# METAL-POOR STARS OBSERVED WITH THE AUTOMATED PlANET FINDER TELESCOPE. I.
DISCOVERY OF FIVE CARBON-ENHANCED METAL-POOR STARS FROM LAMOST

Mohammad K. Mardini Key Lab of Optical Astronomy, National Astronomical
Observatories, Chinese Academy of Sciences, A20 Datun Road, Chaoyang, Beijing
100102, China  School of Astronomy and Space Science, University of Chinese
Academy of Sciences, No.19(A) Yuquan Road, Shijingshan District, Beijing,
100049, China  Haining Li Key Lab of Optical Astronomy, National Astronomical
Observatories, Chinese Academy of Sciences, A20 Datun Road, Chaoyang, Beijing
100102, China  Vinicius M. Placco Department of Physics, University of Notre
Dame, Notre Dame, IN 46556, USA  JINA Center for the Evolution of the
Elements, USA  Sofya Alexeeva Shandong Provincial Key Laboratory of Optical
Astronomy and Solar-Terrestrial Environment, Institute of Space Sciences,
Shandong University, Weihai 264209, China  Daniela Carollo  INAF Astrophysical
Observatory of Turin - Italy  Ali Taani  Physics Department, Faculty of
Science, Al-Balqa Applied University,Jordan  Iminhaji Ablimit  Department of
Astronomy, Kyoto University, Kitashirakawa-Oiwake-cho, Sakyo-ku, Kyoto
606-8502, Japan  Liang WANG  Max-Planck-Institut fur extraterrestrische
Physik, Giessenbachstrasse, D-85748 Garching, Germany  Gang Zhao Key Lab of
Optical Astronomy, National Astronomical Observatories, Chinese Academy of
Sciences, A20 Datun Road, Chaoyang, Beijing 100102, China  School of Astronomy
and Space Science, University of Chinese Academy of Sciences, No.19(A) Yuquan
Road, Shijingshan District, Beijing, 100049, China
Haining Li<EMAIL_ADDRESS>
Gang Zhao<EMAIL_ADDRESS>
###### Abstract

We report on the discovery of five carbon-enhanced metal-poor (CEMP) stars in
the metallicity range of $-3.3<$ [Fe/H] $<-2.4$. These stars were selected
from the LAMOST DR3 low-resolution (R$\sim$ 2,000) spectroscopic database as
metal-poor candidates and followed-up with high-resolution spectroscopy
(R$\sim$110,000) with the LICK/APF. Stellar parameters and individual
abundances for 25 chemical elements (from Li to Eu) are presented for the
first time. These stars exhibit chemical abundance patterns that are similar
to those reported in other literature studies of very and extremely metal-poor
stars. One of our targets, J2114$-$0616, shows high enhancement in carbon
([C/Fe]=1.37), nitrogen ([N/Fe]= 1.88), barium ([Ba/Fe]=1.00), and europium
([Eu/Fe]=0.84). Such chemical abundance pattern suggests that J2114$-$0616 can
be classified as CEMP-r/s star. In addition, the star J1054+0528 can be
classified as a CEMP-rI star, with [Eu/Fe]=0.44 and [Ba/Fe]=$-$0.52. The other
stars in our sample show no enhancements in neutron-capture elements and can
be classified as CEMP-no stars. We also performed a kinematic and dynamical
analysis of the sample stars based on Gaia DR2 data. The kinematic parameters,
orbits, and binding energy of these stars, show that J2114$-$0616 is member of
the outer halo population, while the remaining stars belong to the inner halo
population but with an accreted origin. Collectively, these results add
important constraints on the origin and evolution of CEMP stars as well as on
their possible formation scenarios.

Galaxy: halo— stars:fundamental parameters— stars: abundances—-stars:
kinematics and dynamics — stars: Population II— stars: chemically peculiar

††software: MOOG (Sneden, 1973; Sobeck et al., 2011), IRAF (Tody, 1986, 1993),
NumPy (van der Walt et al., 2011), SciPy (Jones et al., 2001–), Matplotlib
(Hunter, 2007), Astropy (Astropy Collaboration et al., 2013, 2018) DETAIL
(Butler & Giddings, 1985), SynthV_NLTE (Ryabchikova et al., 2016), binmag
(Kochukhov, 2010)

##  1 Introduction

Metal-poor stars play a crucial role in galactic archaeology, since they
represent a fossil record of the nucleosynthesis products of their progenitors
(McWilliam et al., 1995; Beers & Christlieb, 2005; Chiaki et al., 2017; Jeon
et al., 2017). These stars contain detailed information about the past of
their host systems, which can be used to study the early Universe and the
beginning of star and galaxy formation (Frebel & Norris, 2015). One of the
primary tools to study metal-poor stars is the determination of their chemical
abundances, elements heavier than lithium reflect the extent of chemical
enrichment within its natal cloud (Christlieb et al., 2002; Frebel et al.,
2005; Frebel & Norris, 2015).

Many studies have indicated that metal-poor stars may show high carbon-to-iron
ratios, and thus classified as carbon-enhanced metal-poor stars (Beers &
Christlieb, 2005; Aoki et al., 2007; Hansen et al., 2016a; Placco et al.,
2016; Kielty et al., 2017; Roriz et al., 2017; Cruz et al., 2018; Caffau et
al., 2018). Furthermore, CEMP stars can be divided into four sub-classes,
according to their neutron-capture elements nature. The first sub-class is the
CEMP-s, where stars show high enhancements in carbon, together with
enhancements in elements formed mainly by the slow (s-) neutron capture
process (Placco et al. 2013; Abate et al. 2015). The second sub-class is
CEMP-r/s, where stars show enhancements in both slow (s-) and rapid (r-)
process material. The studies of Lucatello et al. (2005), Starkenburg et al.
(2014) and Hansen et al. (2016b) confirmed the binarity of CEMP-s stars, so
their peculiar chemical pattern can be explained by pollution from an
asymptotic giant branch (AGB) star companion, which has since become a white
dwarf. The wide variety enhancement associated with CEMP-r/s can not be
explained by the binarity of involving AGB stars (Beers & Christlieb, 2005;
Aoki et al., 2007).

Recently, the Laser Interferometer Gravitational Wave Observatory (LIGO) and
Virgo Collaboration (LVC) were able to detect the loudest gravitational wave
(GW170817), generated by binary neutron star merger (Abbott et al., 2017).
Shappee et al. (2017) studied the optical counterpart of GW170817 namely Swope
Supernova Survey 2017a (SSS17a), where they showed (in Figure 4A) that SSS17a
is consistent with the expectations for r-process heating, which supports the
idea that the main site of the r-process production is the neutron star
mergers (Qian & Wasserburg, 2007; Arnould et al., 2007; Troja et al., 2017). A
few CEMP stars exhibit overabundances of the elements formed mainly by the
r-process (Burbidge et al., 1957). In all probability, their r-process
enhancement took place in their birth cloud, when it is enriched by r-process
material. These stars are known as CEMP-r. The r-process-rich ejecta mix with
the interstellar medium (ISM), which will become star-forming regions and form
new generations of r-rich stars (Ji et al., 2016).

Recently, it has been recognized that the majority of CEMP stars at [Fe/H]
$\leq-4$ show no enhancement in neutron capture elements (e.g., Christlieb et
al. 2002; Frebel et al. 2005; Caffau et al. 2011; Keller et al. 2014) . These
stars are classified as CEMP-no. While CEMP-s were have been established as
member of binary systems, CEMP-no were found to be inconsistent with the
binary properties of the CEMP-s class (Starkenburg et al., 2014; Hansen et
al., 2016a), thereby strongly indicating a different physical origin of their
carbon-enhancement. These stars most likely appear to exhibit abundance
patterns of its natal cloud, that still makes their unusual chemical
composition a puzzle (Ito et al., 2013; Placco et al., 2014, 2015, 2016;
Roederer et al., 2016; Yoon et al., 2016).

In addition to carbon enhancement, rare CEMP stars exhibit nitrogen
overabundances, typically with [N/Fe] $>$ 1.0 and [N/C] $>$ 0.5, although this
phenomenon depends on metallicity, it appears to be more frequent at [Fe/H]
$<2.8$ ($\sim$ factor of 10) (Bessell & Norris, 1982; Johnson et al., 2007;
Pols et al., 2009, 2012; Roederer et al., 2014a). The existence of a
population of nitrogen-enhanced metal-poor (NEMP) stars was predicted from the
same evolution scenario, since AGB stars efficiently cycle carbon into
nitrogen in their envelopes. Detailed AGB nucleosynthesis models of low
initial mass ($<$2.5 M$\odot$) produce carbon, but do not produce nitrogen
because it is burned during helium shell flashes. On the other hand, AGB
models of higher mass convert the dredged-up carbon into nitrogen by CN
cycling at the bottom of the convective envelope (hot bottom burning, HBB)
(Chiappini et al., 2005; Hirschi, 2007; Ekström et al., 2008; Izzard et al.,
2009; Joggerst et al., 2010).

