD 94028 and HD 140283 (Peterson et al. 2020; Siqueira-Mello et al. 2015).
First of all, Figure 17 indicates that 244819 and B118 are weakly enriched in
r elements. Figure 18 shows A(X)-A(Eu) versus Z for the four member stars from
this paper plus the four stars from Barbuy et al. (2014) for Y, Zr, Ba, and
La. It shows that the sample stars essentially follow the same pattern,
whereas 244819 and B118 show a larger abundance difference between first peak
and second peak elements.

Spite et al. (2018) suggested that the heavy element enrichment should take
place first due to a pure r process, followed by an enrichment of first-peak
elements only, and that this second mechanism would be detectable only in
weak-r-process stars. On the other hand, spinstars could be progenitors of
magneto-rotational supernovae, but in case the conditions do not allow
r-process elements to form in the final explosion, it is also possible that we
can only observe the signature of the s-process production in spinstars today.

There are a number of points that it is important to consider in our analysis.
First of all, the s-process efficiency in spinstars varies greatly if we
consider different theoretical stellar yields. While for instance the
s-process production in models by Frischknecht et al. (2016) and Choplin et
al. (2018) would stop in the Ba mass region, in models by Limongi $\&$ Chieffi
(2018) heavier elements up to lead could be produced. This uncertainty of
course affects Galactic chemical evolution (GCE) predictions (e.g. Cescutti et
al. 2013, Rizzuti et al. 2019, Prantzos et al. 2020, Rizzuti et al. 2021).
Additionally, CCSNe generated from slowly rotating progenitors or spinstars
can also eject other nucleosynthesis components made before the SN explosion
(similarly to the intermediate neutron capture process or i process - see e.g.
Roederer et al. 2016 and Banerjee et al 2018) or by explosion (similarly to
the zoo of neutrino-driven wind components - see e.g. Qian $\&$ Wasserburg
2008, Farouqi et al. 2009, Roberts et al. 2010, Arcones $\&$ Montes 2011). As
we mentioned earlier, all of these processes may contribute to the production
of Sr, Y, and Zr, while at low metallicities the i process can potentially
produce elements across the whole mass region below and beyond Fe, including
Sr, Y, Zr, and Ba, in different types of stars (e.g. Abate et al. 2016,
Roederer et al. 2016, Clarkson et al. 2018, Banerjee et al 2018). For
instance, the high [Ba/La] in stars 256289, 402370, B-107, B-128, and B-130
would be compatible with the i process (see e.g. Hampel et al. 2016).

Alternative possibilities of neutron-capture element enrichment are the
magneto-rotationally driven explosions of core-collapse supernovae (Winteler
et al. 2012), or s process taking place in asymptotic giant branch (AGB) stars
and subsequent mass transfer within a binary system (e.g. Cristallo et al.
2015). A study on possible nucleosynthesis processes is given in Hansen et al.
(2014).

This makes the observation of more elements per stellar target and at high
resolution of different stars in globular clusters such as NGC 6522 crucial.
Within this scenario the large star-to-star variations of heavy-element
enrichment could be a natural outcome of an intrinsic scatter of s-process
efficiencies in spinstars or the varying contribution of different processes
active before and during SN explosions in massive stars. On the other hand,
when abundances of several heavy elements are available it becomes possible to
disentangle the dominant nucleosynthesis component(s) that made the whole
observed abundance pattern (see e.g., Roederer et al. 2016, Peterson et al.
2020). Stars in NGC 6522 are carriers of the same signatures of the
nucleosynthesis processes active in the early galaxy and observed in halo
stars, even if they are more metal-rich as a result of a steeper age-
metallicity relation in the Galactic bulge (as suggested in C11).