Metal poor stars reside primarily in the halo system of the Milky Way, a
complex old component of the Galaxy that comprises at least two diffuse
stellar populations, the inner- and the outer-halo, with different
metallicities, kinematics and spatial density profiles (Carollo et al., 2007,
2010; Beers et al., 2012), several streams and overdensities (Grillmair,
2009), and a recently discovered large structure in the inner region, product
of a past merger event (Helmi et al., 2018). It has been recognized that 15%
$-$ 20% of stars with [Fe/H] $<-$2.0 in the halo system are CEMP and the
fraction increases with declining metallicity becoming $\sim$ 75% at [Fe/H]
$<-$4.0 (see Carollo et al., 2014, and reference therein). There is also
evidence for a significant contrast in the frequency of CEMP stars that are
kinematically assigned to the inner- and outer-halo components (the inner halo
is on-average non rotating, while the outer halo exhibits a significant
retrograde signature). The outer halo exhibits a fraction of CEMP stars twice
the inner halo in the metallicity interval $-$2.5 $<$[Fe/H] $-$2.0 (Carollo et
al., 2012). Such increase in frequency of CEMP stars can be explained as a
population driven effect, due to the fact that the outer halo is the dominant
component at large distance from the galactic plane and at metallicities,
[Fe/H] $<-$2.0. The chemical differences between the two stellar halo
populations were established also in terms of CEMP sub-classes by Carollo et
al. (2014). It was shown that the relative numbers of CEMP-no stars compared
to CEMP-s varies between the inner- and outer-halo and the frequency of the
CEMP-no stars is higher in the outer halo, while the frequency of the CEMP-s
stars is higher in the inner halo. The analyses of kinematics and dynamics of
our sample stars will establish their inner/outer halo membership and on their
origin.

The chemical abundance analysis of metal-poor stars and searching of CEMP and
NEMP stars provides important constrains on the chemistry evolution of the
Galaxy, initial mass function (IMF; Hirano et al. 2014) and the models of
mass-transfer and evolution of components in binary systems. In this paper, we
report on the discovery of five CEMP stars, which add impact constraints on
different stellar and Galactic chemical-evolution scenarios, as well as the
nature of their progenitors.

This paper is outlined as follows: target selection and observations are
discussed in Section 2. Section 3 discusses the determination of stellar
parameters. Our abundance analysis is discussed in Section 4. We discuss our
results, chemical peculiarities, and kinematics in Section 5, and our
conclusions are given in Section 6.

##  2 Target Selection, Observations and Data Reduction

###  2.1 LAMOST database

Our sample was selected from the Large Sky Area Multi-Object Fiber
Spectroscopic Telescope survey (Zhao et al., 2006, 2012; Cui et al., 2012),
where we ran two different methods independently in order to estimate the
total metallicity [Fe/H] of an object. These two methods can be highlighted by
the following points:

  * •

Matching synthesized spectra with the observed data, and calculate the
corresponding Lick indices to find the best parameters 111See
http://astro.wsu.edu/worthey/html/index.table.html.

  * •

Matching observed normalized spectra with synthesized spectra using a
minimization technique to carry out the best-fit for metallicity.

After running these two methods, we identified a star as very metal-poor
(hereafter VMP) candidate, if these two methods yield [Fe/H] $\leq-2.7$,
assuming a typical uncertainty of deriving metallicities from low-resolution
spectra ($0.1$-$0.3$ dex).

###  2.2 High resolution spectroscopy

We obtain high-resolution spectra (R = 110,000 and a slit width of 0.5 mm),
covering the wavelength range of ($3730$-$9989$ Å ), for a sample of 12 stars,
using the 2.4m Lick Automated Planet Finder (thereafter Lick/APF). In
addition, we observed the well studied VMP star HD2796 as a standard. For more
information about this telescope and its instruments, we refer the reader to
Radovan et al. (2014). After removing the dispersion caused by atmospheric
refraction 222This step was achieved using a trombone-style atmospheric
dispersion corrector (ADC), we carried out a standard echelle data reduction
(bias subtraction, flat- fielding, background subtraction, extraction,
wavelength calibration and continuum-normalization) to obtain 1D normalized
spectra. During our analysis we found that five stars showed high carbon
abundances. In this paper, we present these five stars and in a future paper
we will present the remaining sample. Table 1 shows the observational details
of our sample. The seventh column lists signal-to-noise ratios (S/N), measured
using IRAF splot task at $\sim$ 4500Å . For the radial velocity (RV)
measurements, we built a routine to estimate the heliocentric corrections of
our observations and cross-correlate our spectra against synthesized
templates, with the same spectral type of each star, using strong features
(e.g., Mg I triplet). In addition, we checked the validity of our method using
some metal-poor radial velocity standards. These values are also listed in the
last column of Table 1.

###  2.3 Equivalent Widths

We adopted an atomic line-list from Aoki et al. (2013) and Frebel et al.
(2013) (also see Section 4.3), we used Gaussian profiles to measure the
equivalent widths (thereafter EWs) of those isolated atomic lines 333Gaussian
profiles can not fit strong lines very well, but still lines with EWs $\sim$
100mÅ are still gaussian-like shapes.. Section 4 describes in more details our
abundance measurement method. The measured EWs are listed in Table 2.

In order to assess the quality of our EWs measurements we compared our EWs of
HD2796 with Cayrel et al. (2004) and Li et al. (2015), shown in Figure 1. Our
EWs have good agreement with Cayrel’s results ($R^{2}=0.9806$, zero-point
shift $-0.107$ and $\sigma$ of $3.1$ mÅ ), and Li’s results ($R^{2}=0.9879$,
zero-point shift +1.0234 and $\sigma$ = 2.1 mÅ ). The upper panel of Figure 1
shows a direct comparison of our EWs and their measured values for the common
lines, and one-to-one line is used as a reference. The lower panel shows very
small residuals of the EWs (this work - literature) for those $100$ mÅ
$>$EW$>10$ mÅ .

Neglecting the uncertainties of our continuum definition and based on Cayrel
(1988) formula (see Eq. 1), we expected our uncertainties to be around 2.5 mÅ
. However, one should note that those values are for the weakest detected
lines in these spectra. Where, $FWHM$ is the full width at half maximum, $S/N$
is the signal-to-noise ratio and $\delta x$ is the pixel size.

$\sigma_{EW}=\frac{1.5}{S/N}~{}\sqrt{FWHM*\delta x}$ (1)

Figure 1: Comparison of the equivalent widths of HD2796, the blue circles
represent the comparison of this work and Cayrel et al. (2004) ,while the
black triangles represent the comparison of this this work and Li et al.
(2015). Note that the biggest discrepancies between our EWs and the literature
data are for the strongest lines. Thus, the spectrum synthesis approach have
been used for these lines (e.g., Mg I tripletsg).

##  3 Stellar parameters

Despite the fact that iron abundances derived from both excitation levels (Fe
I and Fe II) are affected by uncertainties in the model atmospheres
temperature and NLTE effects (see Mashonkina et al., 2017, and reference
therein), many researches still rely on this method (the so-called traditional
spectroscopic method) to derive their stellar atmospheric parameters. Our
adopted stellar parameters have been determined through this standard
spectroscopic method. Additionally, we were able to determine effective
temperatures from photometry and surface gravities from parallax/distances.

###  3.1 Effective temperature

The effective temperatures ($T_{\rm eff}$) were estimated by minimizing the
trend between Fe I lines abundances and excitation potential ($\chi$).

To estimate $T_{\rm eff}$ from the available colors (B,V,J,H and K), we cross
match our sample with two catalogs from the Virtual Observatory: The fourth US
Naval Observatory CCD Astrograph Catalog (UCAC4, Zacharias et al. 2013 ) and
Two Micron All Sky Survey (2MASS, Skrutskie et al. 2006) and then employ
Ramírez & Meléndez (2005) temperature calibration. These $T_{\rm eff}$
(determined from photometric and spectroscopic methods) together with $T_{\rm
eff}$ taken from Gaia DR2 (Gaia Collaboration et al., 2018) are in good
agreement with each other ($\pm 150$ K).

###  3.2 Surface gravity and microturbulence

We determined the surface gravity $\log\,g$) by forcing Fe I abundances to
agree with Fe II abundances. In addition, we crossed match our sample with
Gaia DR2 distances catalogue (Bailer-Jones et al., 2018), to estimate the
distance modulus from the available parallax of Gaia Collaboration et al.
(2018) (see Eq. 2 and Eq. 3).

$\log\frac{g}{g_{\odot}}=\log\frac{M}{M_{\odot}}+4\log\frac{T_{\rm
eff}}{T_{\rm eff}}+0.4(M_{bol}-M_{bol_{\odot}})$ (2)

$M_{bol}=V+BC+5\log\varpi+5$ (3)

Figure 2: Our sample shown in an H-R-Diagram, based on the $T_{\rm eff}$ and
$\log\,g$ determined from the LIKC/APF spectra (see Table 3). Yale-Yonsei 12
$\mathrm{Gyr}$ isochrones with [$\alpha$/Fe] = +0.4 and [Fe/H] = $-$2.5,
$-$3.0, and $-$3.5 from Demarque et al. (2004) overplotted as reference.

where, $M$ is the stellar mass, $M_{bol}$ is the absolute bolometric
magnitude, $V$ is the visual magnitude, $BC$ is the bolometric correction (see
Alonso et al., 1999, Eq. 18), and $\pi$ is the parallax.

The microturbulence velocities ($\xi$) were determined by removing any trend
between Fe I lines abundances with the EWs of those lines.

The $T_{\rm eff}$ determined from the photometric and spectroscopic methods
show systematically different results. Frebel et al. (2013) have presented an
explicit method to adjust the spectroscopic $T_{\rm eff}$. This scheme
increases the $T_{\rm eff}$ determined for cool red-giants up-to several
hundred degrees, on the other hand the $T_{\rm eff}$ determined for main-
sequence stars are mostly unaffected. Motivated by Frebel et al. (2013)
results, the atmospheric stellar parameters ($T_{\rm eff}$, $\log\,g$ and
$\xi$) determined for our sample stars from the spectroscopic method were
considered as initial parameters, and then corrected following the same scheme
presented in Frebel et al. (2013). It’s worthy to note that our parameters
were not determined independently. Thus, this procedure was iterated to
consistency. The derived stellar atmospheric parameters are listed in Table 3.

The corrected spectroscopic surface gravities (in cgs units) versus the
corrected spectroscopic $T_{\rm eff}$ of our programme stars, with 12
$\mathrm{Gyr}$ Yale-Yonsei isochrones as a reference (Demarque et al., 2004)
are shown in Figure 2. The error bars shown in Figure 2 represent $T_{\rm
eff}$ and $\log\,g$ one-sigma errors ($\pm 150$ K and $\pm 0.3$ cgs,
respectively).

##  4 Abundance Analysis

We used only non-blended lines with reliable continuum normalization to
measure the chemical abundances using EWs analysis. However, for the molecular
bands and blended lines, spectral synthesis was used. In other words, our
chemical abundances were done by a mixture of spectrum synthesis and
equivalent width analysis. The LTE abundances for all elements are listed in
Table 4. Moreover, we considered the deviations from LTE for Li I , Na I and
Mg I lines. The LTE and NLTE abundances are presented in Table 5. We adopted
the solar log$\epsilon_{\odot}$(X) from Asplund et al. (2009) to obtain our
final chemical abundances and [X/Fe] ratios.

Figure 3: Portions of the LICK/APF spectra near the CH and CN bands (are
represented as black filled circles). Best-fit of the selected CH and CN are
also shown (red solid line). The dashed lines show the synthetic spectrum with
no carbon in the atmosphere.

###  4.1 LTE and NLTE calculations

We used stellar atmosphere models from 1D ATLAS NEWODF grid of Castelli &
Kurucz (2003). Our LTE abundances were performed using an updated version of
the stellar code MOOG (Sneden, 1973). In this update, a continuous scattering
will be treated as a source function, in other-words the absorption and
scattering will be summed rather than treated as true absorption (Sobeck et
al., 2011).

The departures from LTE in the stellar atmospheres were considered for three
chemical elements, Li, Na, and Mg. The adopted Na I and Mg I-II model atoms
are described in Alexeeva et al. (2014) and Alexeeva et al. (2018),
respectively. To solve the radiative transfer and statistical equilibrium
equations, we used the code detail (Butler & Giddings, 1985) based on the
accelerated $\Lambda$-iteration method (Rybicki & Hummer, 1991). The obtained
departure coefficients, $b_{\rm{i}}$ = $n_{\rm{NLTE}}$ / $n_{\rm{LTE}}$, were
then used by the codes binmag3 (Kochukhov, 2010) and synthV-NLTE (Ryabchikova
et al., 2016) to calculate the synthetic NLTE line profiles. Here,
$n_{\rm{NLTE}}$ and $n_{\rm{LTE}}$ are the statistical equilibrium and thermal
(Saha-Boltzmann) number densities, respectively.

We constructed the Li I model atom in the same manner as it was described in
Lind et al. (2009). The main difference between our model atom and the model
atom of Lind et al. (2009) is the collision excitation recipe. We adopted the
electron collision data from Osorio et al. (2011), while Lind et al. (2009)
used cross-sections for collisional excitation by electrons from Park (1971).
We tested our model with Li-enhanced stars and have found a good agreement
with Lind et al. (2009).

###  4.2 Lithium

Our LICK/APF spectra showed no obvious features for the lithium abundance
determinations, except in J2216$+$2232 spectrum. Therefore, upper-limits were
determined for the rest of the sample stars. The lithium abundance was derived
from the Li I 6707.7 Å resonance line. The line-list was taken from VALD
database, hyperfine structure and isotope structure taken into account with
the data from Sansonetti et al. (1995).

###  4.3 Carbon and Nitrogen

Our carbon abundances were derived from the molecular CH AX band around 4300 Å
The molecular line data for the spectrum synthesis was taken from VALD (Kupka
et al., 1999) database. We determine carbon abundances using the method
described in Alexeeva & Mashonkina (2015), our synthetic flux profiles were
convolved with a profile that combines a rotational broadening and broadening
by macroturbulence with a radial-tangential profile. The most probable
macroturbulence velocity Vmac was varied between 4 and 9 km s-1for different
CH lines.

The nitrogen abundances were estimated from the CN 4215 Å and 6971Å bands,
using a spectrum synthesis approach. The dissociation energy (D0) of CN was
adopted to be 7.65 eV (Bauschlicher et al., 1988). Only in J2114$-$0616
spectrum, we have found visible CN 4215 Å and 6971Å bands, which can be
measured quite reliably. The best fits of some molecular lines in our sample
are shown in Figure. 3.

###  4.4 Light Elements: from Na to Zn

We were able to measure abundances for Na, Mg, Ca, Ti, Sc, V, Cr, Mn, Co, Ni,
and Zn in the LICK/APF spectra. Our sample stars are metal-poor and only
resonance Na I lines at 5889, 5895 Å are available for measurements. The van
der Waals damping constant, C6 = $-$31.6, for these lines was adopted from
solar line-profile fitting (Zhao et al., 2016). The magnesium abundances are
derived from Mg I lines at 4703, 5172, 5183, 5528, 5711 Å (see Figure 4).

Figure 4: An example of the NLTE fittings of Na and Mg in LICK/APF spectra .
The observed spectra are shown by filled circles, for comparison, the LTE
profiles computed with the corresponding NLTE abundances are shown by the blue
dashed curves.  Figure 5: Portions of the LICK/APF spectrum of HD2796 (shown
as blue filled circles) near the lines of Eu II at 4129 Å (top left), Ba II at
6141 Å (top right), Sr II at 4077 Å (bottom left), and Zr II at 4161 Å (bottom
right). Best-fit synthetic spectra are also shown (yellow solid line),
together with abundance variations (magenta dashed line and cyan dash-dotted
line) of $\pm$0.1 dex (for Eu), $\pm$0.1 dex (for Ba), $\pm$0.3 dex (for Sr),
and $\pm$0.5 dex (for Zr).

For the most of our sample stars, Ca abundances were derived from 18 well-
defined Ca I lines, Ti abundances were obtained from 26 Ti I and Ti II lines
in total, Scandium appeared in more than six Sc II lines, while V has only one
reliable line at V I 4379.23 Å .

Cr I has at least 5 lines, while Mn I and Sc II lines appeared in all sample
spectra. In our standard star HD2796 spectrum we were unable to find any
detectable feature for Co I and Ni I, while they appeared in the remaining
stars of our sample. Moreover, J1645$+$4357 spectrum shows no detectable line
for the heaviest element of the iron peak Zn, see Table 4.

###  4.5 Neutron-capture Elements

The strontium (Sr) and barium (Ba) abundances for metal-poor stars are very
important, since they are the most commonly detected neutron-capture elements
whose abundances are measured in the vast majority of the metal-poor stars,
thus these two elements are the key of understanding the nature of the
neutron-capture processes in our Galaxy.

One of our main aims of the LICK/APF observations was to increase the high-
resolution chemical inventory of metal-poor stars. In addition, to these
elements mentioned previously, we also investigated yttrium, zirconium,
lanthanum, cerium, praseodymium, neodymium, samarium, and europium, see Table
4.

Figure 5 shows an illustrative spectrum-synthesis example of the neutron-
capture elements in our standard star HD2796. These elements show very weak
features compared to those elements with Z$\leq 30$, with abundance
uncertainties varying between 0.1 and 0.3 dex, depending on the quality of our
spectrum-synthesis abundance measurements.