Figure 14:  [Na/Fe] versus [O/Fe] anti-correlation in NGC 6522 stars compared
with stars in NGC 6121; and [Na/Fe] versus [N/Fe], [Al/Fe] and [Mg/Fe] versus
[Al/Fe] correlations. Symbols: blue-filled triangles represent the present
results for NGC 6522; red-filled triangles denote results for NGC 6522 from
Barbuy et al. (2014). Figure 15: [Sr,Y,Zr,Ba,La,Eu/Fe] versus [Fe/H] in bulge
stars. Symbols: Large, blue-filled circles mark present results of NGC 6522
member stars; large, blue open circles represent the present results of non-
member stars. Field stars: Blue open triangles represent red giants by Johnson
et al. (2012); magenta-filled triangles represent red giants by Siqueira-Mello
et al. (2016); cyan-filled triangles denote red giants by van der Swaelmen et
al. (2016); red-filled circles show dwarfs by Bensby et al. (2017); light blue
open circles mark metal-poor giants by Howes et al. (2016); acquamarine open
circles show metal-poor giants by Lamb et al. (2017). Globular cluster stars:
Orange-filled squares show NGC 6522 (Barbuy et al. 2014); green-filled
pentagons mark HP 1 (Barbuy et al. 2016); tomato-filled pentagons represent
NGC 6558 (Barbuy et al. 2018); magenta-filled pentagons show M62 (Yong et al.
2014). Figure 16: [Y/Ba] vs. [Fe/H] (upper panel) and [Y/Ba] versus [Ba/H]
(lower panel). Symbols: Large blue-filled circles show present results on NGC
6522 member stars; large blue open circles show present results on non-member
stars. Field stars: Magenta-filled triangles denote red giants by Siqueira-
Mello et al. (2016); green open circles mark dwarfs by Bensby et al. (2017);
light blue open circles show metal-poor giants by Howes et al. (2016);
acquamarine open circles represent metal-poor giants by Lamb et al. (2017).
Globular cluster stars: Orange-filled squares show NGC 6522 (Barbuy et al.
2014); green-filled pentagons represent HP 1 (Barbuy et al. 2016); tomato-
filled pentagons show NGC 6558 (Barbuy et al. 2018); magenta-filled pentagons
represent M62 (Yong et al. 2014). The underlying coloured model corresponds to
calculations for nucleosynthesis taking place in spinstars and in MRD
supernovae (Cescutti et al. 2018 and references therein).  Figure 17:
A(X)-A(Eu) versus atomic number Z for stars B118 and 244819 compared with
stars CS 31082-001, ROA 276, HD 94028, HD 140283, and HD 122563. All
abundances are normalised to the Sr abundance of B118. Symbols are identified
in the figure panel. Figure 18: A(X)-A(Eu) versus atomic number Z for star
B118 compared with the other seven sample stars. The abundances are not
normalised. Symbols are identified in the figure panel.

###  6.2 The two non-member stars 234816 and 244523

The stars 234816 and 244523 have the correct magnitudes, radial velocities and
metallicities to be members of NGC 6522. However, the Gaia proper motions
reported in Table 3 rule out their membership.

Could these two stars be former members that are evaporating from the cluster?
Madrid et al. (2017) studied evaporation rates as a function of galactocentric
distance RGC and time and predicted a very high evaporation at low RGC due to
the strong tidal field in the central parts of the Galaxy. NGC 6522 is
estimated to have a mass of 5.93$\times$104 M⊙ (Gnedin & Ostriker 1997), which
is not high for a globular cluster. It is located at RGC $\approx$ 1 kpc and
has an age above 12.1 - 12.4 Gyr using Dotter et al. (2008) isochrones and
even older (with 14.1 - 14.2 Gyr) using BaSTI isochrones (Kerber et al. 2018).
It can be seen from Fig. 6 of Madrid et al. (2017) that the evaporation rate
in a bulge cluster like NGC 6522 should be extremely high. It is therefore
acceptable to suggest that the two stars could be evaporating from the
cluster. However, a more detailed orbital calculation would be needed to check
this possibility, such as the one carried out by Hanke et al. (2020). In
particular, star 234816 has different alpha-element abundances and should not
be a member.

##  7 Conclusions

We derived a mean metallicity of [Fe/H]=$-$1.16$\pm$0.05 from the four sample
stars. Combined with the other four stars from Barbuy et al. (2014), the
result is a mean metallicity of [Fe/H]=$-$1.05$\pm$0.20.

Among the six stars analysed, two of them are indicated to be non-members from
Gaia proper motions; still, they have the correct magnitude, radial velocity,
and metallicity to be members. Only a fraction of about 0.5% of stars in the
Galactic bulge have metallicities below [Fe/H]$\mathrel{\hbox{\raise
2.15277pt\hbox{$<$}\hbox to0.0pt{\hss\lower 2.15277pt\hbox{$\sim$}}}}$-1.0
(Barbuy et al. 2018a). Therefore, we suggest that these stars could be
evaporating from the cluster; but even so, we do not include their abundances
in the discussion below. Star 244523 has abundances compatible with the member
stars. Star 234816 shows different high alpha-element abundances, and it could
be an intruder; hence, it could have been a bulge star with the correct
magnitude and metallicity to be considered a member before we had Gaia
measurements, but it was eventually revealed as a non-member star.

For the present results on the four confirmed member stars, together with
those by Barbuy et al. (2014), the alpha-elements show enhancements of
[O/Fe]=+0.38, [Mg/Fe]=$\approx$+0.28, [Si/Fe]$\approx$+0.19, and
[Ca/Fe]$\approx$+0.13, [Ti/Fe]$\approx$+0.13. A higher enhancement in O and
Mg, and a lower one in Si, Ca, and Fe can be explained by their formation in
hydrostatic conditions for the former, and in explosive nucleosynthesis for
the latter (e.g. Woosley & Weaver 1995; McWilliam 2016).