###  4.6 Uncertainties of stellar parameters and abundances

In Section 2.3 we discussed the uncertainties arising from the EWs. The
stellar parameters are subject of another uncertainties arising from our
ionization equilibrium process (the traditional spectroscopic method), where
we expect that this process will provide typical internal accuracy in our
estimated surface gravity and microturbulence of 0.1 dex and 0.2 km s-1
respectively. In addition, these uncertainties will affect our abundances
determination. The major uncertainties in our abundance determinations are
related to the estimated effective temperatures (usually in the order of 100
K).

Figure 6: Abundances comparison between this work and Cayrel et al. (2004, for
elements with Z$\leq 30$) and François et al. (2007, for neutron-capture
elements). Excluding Sr abundance, the agreement is highly acceptable. The
dashed line refers to $\Delta$= 0 and the dash-dotted lines refer to
$\Delta=\pm 0.3$ dex.

Table 6 lists these abundance uncertainties for HD2796 (as an example), where
we varied our atmospheric models effective temperature, surface gravity, and
microturbulence by 150K, 0.3dex, and 0.3km s-1 respectively. Figure 6 shows a
comparison between our elemental abundances and Cayrel et al. (2004, for
elements with Z$\leq 30$) and François et al. (2007, for neutron-capture
elements). Apart from Sr (with $\Delta$ [Sr/Fe]=$-0.62$ relative to François
et al.), the differences between our work and their work lie within $\pm 0.3$
dex (dash-dotted line in Figure 6). We impute our scattering to use cooler
temperature, lower surface gravity, and lower S/N ratio than those used in
Cayrel et al. (2004) and François et al. (2007) ($T_{\rm eff}$= 4950 K,
$\log\,g$= 1.5 dex, and S/N ratio = 250-550).

##  5 Results and discussion

###  5.1 Chemical abundance comparison with literature data

We provided stellar parameters and detailed chemical abundances for five
metal-poor red giant stars, reported for the first time using high resolution
spectroscopy. These stars exhibit similar chemical abundance patterns to,
reported in other, very and extremely metal-poor stars (e.g., François et al.,
2007; Yong et al., 2013).

Figure 7: log$\epsilon$ lithium abundances as a function of effective
temperature. Red filled star refers to J2216$+$2232, blue filled stars refer
to the determined upper limits, and small open circles refer to the full
sample of Roederer et al. (2014b). The dotted line shows the predicted
primordial lithium abundance, log$\epsilon_{Li}$ = 2.64 (Spergel et al.,
2007).

Lithium is considered as a key diagnostic, to test and constrain our
understanding of the early Galaxy, of stellar interiors and evolution. Figure
7 illustrates the evolution of lithium as a function of $T_{\rm eff}$ using
the halo star sample from Roederer et al. (2014b, including upper limits). The
dotted line refers to the primordial lithium abundance predicted by the
Standard Big Bang Nucleosynthesis (Spergel et al., 2007). Among the program
stars, we could only detect Li in 2216$+$2232 (shown as red filled star) with
A(Li) = 0.95. For completeness the upper limits for the rest of our sample
stars have been provided (blue filled stars). This is not unexpected, as our
sample stars are red-giants, whose Li content in the outer layers have been
diluted by the canonical extra mixing and the first dredge-up (FDU) process.

Our sample stars exhibit relatively high [C/Fe] ratios, as shown in Figure 8
(left panel), which represents [C/Fe] ratio as a function of luminosity. We
adopted a classification of Aoki et al. (2007), who suggest a scheme that
takes into consideration the nucleosynthesis and mixing effects in giants. We
define the stars that satisfy the following criteria as CEMP stars:
[C/Fe]$\geqslant$0.7 for stars with $\log$(L/L⊙)$\leqslant$2.3 and
[C/Fe]$\geqslant$3.0$-$ $\log$(L/L⊙) for stars with $\log$(L/L⊙)$>$2.3. The
luminosities of our stars were calculated based on the prescription of Aoki et
al. (2007), assuming stellar mass of 0.8 M⊙, following Aoki et al. (2005) and
Ryan et al. (2005). For completeness, and due to the fact that our sample
stars are gaints (see Figure 2), we use the carbon evolutionary correction
described in Placco et al. (2014) to assess whether our sample stars could
indeed be classified as CEMP, this method suggests that carbon levels decrease
as stars evolve into the giant branch phase, due to some level of internal
mixing. As a result, the correction increases the C abundances up to several
dex, which support our claims that these stars are CEMP stars.

Figure 8: Left panel: [C/Fe] versus luminosity for our sample stars. The dash-
dotted line indicates the dividing line between carbon-enhanced and carbon-
normal stars as defined in Aoki et al. (2007). The dashed line corresponds to
[C/Fe] = 1.0. Right panel: [C/Fe] vs. [N/Fe]. The two criteria for NEMP stars
suggested by Pols et al. (2012) are respectively shown in dotted ([N/Fe]
$\geq$ 1.0 and [N/C] $\geq$ 0.5) and dash-dotted lines ([(C+N)/Fe] $>$ 0.9).
The filled red stars refer our sample stars. Non-carbon-enhanced objects
studied by previous works (Gratton et al., 2000; Cayrel et al., 2004; Honda et
al., 2004; Aoki et al., 2005) are shown by triangles up. Metal-poor stars from
Spite et al. (2005) (filled diamonds), and CEMP stars from Aoki et al. (2007)
(crosses) are also plotted for comparison. The candidate to CNEMP objects
J2114-0616 is marked.

With the Aoki et al. (2007) definition of CEMP stars and the carbon
evolutionary correction described in Placco et al. (2014) in mind, Figure 8
(left panel) shows that our program stars are located above the limit. Thus,
we point J1054$+$0528, J1529$+$0804, J1645$+$4357, J2114$-$0616, and
J2216$+$2232 as CEMP stars.

For most of our program stars CN bands are not measurable, we could only
measure N abundance for J2114-0616, which exhibits high nitrogen abundance
with [N/Fe]=1.88, [N/C]$>$0.51, and [(C+N)/Fe]=1.53. Figure 8 (right panel)
shows [C/Fe] as a function of [N/Fe], with the dotted and dash-doted lines
referring to Pols et al. (2012) NEMP stars criteria. We classify J2114$-$0616
as a potential nitrogen-enhanced metal-poor (NEMP) star. Since J2114$-$0616
satisfies both criteria (star with [C/Fe]$\geqslant$1.0 and
[N/C]$\geqslant$0.5), it can be designated as a carbon and nitrogen-enhanced
metal-poor (CNEMP) star.

Moreover, we investigated [Eu/Fe] as a function of [Ba/Fe] to study s-process
and r-process enrichment, under the pretext that J2114$-$0616 shows 0.0$<$
[Ba/Eu] $<$ +0.5, and [Ba/Fe] $>$ 0.5 (see figure 9), we regard it as a
CEMP-r/s star. In addition to the enhancements in both slow (s-) and rapid
(r-) process species, J2114$-$0616 shows high [N/Fe] ratio, along with its
high [C/Fe], which suggest that its peculiar chemical pattern may come from
mass transfer from an AGB companion, before it tuned to a white dwarf.

Figure 9: [Eu/Fe] as a function of [Ba/Fe], the open circles are data take
from saga database (Suda et al., 2008) and the magenta star refer to the
position of J2114$-$0616\.

.

Figure 10: Comparison between the light-element abundances in our sample stars
(red filled star) and the literature C-normal metal-poor stars from Yong et
al. (2013) (blue open circles). Figure 11: Selected neutron-capture elements
abundances in our sample stars (red filled stars), as a function of
metallicity, compared to literature data adopted from François et al. (2007)
(blue open circles). Figure 12: Abundance patterns for the neutron-capture
elements in our sample. The solid line represents the Solar-system s-process
abundance pattern from Burris et al. (2000), scaled to match the observed
abundance of Ba in each star. The dotted line represents the Solar-system
r-process abundance pattern from Burris et al. (2000), scaled to match the
observed abundance of Eu in each star. Note that we only include these stars
with measured Ba and Eu.

Light element distributions in CEMP stars are quite similar to those in non-
carbon-enhanced stars (C-normal). The abundance ratios [X/Fe] as a function of
[Fe/H] of our sample stars (red filled stars) are presented in Figure 10, for
the elements from C through Zn, compared with literature data adopted from
Yong et al. (2013) (blue open circles). In general, the abundance ratios seen
in our sample show good agreement with the abundance ratio trends defined by
the literature sample. On the other hand, J1529+0804, which shows enhancement
in manganese with [Mn/Fe]=0.43, is not a good example of this agreement, this
enhancement may also be true for J1645+4357 ([Mn/Fe]=0.12), keeping in mind
that at low metallicities, the NLTE behavior may systematically increase the
manganese abundance up to 0.7 dex (which we will explore in future work)
(e.g., Bergemann & Gehren, 2008).

###  5.2 Nucleosynthetic signatures of s- and r-process

Only elements lighter than zinc can be produced via nuclear-fusions, on the
other hand, heavier elements can be synthesized by either the rapid neutron
capture process, r-process, and the slow neutron capture process, s-process
(e.g., Meyer 1994; Arnould et al. 2007 and references therein). Metal-poor
stars provide unique opportunities to attain nucleosynthetic signatures, thus
better understanding of the chemical evolution of these elements and the
nucleosynthesis occurred in the early Universe.

As mentioned previously, we were able to determine abundances for up to 10
heavy elements, including the light trans-iron elements ($38\leq$Z$\leq 46$)
and the second r-process peak elements. The abundances of selected neutron-
capture elements for our sample stars (red filled stars), as a function of the
metallicity, overlaid with literature data adopted from François et al. (2007)
(blue open circles) are shown in Figure 11. No significant discrepancies are
found between the selected neutron-capture elements abundances of our sample
stars and the literature data

Figure 12 shows the neutron-capture element abundances for HD2796 and three
sample stars, compared with the Solar System s-process (normalized to Ba -
solid line) and r-process (normalized to Eu - dotted line) components. The s-
and r- fractions were taken from Burris et al. (2000). Abundances for the
first-peak s-process elements (Sr, Y, and Zr) are well described by the Solar
s-process for HD2796 and J2216+2232, while for J1054+0528 and J2114-0616 they
are roughly consistent with the r-process component. At the same time, the
noticeable deviations from for the light elements might interestingly be
related to the effects of core-collapse supernovae. In contrast, for the
second-peak s-process elements, there is an excellent agreement between
measurements and the Solar s-process component for J2114-0616 which, combined
with its enhancements in carbon and nitrogen, supports the hypothesis of mass
transfer in a binary system from an AGB companion. For J1054+0528, it appears
that all the neutron-capture element abundances are a result of an r-process
event, with no contributions from the s-process component.

We could only detect weak Ba II lines in the spectrum of J1645$+$4357, which
results in a relatively low Ba abundance while no other s-process elements can
be detected. Such chemical patterns (high carbon, low barium, and absence of
other s-process elements) suggests that J1645$+$4357 was formed out of
pristine gas. The common definition of strongly r-process-enhanced star is
[Eu/Fe] $>+1.0$ and [Ba/Eu] $<0$ and moderately r-process-enhanced star is
(r-II stars) $+0.3\geq$[Eu/Fe]$\leq+1.0$ and [Ba/Eu] $<0$ (e.g., Frebel,
2018), adopting this criterion we suggest that J1054$+$0528 ([Eu/Fe]=0.44) is
a new member of the moderately r-process-enhanced stars (r-I stars). On the
other hand, J2114$-$0616 exhibits different chemical abundance patterns,
enhancement of s-process species along with relatively high magnesium
abundance suggesting that 22Ne($\alpha$, n)25Mg may have operate as a main
neutron source in J2114$-$0616 (Masseron et al., 2010). Moreover, it turns out
that the neutron density linked to this reaction favors the production of
cerium (81% synthesized by s-process) and europium (97% synthesized by
r-process), suggesting that these elements can’t be explained by s-process
only, and additional r-process is required to describe this behavior (Gallino
et al., 1998; Goriely & Mowlavi, 2000).

###  5.3 Kinematics and dynamics

Figure 13: Total energy vs. angular momentum in the z direction for the stars
in our sample. Top panel: C-normal stars are represented by black stars, while
CEMP stars are denoted by red stars. Bottom panel: the same sample of stars,
with blue filled star symbols indicating CEMP-no members of the inner halo,
the orange filled square symbol represents the CEMP-no and CEMP-s members of
the outer halo. The green dashed curve denotes the locus of the points that
possess constant apo-Galactic radius, rapo = 15 kpc, while the golden dot-
dashed horizontal lines shows the values of the energies delimiting the
transition zone.

The full space motion is derived by combining the observables obtained by Gaia
DR2, positions and proper motions ($\alpha$, $\delta$, $\mu_{\alpha}$,
$\mu_{\delta}$). We utilize the software TOPCAT to cross match our sample with
two catalogs from the Virtual Observatory: Gaia DR2 (Gaia Collaboration et
al., 2018, for proper motions) and Gaia DR2 distances (Bailer-Jones et al.,
2018, for distances), see Table 7. Radial velocities are obtained through
cross-correlation with synthetic spectra after the heliocentric corrections to
the observed spectra are applied. The velocities calculated in the Local
Standard of Rest (LSR) are referred to as $(U,V,W)$ which are corrected for
the motion of the Sun by adopting the values ($U,V,W$) = ($-$9,12,7) km s-1
(Mihalas & Binney, 1981). The velocity component $U$ is taken to be positive
in the direction towards the Galactic anti-centre, the $V$ component is
positive in the direction towards Galactic rotation, and the $W$ component is
positive toward the north Galactic pole. We also compute the rotational
velocity component about the Galactic centre in a cylindrical frame, denoted
as Vϕ , and is calculated assuming that the LSR is on a circular orbit with a
value of 220 km s-1 (Kerr & Lynden-Bell, 1986). The orbital parameters are
derived by adopting a Stäckel type gravitational potential (which consists of
a flattened, oblate disk, and a nearly spherical massive dark-matter halo; a
complete description is given by Chiba & Beers (2000, Appendix A) and
integrating their orbital paths based on the starting point obtained from the
observations.
In addition, we evaluate the integrals of motion for any given orbit, deriving
the energy, E, and the angular momentum in the vertical direction, LZ = R x
Vϕ. Note that R represents the distance from the Galactic center projected
onto the disk plane. Typical errors on the orbital parameters (at Zmax $<$ 50
kpc; Carollo et al., 2010) are: $\sigma$rperi $\sim$ 1 kpc, $\sigma$rapo
$\sim$ 2 kpc, $\sigma$ecc $\sim$ 0.1, $\sigma$Zmax $\sim$ 1 kpc.
Carollo et al. (2014) established a method for assigning the membership to the
inner- and outer-halo stellar populations based on the integrals of motion
(total energy and vertical angular momentum) of a large sample of SDSS/SEGUE
DR7 calibration stars. Inner halo stars are mostly highly bound to the Galaxy
(lower energy values, $E$ $<$ $-$1.1 km2 s-2) and possess orbits with apo-
galactic distance r${}_{\rm apo}<$ 15 kpc, while outer halo stars are less
bound to the Galaxy (higher energy values, $E$ $>$ $-$0.9 km2 s-2 ) and
possess orbits with r${}_{\rm apo}>$ 15 kpc. Stars with r${}_{\rm apo}>$ 15
kpc and $E$ $<$ $-$1.1 km2 s-2 can be also considered pure inner halo stars.
In general, stars in the outer halo are dominated by retrograde orbits but can
also possess rotational velocities less retrograde or higlhy prograde, due to
the large velocity dispersion of the outer halo ($\sim$ 165 km s-1; Carollo et
al., 2010). This is clearly evident in the right panel of Figure 4 in (Carollo
et al., 2014) .
Figure 13 shows the total energy, E, as a function of the angular momentum in
the vertical direction, LZ, for the program stars. In the top panel, the black
filled star symbol represents HD2796, while the red filled star symbols denote
the CEMP stars. The grey horizontal area shows the range of binding energy
values defining the transition zone between the inner- and the outer-halo
components ($-$1.1 km2 s-2 $<E$ $<$ $-$0.9 km2 s-2), which is defined as the
energy range where stars have similar probability to be members of these
components. The green dashed curve represents the locus of stars possessing
orbits with constant apo-galactic radius rapo = 15 kpc. In the bottom panel
the magenta star symbols denote the CEMP-no stars in the inner halo,
classified according to their value of binding energy and apo-galactic
distance, the CEMP-r/s star (J2114$-$0616) is represented by an orange filled
square and it is member of the outer halo.
Figure 14 shows the galactocentric rotational velocity as a function of the
metallicity for the program stars. It is interesting to note that J2114$-$0616
possesses a prograde motion (rotate in the same direction of the galactic
disk) with velocities within $\sim$ 2$\sigma$ (CEMP-r/s; J2114$-$0616) of the
mean rotational velocity of the outer halo population ($-$80 km s-1). Highly
prograde stars in the outer halo were also found in the sample of CEMP stars
reported in Carollo et al. (2014, ; Figure 4).
Numerical cosmological simulations of MW-mass galaxies predict that stars in
the inner halo of the MW formed mainly from massive subgalactic fragments that
experienced an extended star formation activity (Zolotov et al., 2009; Font et
al., 2011; McCarthy et al., 2012; Tissera et al., 2013, 2014), while outer
halo stars formed predominantly in lower-mass subgalactic fragments with short
or truncated star formation history (Carollo et al., 2007, 2010; Beers et al.,
2012; Tissera et al., 2013, 2014; Carollo et al., 2016, 2018). The central
regions of simulated halos (within $\sim$ 15 kpc) have an important
contribution of in-situ stars (formed in the main progenitor galaxy) which
have various possible origins (Brook et al., 2004; Zolotov et al., 2009; Font
et al., 2011; House et al., 2011; Tissera et al., 2013; Cooper et al., 2015;
Pillepich et al., 2015; Monachesi et al., 2016). On the contrary, stars in the
outer halo formed primarily in low-mass subgalactic systems which were
subsequently accreted. The origin of halo stars can be understood by
inspecting a combination of their orbital parameters and integrals of motion.
In case of our sample, J2114$-$0616, possesses orbital parameters, energy and
vertical angular momentum that place it in the outer halo population and it
likely were formed in low-mass systems outside the virial radius of the
progenitor galaxy and accreted later on. The orbital parameters and binding
energy of the remaining CEMP stars, J1054+0528, J1529+0804, J1645+4357 and
J2216+2232, suggest that they are members of the inner halo population.
However, their metallicity and C-enhancement indicate that they may have
formed not in situ but in small mass subgalactic fragments which were accreted
very early on and contributed to the old central regions of the halo system
(Tissera et al., 2018; Carollo et al., 2018).