The r-process element Eu is enhanced by [Eu/Fe]=+0.40. The $\alpha$-element
enhancements in O and Mg, together with that of the r-process element Eu, are
indicative of a fast early enrichment by type II supernovae. With regard to
the indicators of multiple stellar populations, we suggest that Na shows an
anti-correlation with O, and more clearly a correlation with N and Al, whereas
Mg and Al are also correlated.

A main objective of this study is the verification of the enhancement of
s-element abundances, and the possibility of an early enrichment by spinstars.
The neutron-capture elements typically indicated as s-process elements are
enhanced with [Y/Fe]=+0.33, [Zr/Fe]=+0.23, [Ba/Fe]=+0.35, and [La/Fe]=+0.23.
In addition to this observation we find the following:

a) There are significant relative abundance variations between neutron-capture
elements, where [Y/Ba] is particularly enhanced in two stars.

b) [Ba/Eu] = $-$0.14, $-$0.07, $-$0.40, +0.16, in the four member stars, as a
measure of the s- to r- process nucleosynthesis, tends to be slightly below
solar. This result is still compatible with the interpretation given in B09
and C11 that their production cannot be attributed to the r-process only, as
first suggested by Truran (1981) for very old stars.

As discussed in C11, the presence of s-process element enhancements in very
old stars could be due to an s-process enrichment of the primordial matter
from which the cluster formed, processed in spinstars (e.g. Frischknecht et
al. 2016). Alternatively, the production of heavy elements could be due to a
combination of different nucleosynthesis processes, in particular for the
atomic mass region of Sr. Another possibility would be to have spinstars
producing the s-process elements during its hydrostatic phase and producing
the r-process elements at the supernova explosion, and therefore to be the
source of both. This is possible if the spinstars rotate fast enough to
produce an MHD explosion with the right conditions to produce an r process
(Nishimura et al. 2017 and references therein). However, within this scenario
it is extremely uncertain to predict the observed relative contribution of the
s-process and r-process elements, since the r-process-rich material could be
ejected asymmetrically and/or could carry a large range of efficiency in
r-process production. Therefore, a possible outcome could be that the final
enrichment produced by such a spinstar and magnetic-rotationally driven (MRD)
SN is dominated by the r-process signature, because of higher yields of the
MRD SNe compared to those of the s process (Spite et al. 2018). On the other
hand, the enrichment of the local interstellar medium could also be s-process
rich, depending on the spatial distribution of different nucleosynthesis
products in the SN ejecta. Finally, nucleosynthesis taking place in AGB stars
and the i process might be alternative possibilities that should be further
inspected.

Taking into account the different uncertainties at play, we confirm the
conclusions from Barbuy et al. (2014) that the observed abundances are
compatible with the s-process production in spinstars. However, we cannot rule
out that the same enrichment signature could be produced by a combination of
nucleosynthesis processes active in the early generations of stars.

###### Acknowledgements.

BB, EC, LM, SO acknowledge partial financial support from CAPES- Finance code
001, CNPq and Fapesp. SOS acknowledges the FAPESP PhD fellowship no.
2018/22044-3 and the DGAPA-PAPIIT grant IG100319. RH and MP acknowledges
support from the IReNA AccelNet Network of Networks, supported by the National
Science Foundation under Grant No. OISE-1927130 and from the World Premier
International Research Centre Initiative (WPI Initiative), MEXT, Japan. CC,
RH, MP and GC acknowledge support from the ChETEC COST Action (CA16117),
supported by COST (European Cooperation in Science and Technology). MP thanks
the support to NuGrid from STFC (through the University of Hull Consolidat ed
Grant ST/R000840/1), and access to viper, the University of Hull High
Performance Computing Facility. MP acknowledges the support from the
Ledulet-2014 Program of the Hungarian Academy of Sciences (Hungary). SO
acknowledges the Italian Ministero dell’Università e della Ricerca Scientifica
e Tecnologica (MURST), Italy. This work has made use of data from the European
Space Agency (ESA) missionGaia (http://www.cosmos.esa.int/gaia), processed by
the Gaia Data Processing and Analysis Consortium
(DPAC,http://www.cosmos.esa.int/web/gaia/dpac/consortium). Funding for the
DPAC has been provided by national institutions, in particular the
institutions participating in the Gaia Multilateral Agreement.

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##  Appendix A Candidate members observed with GIRAFFE

The radial velocities of GIRAFFE observations were measured by fitting the
cross-correlation peak with the solar spectra available at the ESO
portal888http://www.eso.org/observing/dfo/quality/GIRAFFE/pipeline/solar.html,
and the heliocentric radial velocities were derived. The final heliocentric
radial velocities of GIRAFFE spectra were determined by taking the average of
the mean heliocentric velocity of each setup. This method showed a better
agreement with their counterparts in the more reliable UVES spectra
velocities. Signal to noise was measured in points of continua in each of the
setups, resulting in a mean S/N$\sim$ 70 in the HR11 setup and S/N$\sim$ 93 in
the HR12 setup.