Figure 14: Galactocentric rotational velocity vs metallicity for the program
stars color coded as in Figure 12.

##  6 conclusion

In this work we reported on the discovery of five CEMP stars selected from the
LAMOST DR3 database as metal-poor candidates. High-resolution spectra are
obtained for the first time with APF /LICK. We confirmed that J1054$+$0528,
J1529$+$0804, J1645$+$4357, J2216$+$2232 and J2114$-$0616 show high
enhancement in carbon with [C/Fe]= 0.82, 0.80, 0.70, 0.62, and 1.37,
respectively, taking an advantage of their observed high-resolution spectra
and the correction for evolutionary mixing. We provided stellar parameters and
chemical abundances for up to 25 elements: three light elements (Li, C, N), 12
elements from Na to Zn, 10 neutron-capture elements (Sr, Y, Zr, Ba, La, Ce,
Pr, Nd, Sm, and Eu). Our results show no significant abundance differences
with literature, thus can be used to study the chemical enrichment at the
earliest time of the Galaxy.

J1054+0528 shows moderate enhancement in europium [Eu/Fe]=0.44, with low
barium-to-iron ratio [Ba/Fe]=$-$0.52. These abundances indicate that this star
can be classified as CEMP-rI and future higher signal-to-noise observations
should be carried out to obtain abundances of other r-process elements, such
as Gd, Th, Os, Pt, and Pb.

No reliable s-process elements lines were found in J1645$+$4357 spectrum. More
accurate measurements of more heavy elements would be needed for further
deduction. Thus, we strongly recommend future observation of J1645$+$4357 in
the near-ultraviolet spectroscopy, to determine many key elements such as Ge,
Zr, Os, Pt, and Pb.

J2114$-$0616 can be classified as CNEMP-r/s (carbon nitrogen-enhanced metal-
poor) star, due to the high enhancement in nitrogen with [N/Fe]= 1.88, [N/C]=
0.5, barium [Ba/Fe]=1.00, and europium [Eu/Fe]=0.84. Other r-process elements,
such as Gd, Th, Os, Pt, and Pb can be obtained in future high resolution and
high signal-to-noise observations.

Kinematic and dynamic analysis based on Gaia DR2 parameters reveals that
J2114$-$0616 is stellar members of the outer halo population and suggests that
it may have formed in low-mass systems outside the virial radius of the
progenitor galaxy and accreted later on during the merging and disruption of
such systems. The stars J1054+0528, J1529+0804, J1645+4357 and J2216+2232 are
stellar members of the inner halo population, however their very low
metallicity and enhancement in carbon abundance strongly indicate that they
were born in low-mass sub-galactic systems which were accreted during the
initial phases of Galaxy assembly and contributed to the old stellar
populations of the inner halo.

## Acknowledgement

M. K. M. would like to thank The World Academy of Sciences and the Chinese
Academy of Sciences for the CAS-TWAS fellowship. M. K. M. thanks Ian Roederer,
Richard de Grijs, and Rene A. Mendez for their comments on earlier versions of
the manuscript, and the anonymous referee, who made valuable suggestions that
helped improve the paper. V. M. P. acknowledges partial support for this work
from grant PHY 1430152; Physics Frontier Center/ JINA Center for the Evolution
of the Elements (JINA-CEE), awarded by the US National Science Foundation
(NSF). S.A. is grateful to China Postdoctoral international exchange program
(ISS-SDU) for financial support (Weihai, China). This study is supported by
the National Natural Sciences Foundation of China under grant No. $11890694$,
$11573032$, and $11390371$.

.

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Table 1: Log of the Lick/APF observations.

| ID | Date | RA | DEC | $r$ | Exptime | S/N | $V_{r}$
---|---|---|---|---|---|---|---|---
|  |  |  |  | (mag) | (s) | (pixel${}^{-1})$ | (km s-1)
1 | HD2796 | 18 Nov 2015 | 00 31 16.91 |  $-$16 47 40.8 | 8.51 | 900*2 | 43 |  $-$60.51
2 | J1054$+$0528 | 28 May 2015 | 10 54 33.10 |  $+$05 28 12.7 | 12.65 | 1800*4 | 42 | 82.36
3 | J1529$+$0804 | 30 May 2015 | 15 29 53.94 |  $+$08 04 48.1 | 12.49 | 1800*4 | 46 | 23.40
4 | J1645$+$4357 | 28 May 2015 | 16 45 14.95 |  $+$43 57 12.0 | 12.79 | 1800*4 | 40 |  $-$83.20
5 | J2114$-$0616 | 23 Sep 2015 | 21 14 01.52 |  $-$06 16 10.3 | 10.81 | 1200*4 | 46 |  $-$160.32
6 | J2216$+$2232 | 27 July 2015 | 22 16 39.31 |  $+$22 32 50.4 | 12.03 | 1800*4 | 47 |  $-$339.38

Note. — The S/N ratio per pixel was measured at $\lambda\sim 4500$ Å .

Luminosity was derived based on Aoki et al. (2007) relation

Table 2: Equivalent widths of our sample.

$\lambda$ | Species | $\chi$ | $\log\,gf$ | HD2796 | J1054+0528 | J1529+0804 | J1645+4357 | 2114$-$0616 | J2216+2232
---|---|---|---|---|---|---|---|---|---
(Å ) |  | eV |  |  mÅ |  mÅ |  mÅ |  mÅ |  mÅ |  mÅ
4049 | C(CH) | — | — | — | syn | — | — | syn | —
4218 | C(CH) | — | — | syn | syn | — | — | syn | syn
4248 | C(CH) | — | — | — | — | syn | syn | syn | —
4253 | C(CH) | — | — | syn | — | syn | — | syn | syn
4261 | C(CH) | — | — | — | — | — | syn | — | —
4273 | C(CH) | — | — | — | syn | — | syn | syn | syn
4279 | C(CH) | — | — | — | — | — | syn | syn | syn
4281 | C(CH) | — | — | syn | — | — | — | syn | syn
4292 | C(CH) | — | — | — | syn | syn | — | syn | —
4310 | C(CH) | — | — | — | — | syn | — | — | syn
4313 | C(CH) | — | — | syn | — | — | syn | syn | —
4362 | C(CH) | — | — | — | syn | — | — | — | syn
4366 | C(CH) | — | — | — | — | — | syn | — | syn
4214 | N(CN) | — | — | — | — | — | — | syn | —
6970 | N(CN) | — | — | — | — | — | — | syn | —
5889.95 | Na I | 0.00 | 0.10 | 169.6 | 118.7 | 66.4 | 174.2 | 195.0 | 140.0
5895.92 | Na I | 0.00 | -0.20 | 164.8 | 120.4 | 65.1 | 161.1 | 181.8 | 159.4
4702.99 | Mg I | 4.33 | -0.44 | 77.2 | 26.1 | 8.8 | 64.5 | 73.0 | 51.3
5172.68 | Mg I | 2.71 | -0.45 | 217.9 | 124.4 | 209.7 | 211.0 | 183.6 |
5183.60 | Mg I | 2.72 | -0.24 | 260.2 | 145.6 | 147.8 | 232.1 | 267.0 | 218.5
5528.40 | Mg I | 4.35 | -0.50 | 68.7 | 36.2 | 33.7 | 72.7 | 81.3 | 48.0
5711.09 | Mg I | 4.35 | -1.72 | 3.5 | 10.5 | — | — | 9.7 | 5.3
4283.01 | Ca I | 1.89 | -0.22 | 60.9 | — | 15.2 | 54.9 | 78.5 | 59.9
4318.65 | Ca I | 1.89 | -0.21 | 64.7 | — | — | — | 22.2 | 40.0
4425.44 | Ca I | 1.88 | -0.36 | 50.9 | 27.2 | 19.3 | — | 52.9 | —
4454.78 | Ca I | 1.90 | 0.26 | 81.9 | 31.9 | 46.5 | 42.1 | 76.4 | 64.9
4455.89 | Ca I | 1.90 | -0.53 | 43.9 | — | 18.8 | — | 38.7 | 15.9

Note. — Table 2 is published in its entirety in the machine-readable format. A
portion is shown here for guidance regarding its form and content.