We applied a selection of stars with radial velocities within $\pm$12 Km/s of
that of NGC 6522 to both the 2012 and the 2016 observations. On the selected
sample, we combined the radial velocities with Gaia proper motions to derive
membership probabilities. Table 14 reports the selected stars, their
identification, coordinates, magnitudes, and respective heliocentric radial
velocities.

Table 14: OGLE and 2MASS numbers, coordinates, magnitudes including $JHK_{s}$ from both 2MASS and VVV surveys, and final radial velocities for stars observed with GIRAFFE in 2012 and 2016.  OGLE |  ---2MASS ID | $\alpha_{2000}$ |  ---$\delta_{2000}$ |  ---$V$ |  ---$I$ |  ---$J$ |  ---$H$ |  ---$K_{\rm s}$ |  ---$J_{\rm VVV}$ |  ---HVVV |  KVVV | v${}_{\rm r}^{hel}$ |  BJohnson | member
---|---|---|---|---|---|---|---|---|---|---|---|---|---|---
GIRAFFE SAMPLE

|  |  |  | GIRAFFE sample (2012) |  |  |  |  |  |  |
244853 | 18033424-3002109 | 18:03:34.13 |  $-$30:02:11.1 | 16.085 | 14.457 | 11.304 | 10.994 | 11.374 | — | — | — |  $-$11.79 | 17.363 | 99.5
402384 | 18034256-3001409 | 18:03:42.55 |  $-$30:01:40.5 | 16.055 | 14.388 | 13.259 | 12.506 | 12.336 | 13.0877 | 12.4530 | 12.2813 | -20.04 | 17.402 | 99.0
|  |  |  | GIRAFFE sample (2016) |  |  |  |  |  |  |
244551 | 18033361-3002389 | 18:03:33.61 |  $-$30:02:38.9 | 16.134 | 14.455 | 13.321 | 12.551 | 12.523 | 13.1628 | 12.5364 | 12.3582 | -8.63 | 17.486 | 97.8
244555 | 18033467-3002305 | 18:03:34.67 |  $-$30:02:32.2 | 16.536 | 14.480 | 11.436 | 10.595 | 10.478 | — | — | — |  $-$2.33 | 18.231 | 99.3
244813 | — | 18:03:29.00 |  $-$30:02:28.3 | 16.147 | 14.578 | — | — | — | 13.1904 | 12.5908 | 12.4422 |  $-$10.80 | 17.365 | 43.7
256298 | 18033214-3000350 | 18:03:32.13 |  $-$30:00:34.9 | 16.038 | 14.430 | 13.133 | 12.306 | 12.173 | 13.1264 | 12.4863 | 12.2987 |  $-$22.39 | 17.329 | 99.8
402371 | 18033854-3002075 | 18:03:38.57 |  $-$30:02:07.3 | 16.065 | 14.417 | 12.932 | 12.218 | 12.254 | 13.1894 | 12.6119 | 12.4397 |  $-$17.52 | 17.394 | 99.6
402508 | 18034025-3003178 | 18:03:40.16 |  $-$30:03:18.1 | 16.278 | 14.590 | 11.112 | 11.373 | 11.216 | 12.8682 | 12.0634 | 12.1385 |  $-$4.32 | 17.593 | 91.9

##  Appendix B  Hyperfine structure of Ba II lines

Table 15: Atomic constants for BaII used to compute hyperfine structure. A constants are from Rutten (1978), B constants from Biehl (1976), and B constants not available in the literature are assumed as null.  species |  $\lambda$ (Å) | Lower level | J | A(mK) | A(MHz) | B(mK) | B(MHz) | Upper level | J | A(mK) | A(MHz) | B(mK) | B(MHz)
---|---|---|---|---|---|---|---|---|---|---|---|---|---
135BaII | 5853.668 |  5d 2D3/2 | 3/2 | 3.56 | 106.7261 | 0 | 0 |  6p 2P∘3/2 | 3/2 | +3.47 | 104.028 | +2.2 | 65.9544
137BaII | 5853.668 |  5d 2D3/2 | 3/2 | 3.97 | 119.0176 | 0 | 0 |  6p 2P∘3/2 | 3/2 | +3.88 | 116.3195 | +3.25 | 97.4326
135BaII | 6141.713 |  5d 2D5/2 | 5/2 | 1.49 | 44.6691 | 0 | 0 |  6p 2P∘3/2 | 3/2 | +3.47 | 104.028 | +2.2 | 65.9544
137BaII | 6141.713 |  5d 2D5/2 | 5/2 | 1.66 | 49.7655 | 0 | 0 |  6p 2P∘3/2 | 3/2 | +3.88 | 116.3195 | +3.25 | 97.4326
Table 16: Hyperfine structure for Ba II 5853.675 Å line.  5853.675 $\rm\AA$; $\chi$=0.604321 eV |  |
---|---|---
log gf(total) = $-$1.10 |  |
$\lambda$ ($\rm\AA$) | log gf | iso |  |
5853.673 |  $-$2.3441 | 135 |  |
5853.674 |  $-$2.3441 | 135 |  |
5853.674 |  $-$2.7421 | 135 |  |
5853.673 |  $-$2.1400 | 135 |  |
5853.677 |  $-$2.1400 | 135 |  |
5853.676 |  $-$2.0431 | 135 |  |
5853.671 |  $-$2.1980 | 135 |  |
5853.679 |  $-$2.1980 | 135 |  |
5853.675 |  $-$1.5959 | 135 |  |
5853.673 |  $-$2.3441 | 137 |  |
5853.674 |  $-$2.3441 | 137 |  |
5853.674 |  $-$2.7421 | 137 |  |
5853.673 |  $-$2.1400 | 137 |  |
5853.677 |  $-$2.1400 | 137 |  |
5853.676 |  $-$2.0431 | 137 |  |
5853.670 |  $-$2.1980 | 137 |  |
5853.680 |  $-$2.1980 | 137 |  |
5853.674 |  $-$1.5959 | 137 |  |
5853.675 |  $-$2.7562 | 134 |  |
. 5853.675 |  $-$2.2451 | 136 |  |
5853.675 |  $-$1.2845 | 138 |  |

In Table 15, we give the hyperfine structure constants for the BaII 5853.675 Å
line, and in Table 16 we report the list of lines sub-divided due to the
hyperfine structure, thus completing the similar results for the 6141.713 and
6496.897 AA lines reported in Barbuy et al. (2014).

##  Appendix C Equivalent widths and atomic data

[x]l@rrrrl@r@rrrrrrrrr  Equivalent widths of Fe I and Fe II lines employed,
their wavelengths, excitation potential (eV), oscillator strengths,and van der
Waals broadening constant C6 adopted.
continued.
species $\lambda$ $\chi_{ex}$ loggf --C6 234816 244523 244819 256289 402322
402370
species $\lambda$ $\chi_{ex}$ loggf --C6 234816 244523 244819 256289 402322
402370
(Å) (eV) KUR VALD NIST ADOPT

Fe I 5853.15 1.48 —$-$5.280— --$-$5.13 --0.174E-31 43.4 13.9 — — 9.6 —
Fe I 5856.08 4.29 $-$1.640$-$1.328$-$1.328 $-$1.64 -0.423E-30 38.3 — — — 17.4
—
Fe I 5858.78 4.22 $-$2.260$-$2.260— $-$2.26 0.376E-30 5.0 — — — — —
Fe I 5859.60 4.55 $-$0.386$-$0.418— $-$0.60 0.674E-30 66.8 66.8 — — 46.7 —
Fe I 5861.11 4.28 $-$2.450$-$2.450— $-$2.45 0.415E-30 7.0 — — — — —
Fe I 5881.28 4.61 $-$1.840$-$1.840— $-$1.84 0.415E-30 99. — — — 10.4 —
Fe I 5902.48 4.59 $-$1.810$-$1.810— $-$1.81 0.708E-30 6.6 — — — 10.2 —
Fe I 5905.67 4.65 $-$0.730$-$0.730$-$0.77 $-$0.73 0.623E-30 65.9 38.7 32.9
30.3 31.2 —
Fe I 5916.25 2.45 $-$2.994$-$2.994$-$2.99 $-$2.99 0.428E-31 98.4 43.2 65.9
61.5 46.3 56.5
Fe I 5927.79 4.65 $-$1.090$-$1.090$-$1.07 $-$1.09 0.786E-30 27.6 42.0 9.7 —
16.1 12.4
Fe I 5929.67 4.55 $-$1.410$-$1.410$-$1.38 $-$1.38 0.644E-30 45.4 35.7 26.1 —
13.0 10.6
Fe I 5930.18 4.65 $-$0.230$-$0.230— $-$0.23 0.784E-30 110.5 — — 60.8 54.9 46.7
Fe I 5934.65 3.93 $-$1.170$-$1.170$-$1.12 $-$1.12 0.234E-30 64.0 71.7 57.1 —
73.0 41.7
Fe I 5952.72 3.98 $-$1.440$-$1.440$-$1.39 $-$1.39 0.249E-30 50.9 46.7 47.2
39.6 30.5 43.7
Fe I 5956.69 0.86 $-$4.605$-$4.605$-$4.61 $-$4.60 0.948E-32 133.2 97.5 70.8 —
72.2 70.3
Fe I 5975.35 4.84 $-$0.822$-$1.292— $-$0.82 0.300E-31 37.1 42.0 26.3 39.5 44.3
36.3
Fe I 5983.68 4.55 $-$1.521$-$0.485— $-$0.78 0.622E-30 99. 119.2 42.0 45.3 28.0
23.1
Fe I 5987.06 4.79 $-$0.654$-$0.432— $-$0.42 0.102E-29 99. — 40.0 — 33.9 37.2
Fe II 5991.38 3.15 $-$3.557$-$3.540$-$3.60 $-$3.54 0.775E-32 18.4 18.4 19.6
29.0 27.2 8.9
Fe I 6003.01 3.88 $-$1.120$-$1.120— $-$1.12 0.211E-30 91.2 81.1 62.5 63.8 75.3
44.2
Fe I 6005.54 2.59 $-$2.922$-$3.602— $-$3.61 0.300E-31 18.4 29.7 18.4 24.3 20.8