Table 3: Stellar Parameters of the Program Stars.

| Lick/APF (adopted) | LAMOST |  | Gaia DR2 |  | Photometry |  |  | Luminosity |  |
---|---|---|---|---|---|---|---|---|---|---|---
ID | $T_{\rm eff}$ | $\log\,g$ | [Fe/H] | $\xi$ |  | $T_{\rm eff}$ | $\log\,g$ | [Fe/H] |  |  | $T_{\rm eff}$ | $\log\,g$ |  |  |  | $T_{\rm eff}$ |  |  | $~{}~{}L_{\ast}/L_{\odot}$ |  |
| (K) | (cgs) |  | (kms-1) |  | (K) | (cgs) |  |  |  | (K) | (cgs) |  |  |  | (K) |  |  |  (ergs-1) |  |
HD2796 | 4869 | 1.04 |  $-$2.51 | 2.01 |  | … | … | … |  |  | 4995 | 2.15 |  |  |  |  … |  |  |  1014.0 |  |
J1054$+$0528 | 5030 | 1.88 |  $-$3.30 | 1.94 |  | 5094 | 1.76 |  $-$3.23 |  |  | 5080 | 1.89 |  |  |  |  4981 |  |  |  166.94 |  |
J1529$+$0804 | 5085 | 2.00 |  $-$3.18 | 2.34 |  | 5026 | 1.54 |  $-$3.27 |  |  | 5041 | 1.90 |  |  |  |  4913 |  |  |  132.27 |  |
J1645$+$4357 | 4810 | 1.39 |  $-$2.97 | 2.93 |  | 4715 | 2.14 |  $-$3.05 |  |  | 4886 | 1.39 |  |  |  |  4652 |  |  |  431.38 |  |
J2114$-$0616 | 4999 | 1.48 |  $-$2.43 | 2.11 |  | 4377 | 1.55 |  $-$2.95 |  |  | 4870 | 1.66 |  |  |  |  4831 |  |  |  409.09 |  |
J2216$+$2232 | 4842 | 1.40 |  $-$2.91 | 1.79 |  | 4902 | 2.65 |  $-$3.13 |  |  | 5018 | 2.05 |  |  |  |  4815 |  |  |  432.89 |  |

Note. — The tenth column was measured using parallaxes adopted from Gaia DR2
and stellar mass of $\sim$ 0.8 M⊙

Table 4: LTE Abundances of Individual Elements for the Program Stars.

| HD2796 |  | J1054$+$0528 |  | J1529$+$0804 |  |
---|---|---|---|---|---|---|---
| log $\epsilon$(X) | [X/Fe] | $\sigma$ | $N$ |  | log $\epsilon$(X) | [X/Fe] | $\sigma$ | $N$ |  | log $\epsilon$(X) | [X/Fe] | $\sigma$ | $N$ |  |
Li I | … | … | … | … |  | … | … | … | … |  | … | … | … | … |  |
C(CH) | 5.60 |  $-$0.31 | 0.05 | 5 |  | 5.95 | 0.82 | 0.09 | 6 |  | 6.05 | 0.80 | 0.11 | 4 |  |
(CH)corr | … | 0.74 | … | … |  | … | 0.04 | … | … |  | … | 0.01 | … | … |  |
N(CN) | … | … | … | … |  | … | … | … | … |  | … | … | … | … |  |
Na I | 4.40 | 0.67 | 0.05 | 2 |  | 3.25 | 0.31 | 0.05 | 2 |  | 2.76 |  $-$0.30 | 0.24 | 2 |  |
Mg I | 5.68 | 0.59 | 0.05 | 4 |  | 4.71 | 0.41 | 0.12 | 5 |  | 4.64 | 0.22 | 0.10 | 4 |  |
Ca I | 4.15 | 0.32 | 0.10 | 21 |  | 3.39 | 0.35 | 0.29 | 11 |  | 3.46 | 0.30 | 0.18 | 17 |  |
Sc II | 0.55 |  $-$0.09 | 0.05 | 13 |  |  $-$0.12 | 0.03 | 0.24 | 6 |  |  $-$0.10 |  $-$0.07 | 0.05 | 3 |  |
Ti I | 2.64 | 0.20 | 0.09 | 18 |  | 2.08 | 0.43 | 0.25 | 13 |  | 1.87 | 0.10 | 0.30 | 11 |  |
Ti II | 2.67 | 0.23 | 0.12 | 33 |  | 1.84 | 0.19 | 0.30 | 25 |  | 1.95 | 0.18 | 0.32 | 24 |  |
V I | 1.36 |  $-$0.06 | 0.11 | 1 |  | 0.93 | 0.30 | 0.13 | 1 |  | … | … | … | … |  |
Cr I | 2.97 |  $-$0.16 | 0.10 | 9 |  | 2.34 | 0.00 | 0.30 | 5 |  | 2.27 |  $-$0.19 | 0.29 | 8 |  |
Mn I | 2.57 |  $-$0.35 | 0.12 | 3 |  | 1.67 |  $-$0.46 | 0.10 | 1 |  | 2.68 | 0.43 | 0.02 | 2 |  |
Fe I | 4.99 | 0.00 | 0.12 | 217 |  | 4.20 | 0.00 | 0.16 | 84 |  | 4.32 | 0.00 | 0.26 | 146 |  |
Fe II | 4.99 | 0.00 | 0.09 | 24 |  | 4.20 | 0.00 | 0.18 | 8 |  | 4.32 | 0.00 | 0.25 | 17 |  |
Co I | … | … | … | … |  | 1.61 |  $-$0.08 | 0.20 | 2 |  | 2.29 | 0.48 | 0.06 | 3 |  |
Ni I | … | … | … | … |  | 3.10 | 0.18 | 0.13 | 4 |  | 3.66 | 0.62 | 0.32 | 7 |  |
Zn I | 2.33 | 0.28 | 0.03 | 2 |  | … | … | … | … |  | 1.95 | 0.57 | 0.12 | 1 |  |
Sr II |  $-$0.01 |  $-$0.37 | 0.10 | 1 |  |  $-$0.73 |  $-$0.30 | 0.13 | 1 |  | … | … | … | … |  |
Y II |  $-$0.66 |  $-$0.36 | 0.06 | 3 |  |  $-$0.81 | 0.28 | 0.16 | 1 |  |  $-$1.04 |  $-$0.07 | 0.19 | 2 |  |
Zr II | 0.19 | 0.12 | 0.12 | 1 |  | … | … | … | … |  | … | … | … | … |  |
Ba II |  $-$0.57 |  $-$0.24 | 0.22 | 3 |  |  $-$1.64 |  $-$0.52 | 0.10 | 3 |  |  $-$2.01 |  $-$1.01 | 0.13 | 2 |  |
La II |  $-$1.69 |  $-$0.28 | 0.04 | 2 |  | … | … | … | … |  |  $-$1.63 | 0.45 | 0.13 | 1 |  |
Ce II |  $-$1.18 |  $-$0.25 | 0.10 | 3 |  |  $-$0.74 | 0.98 | 0.20 | 2 |  |  $-$0.77 | 0.84 | 0.11 | 2 |  |
Pr II |  $-$1.36 | 0.43 | 0.17 | 1 |  | … | … | … | … |  |  $-$0.88 | 1.58 | 0.19 | 1 |  |
Nd II |  $-$1.34 |  $-$0.25 | 0.26 | 7 |  |  $-$0.91 | 0.97 | 0.19 | 9 |  |  $-$0.71 | 1.05 | 0.58 | 5 |  |
Sm II | … | … | … | … |  | … | … | … | … |  |  $-$0.60 | 1.62 | 0.05 | 2 |  |
Eu II |  $-$2.07 |  $-$0.08 | 0.11 | 1 |  |  $-$2.34 | 0.44 | 0.14 | 1 |  | … | … | … | … |  |

Table 4: Continued.