14.9
Fe I 6008.56 3.88 $-$1.291$-$0.982— $-$0.99 0.300E-31 89.3 81.7 — 61.3 — 55.1
Fe I 6012.21 2.22 $-$4.200$-$4.038$-$4.04 $-$3.94 0.362E-31 — — — 17.8 — —
Fe I 6020.17 4.61 $-$0.270$-$0.270— $-$0.27 0.300E-31 62.8 92.0 64.3 — 90.2 —
Fe I 6024.05 4.55 $-$0.120$-$0.120— $-$0.11 0.606E-30 99.9 — 81.9 — 68.7 77.8
Fe I 6027.06 4.07 $-$1.210$-$1.089$-$1.09 $-$1.09 0.273E-30 65.6 54.1 — 52.3
42.8 39.0
Fe I 6056.01 4.73 $-$0.460$-$0.460— $-$0.46 0.849E-30 51.2 39.2 51.0 42.7 42.5
25.1
Fe I 6065.49 2.61 $-$1.530$-$1.530$-$1.530 $-$1.53 0.477E-31 145.0 116.9 116.0
— 131.3 112.4
Fe I 6078.50 4.79 $-$0.481$-$0.323— $-$0.40 0.951E-30 60.1 47.5 50.5 — 60.2
53.0
Fe I 6079.00 4.65 $-$1.120$-$1.120$-$1.10 $-$1.10 0.710E-30 47.5 18.1 20.0 —
28.2 12.7
Fe I 6082.71 2.22 $-$3.573$-$3.573$-$3.57 $-$3.57 0.300E-31 59.9 59.1 42.5 —
36.9 26.3
Fe II 6084.11 3.20 $-$3.808$-$3.780 $-$3.90 $-$3.79 0.787E-32 7.9 19.1 9.1
15.1 12.0 5.2
Fe I 6093.64 4.61 $-$1.500$-$1.500$-$1.47 $-$1.47 0.638E-30 17.0 10.8 3.7 14.8
14.2 10.2
Fe I 6094.37 4.65 $-$1.940$-$1.940$-$1.92 $-$1.92 0.703E-30 11.7 99.0 — — —
2.3
Fe I 6136.99 2.20 $-$2.950$-$2.950$-$2.950 $-$2.95 0.282E-31 — — — 75.8 — —
Fe I 6137.70 2.59 $-$1.403$-$1.403$-$1.403 $-$1.40 0.457E-31 110.4 158.0 131.0
122.2 — 120.7
Fe II 6149.25 3.89 $-$2.724$-$2.720$-$2.80 $-$2.69 0.943E-32 15.6 34.3 17.6
20.0 18.9 12.1
Fe I 6151.62 2.18 $-$3.299$-$3.299$-$3.30 $-$3.30 0.305E-31 84.0 108.9 60.0
57.8 68.4 50.7
Fe I 6157.73 4.08 $-$1.260$-$1.260$-$1.22 $-$1.25 0.261E-30 71.1 24.5 59.8
43.4 59.8 43.2
Fe I 6159.38 4.61 $-$1.970$-$1.970— $-$1.85 0.625E-30 12.1 99.0 — — 11.5 5.0
Fe I 6165.36 4.14 $-$1.550$-$1.474$-$1.47 $-$1.55 0.284E-30 40.7 32.4 33.7
24.1 26.3 23.8
Fe I 6173.33 2.22 $-$2.880$-$2.880$-$2.88 $-$2.88 0.882E-30 — — — 75.8 80.1 —
Fe I 6180.21 2.73 $-$2.780$-$2.586$-$2.65 $-$2.65 0.519E-31 82.6 94.8 60.0
52.1 81.3 49.0
Fe I 6187.99 3.94 $-$2.204$-$1.720$-$4.16 $-$1.72 0.211E-30 63.6 42.9 30.0 —
20.2 29.4
Fe I 6200.31 2.61 $-$2.437$-$2.437$-$1.67 $-$2.44 0.507E-31 — — — 75.5 134.1 —
Fe I 6213.43 2.22 $-$2.660$-$2.482$-$2.48 $-$2.49 0.311E-31 111.6 89.2 — 92.0
99.4 87.7
Fe I 6219.28 2.20 $-$2.433$-$2.433$-$2.43 $-$2.43 0.305E-31 83.0 143.6 98.4
98.3 — 97.3
Fe I 6220.78 3.88 $-$2.460$-$2.460— $-$2.46 0.192E-30 19.7 — — — 9.8 6.7
Fe I 6226.73 3.88 $-$2.220$-$2.220— $-$2.20 0.191E-30 36.6 13.9 17.3 22.1 21.1
9.2
Fe I 6229.23 2.84 $-$2.970$-$2.805$-$2.805 $-$2.97 0.571E-31 54.5 49.4 28.4
37.3 24.0 28.8
Fe I 6240.65 2.22 $-$3.380$-$3.233$-$3.17 $-$3.21 0.309E-31 75.8 66.0 56.0
51.1 50.2 49.7
Fe I 6246.32 3.60 $-$0.960$-$0.733$-$0.88 $-$0.88 0.133E-30 119.2 117.3 92.4
93.0 — 91.6
Fe II 6247.56 3.89 $-$2.329$-$2.310$-$2.40 $-$1.98 0.881E-32 22.4 51.1 31.8
47.8 50.1 22.3
Fe I 6252.56 2.40 $-$1.687$-$1.687$-$1.687 $-$1.69 0.366E-31 51.3 126.6 130.0
118.7 145.6 119.7
Fe I 6254.25 2.28 $-$2.480$-$2.443$-$2.426 $-$2.43 0.326E-31 166.0 129.7 —
113.1 113.5 107.9
Fe I 6265.14 2.18 $-$2.550$-$2.550$-$2.55 $-$2.53 0.295E-31 170.3 118.1 104.2
— 93.5 85.4
Fe I 6270.23 2.86 $-$2.710$-$2.464$-$2.61 $-$2.61 0.575E-31 76.1 70.1 57.5
57.5 61.4 54.8
Fe I 6271.28 3.33 $-$2.950$-$2.703$-$2.70 $-$2.81 0.945E-31 30.1 34.8 11.4 —
8.2 17.3
Fe I 6297.80 2.22 $-$2.740$-$2.740$-$2.74 $-$2.74 0.304E-31 106.6 100.4 98.6
81.8 90.4 74.6
Fe I 6301.50 3.65 $-$0.672$-$0.718$-$0.72 $-$0.60 0.138E-30 109.0 108.3 101.6
95.3 182.0 95.1
Fe I 6302.50 3.69 $-$1.131$-$0.968— $-$0.91 0.145E-30 90.2 82.9 — 72.4 — —
Fe I 6311.50 2.83 $-$3.230$-$3.141$-$3.14 $-$3.22 0.551E-31 71.9 31.1 20.0
26.0 22.0 13.2
Fe I 6315.31 4.14 $-$1.232$-$1.232$-$1.232 $-$1.23 0.265E-30 44.4 27.0 39.2
42.2 44.6 35.5
Fe I 6315.81 4.08 $-$1.710$-$1.710$-$1.66 $-$1.66 0.243E-30 38.8 40.8 25.0 —
31.6 15.4
Fe I 6322.69 2.59 $-$2.426$-$2.426$-$2.43 $-$2.43 0.485E-31 — — — 74.6 96.5 —
Fe I 6335.33 2.20 $-$2.230$-$2.177$-$2.18 $-$2.18 0.295E-31 135.3 96.1 120.0 —
— 105.0
Fe I 6336.83 3.69 $-$1.050$-$0.856$-$0.86 $-$1.05 0.143E-30 94.9 107.9 98.3
83.9 — 79.7
Fe I 6344.15 2.43 $-$2.923$-$2.923$-$2.923 $-$2.92 0.366E-31 81.1 81.8 62.0
61.9 68.5 61.8
Fe I 6355.03 2.85 $-$2.420$-$2.350$-$2.29 $-$2.29 0.549E-31 90.9 93.2 90.0
68.7 83.0 48.2
Fe II 6369.46 2.89 $-$4.253$-$4.160$-$4.29 $-$4.11 0.742E-32 …. — — 15.5 23.1
—
Fe I 6380.74 4.19 $-$1.400$-$1.376$-$1.38 $-$1.38 0.277E-30 74.3 141.1 — 35.8
27.2 14.7
Fe I 6392.54 2.28 $-$4.030$-$4.030— $-$4.03 0.313E-31 42.2 26.6 — — 28.6 11.5
Fe I 6393.60 2.43 $-$1.620$-$1.432$-$1.576 $-$1.58 0.361E-31 156.4 134.6 —
120.9 135.0 123.1
Fe I 6400.00 3.60 $-$0.520$-$0.290$-$0.290 $-$0.29 0.402E-30 — — — 111.7 — —
Fe I 6408.02 3.69 $-$0.970$-$1.018$-$1.02 $-$1.00 0.139E-30 88.8 75.0 — 82.4
88.7 —
Fe I 6411.11 4.73 $-$2.026$-$1.935— $-$2.21 0.679E-30 — — — — — —
Fe I 6411.65 3.65 $-$0.820$-$0.595$-$0.72 $-$0.72 0.132E-30 128.7 111.8 102.0
95.6 109.0 96.2
Fe II 6416.92 3.89 $-$2.740$-$2.650$-$2.90 $-$2.64 0.930E-32 28.3 36.4 20.5
26.1 44.0 19.8
Fe I 6419.94 4.73 $-$0.240$-$0.240$-$0.27 $-$0.25 0.675E-30 54.0 60.2 54.2
52.9 66.6 45.7
Fe I 6421.35 2.28 $-$2.027$-$2.027$-$2.027 $-$2.03 0.310E-31 122.5 121.9 120.0
— — 108.1
Fe I 6430.85 2.18 $-$2.006$-$2.006$-$2.006 $-$2.01 0.281E-31 137.9 129.9 131.0
— 130.0 113.3
Fe II 6432.68 2.89 $-$3.708$-$3.520$-$3.50 $-$3.57 0.742E-32 35.3 28.6 32.3
32.8 36.0 29.0
Fe II 6456.38 3.90 $-$2.075$-$2.100$-$2.20 $-$2.05 0.930E-32 41.4 33.9 27.2
51.9 54.7 24.7
Fe I 6469.20 4.83 $-$0.770$-$0.770$-$0.81 $-$0.81 0.802E-30 63.1 18.1 30.3 —
33.4 27.3
Fe I 6475.62 2.56 $-$2.940$-$2.942$-$2.94 $-$2.94 0.400E-31 95.2 75.1 76.2
57.8 88.0 49.8
Fe I 6481.87 2.28 $-$2.984$-$2.984$-$2.98 $-$2.98 0.305E-31 98.9 50.2 76.2
67.3 78.8 —
Fe I 6494.98 2.40 $-$1.273$-$1.273$-$1.273 $-$1.27 0.340E-31 155.2 138.2 157.4
138.7 140.0 140.1
Fe II 6516.08 2.89 $-$3.450$-$3.320$-$3.37 $-$3.31 0.742E-32 52.2 55.3 29.3
55.2 48.9 35.1
Fe I 6518.37 2.83 $-$2.750$-$2.460$-$2.30 $-$2.30 0.516E-31 58.1 102.4 — 48.2
56.4 —
Fe I 6533.93 4.56 $-$1.460$-$1.460$-$1.430 $-$1.43 0.497E-30 34.7 6.1 13.7 — —
18.6
Fe I 6546.24 2.76 $-$1.650$-$1.536$-$1.54 $-$1.54 0.472E-31 131.8 106.9 112.2
106.3 — 111.2
Fe I 6569.21 4.73 $-$0.420$-$0.127$-$0.45 $-$0.45 0.622E-30 73.6 82.9 60.6
59.3 — 38.1
Fe I 6574.23 0.99 — $-$5.023$-$5.004 $-$5.00 0.129E-31 — — — 48.2 — —
Fe I 6575.02 2.59 $-$2.820$-$2.710$-$2.710 $-$2.71 0.468E-31 — — — 66.9 67.3 —
Fe I 6581.21 1.48 $-$4.860$-$4.679$-$4.68 $-$4.85 0.142E-31 75.2 43.8 33.9 —
41.2 16.2
Fe I 6591.31 4.59 $-$2.070$-$2.070— $-$2.00 0.476E-30 4.1 — — — — 2.9
Fe I 6593.87 2.43 $-$2.422$-$2.422$-$2.42 $-$2.42 0.341E-31 115.0 99.0 — 86.9
87.2 82.4
Fe I 6597.56 4.79 $-$1.070$-$1.070$-$1.05 $-$1.05 0.701E-30 46.0 31.5 — 19.8 —
8.4
Fe I 6608.04 2.28 $-$4.030$-$4.030— $-$4.03 0.294E-31 16.7 33.8 7.4 16.3 18.2
10.5
Fe I 6609.11 2.56 $-$2.692$-$2.692$-$2.69 $-$2.69 0.385E-31 94.6 29.4 — 69.3
54.9 67.1
Fe I 6627.54 4.55 $-$1.680$-$1.680— $-$1.68 0.437E-30 13.3 — 10.7 — — 9.4
Fe I 6678.00 2.69 $-$1.470$-$1.418$-$1.418 $-$1.42 0.428E-31 160.0 200.0 124.0
120.9 130.0 117.4
Fe I 6699.14 4.59 $-$2.190$-$2.101$-$2.101 $-$2.10 0.452E-30 5.2 — — — — —
Fe I 6705.10 4.61 $-$1.470$-$1.382— $-$1.06 0.467E-30 38.1 30.1 — 26.4 — 3.3
Fe I 6726.67 4.61 $-$0.952$-$1.094— $-$1.09 0.447E-30 32.5 36.4 — — — 18.7
Fe I 6739.52 1.56 $-$4.950$-$4.794$-$4.79 $-$4.80 0.147E-31 42.7 18.5 — — —
8.9