| J1645$+$4357 |  | J2114$-$0616 |  | J2216$+$2232 |  | Sun |  |
---|---|---|---|---|---|---|---|---|---
| log $\epsilon$(X) | [X/Fe] | $\sigma$ | $N$ |  | log $\epsilon$(X) | [X/Fe] | $\sigma$ | $N$ |  | log $\epsilon$(X) | [X/Fe] | $\sigma$ | $N$ |  |  | log $\epsilon$(X) |  |  |  |
Li I | … | … | … | … |  | … | … | … | … |  | 0.96 | 2.82 | 0.04 | 1 |  |  | 1.05 |  |  |  |
C(CH) | 6.16 | 0.70 | 0.04 | 6 |  | 7.37 | 1.37 | 0.05 | 10 |  | 6.14 | 0.62 | 0.05 | 8 |  |  | 8.43 |  |  |  |
(CH)corr | … | 0.45 | … | … |  | … | 0.17 | … | … |  | … | 0.45 | … | … |  |  | … |  |  |  |
N(CN) | … | … | … | … |  | 7.28 | 1.88 | 0.17 | 2 |  | … | … | … | … |  |  | 7.83 |  |  |  |
Na I | 3.45 | 0.18 | 0.04 | 2 |  | 4.69 | 0.88 | 0.05 | 2 |  | 3.68 | 0.35 | 0.05 | 2 |  |  | 6.24 |  |  |  |
Mg I | 5.12 | 0.49 | 0.05 | 4 |  | 5.69 | 0.52 | 0.12 | 5 |  | 5.34 | 0.65 | 0.02 | 4 |  |  | 7.60 |  |  |  |
Ca I | 3.44 | 0.07 | 0.28 | 11 |  | 4.17 | 0.26 | 0.29 | 20 |  | 3.70 | 0.27 | 0.17 | 18 |  |  | 6.34 |  |  |  |
Sc II |  $-$0.48 |  $-$0.66 | 0.23 | 4 |  | 0.75 | 0.04 | 0.25 | 13 |  | 0.22 |  $-$0.02 | 0.14 | 8 |  |  | 3.15 |  |  |  |
Ti I | 1.88 |  $-$0.10 | 0.26 | 11 |  | 2.78 | 0.26 | 0.16 | 15 |  | 2.26 | 0.22 | 0.13 | 15 |  |  | 4.95 |  |  |  |
Ti II | 1.76 |  $-$0.21 | 0.25 | 25 |  | 2.65 | 0.13 | 0.31 | 30 |  | 2.05 | 0.01 | 0.27 | 26 |  |  | 4.95 |  |  |  |
V I | … | … | … | … |  | … | … | … | … |  | … | … | … | … |  |  | 3.93 |  |  |  |
Cr I | 2.24 |  $-$0.43 | 0.30 | 6 |  | 3.13 |  $-$0.08 | 0.09 | 8 |  | 2.35 |  $-$0.38 | 0.22 | 10 |  |  | 5.64 |  |  |  |
Mn I | 2.58 | 0.12 | 0.08 | 2 |  | 2.49 |  $-$0.51 | 0.03 | 2 |  | 2.24 |  $-$0.28 | 0.21 | 2 |  |  | 5.43 |  |  |  |
Fe I | 4.53 | 0.00 | 0.22 | 150 |  | 5.07 | 0.00 | 0.14 | 156 |  | 4.58 | 0.00 | 0.18 | 181 |  |  | 7.50 |  |  |  |
Fe II | 4.53 | 0.00 | 0.20 | 18 |  | 5.07 | 0.00 | 0.17 | 24 |  | 4.58 | 0.00 | 0.19 | 20 |  |  | 7.50 |  |  |  |
Co I | 2.04 | 0.02 | 0.09 | 1 |  | 2.51 |  $-$0.05 | 0.03 | 2 |  | 2.27 | 0.19 | 0.11 | 3 |  |  | 4.99 |  |  |  |
Ni I | 3.60 | 0.35 | 0.16 | 8 |  | 3.85 | 0.06 | 0.10 | 11 |  | 3.40 | 0.09 | 0.15 | 9 |  |  | 6.22 |  |  |  |
Zn I | … | … | … | … |  | 2.43 | 0.30 | 0.12 | 2 |  | 1.92 | 0.27 | 0.04 | 2 |  |  | 4.56 |  |  |  |
Sr II | … | … | … | … |  | … | … | … | … |  |  $-$0.66 |  $-$0.62 | 0.08 | 1 |  |  | 2.87 |  |  |  |
Y II | … | … | … | … |  |  $-$0.12 | 0.10 | 0.03 | 2 |  |  $-$0.89 |  $-$0.19 | 0.09 | 2 |  |  | 2.21 |  |  |  |
Zr II | … | … | … | … |  | 0.79 | 0.64 | 0.13 | 1 |  |  $-$0.37 |  $-$0.04 | 0.10 | 1 |  |  | 2.58 |  |  |  |
Ba II |  $-$1.80 |  $-$1.01 | 0.15 | 3 |  | 0.75 | 1.00 | 0.11 | 3 |  |  $-$1.35 |  $-$0.62 | 0.25 | 3 |  |  | 2.18 |  |  |  |
La II |  $-$2.25 |  $-$0.38 | 0.11 | 1 |  |  $-$0.17 | 1.16 | 0.02 | 2 |  | … | … | … | … |  |  | 1.10 |  |  |  |
Ce II | … | … | … | … |  | 0.08 | 0.93 | 0.12 | 3 |  | … | … | … | … |  |  | 1.58 |  |  |  |
Pr II | … | … | … | … |  |  $-$0.49 | 1.22 | 0.15 | 1 |  | … | … | … | … |  |  | 0.72 |  |  |  |
Nd II | … | … | … | … |  | 0.10 | 1.11 | 0.17 | 8 |  |  $-$1.46 | 0.03 | 0.22 | 5 |  |  | 1.42 |  |  |  |
Sm II | … | … | … | … |  |  $-$0.18 | 1.29 | 0.12 | 1 |  | … | … | … | … |  |  | 0.96 |  |  |  |
Eu II | … | … | … | … |  |  $-$1.07 | 0.84 | 0.13 | 1 |  |  $-$2.84 |  $-$0.44 | 0.12 | 1 |  |  | 1.07 |  |  |  |

Note. — N refers to the number of lines adopted for determination of the
elemental abundances.

Table 5: Li, Na, Mg NLTE abundances

Star | log$\epsilon$(Li) | log$\epsilon$(Na) | [Na/Fe]NLTE | log$\epsilon$(Mg) | [Mg/Fe]NLTE
---|---|---|---|---|---
| LTE | NLTE | LTE | NLTE |  | LTE | NLTE |
HD2796 |  $<$ 0.14 |  $<$ 0.15 | 4.40 | 3.65 |  $-$0.08 | 5.68 | 5.69 | 0.60
J1054$+$0528 |  … |  … | 3.25 | 2.91 |  $-$0.03 | 4.71 | 4.81 | 0.51
J1529$+$0804 |  $<$ 0.30 |  $<$ 0.32 | 2.76 | 2.53 |  $-$0.53 | 4.64 | 4.75 | 0.33
J1645$+$4357 |  $<$ 0.24 |  $<$ 0.26 | 3.45 | 2.97 |  $-$0.30 | 5.12 | 5.10 | 0.47
J2114$-$0616 |  $<$ 0.44 |  $<$ 0.44 | 4.69 | 3.88 | 0.07 | 5.69 | 5.69 | 0.52
J2216$+$2232 |  0.95 |  0.96 | 3.68 | 3.06 |  $-$0.27 | 5.34 | 5.33 | 0.64

Note. — J1054$+$0528 has defect in the spectrum at the region of the Li I
line.

Table 6: Uncertainties of log $\epsilon$(X) Propagated from the Stellar Parameters (as Described in Section 4.6), Computed for HD2796 as an Example. Ion |  $\Delta$Teff |  $\Delta$$\log\,g$ | $\Delta\xi$
---|---|---|---
|  $+$ 150 K |  $+$ 0.3 dex |  $+$ 0.3 km s-1
CH(C) | 0.32 |  $-$0.12 | 0.02
CN(N) | 0.50 |  $-$0.10 | 0.01
Na I | 0.16 |  $-$0.05 |  $-$0.03
Mg I | 0.14 |  $-$0.11 |  $-$0.03
Ca I | 0.05 |  $-$0.02 |  $-$0.11
Sc II | 0.09 |  $-$0.12 |  $-$0.13
Ti I |  $-$0.03 | 0.00 |  $-$0.02
Ti II | 0.11 | -0.12 | -0.13
V II |  $-$0.04 | 0.00 |  $-$0.02
Cr II |  $-$0.02 |  $-$0.01 |  $-$0.01
Mn I | 0.01 |  $-$0.01 |  $-$0.08
Fe I | 0.16 |  $-$0.03 |  $-$0.15
Fe II | 0.15 |  $-$0.13 |  $-$0.219
Co I |  $-$0.04 | 0.01 | 0.03
Ni I |  $-$0.01 |  $-$0.05 |  $-$0.13
Zn I | 0.06 |  $-$0.06 |  $-$0.18
Sr II | 0.06 |  $-$0.07 | 0.08
Y II | 0.07 |  $-$0.13 |  $-$0.17
Zr I | 0.07 |  $-$0.12 |  $-$0.16
Ba II | 0.02 |  $-$0.12 | 0.08
La II | 0.04 |  $-$0.12 |  $-$0.14
Ce II | 0.05 |  $-$0.12 |  $-$0.15
Pr II | 0.03 |  $-$0.12 |  $-$0.13
Nd II | 0.03 |  $-$0.12 |  $-$0.13
Eu II | 0.05 |  $-$0.12 |  $-$0.14
Table 7: Parallaxes, Proper Motions and Distances

Star |  Gaia DR2 source ID | $\varpi$ | error |  | pmra | error |  | pmdec | error |  | Distance | d1 | d2 |
---|---|---|---|---|---|---|---|---|---|---|---|---|---|---
|  | (mas) |  | (mas yr-1) |  | (mas yr-1) |  | (kpc) |
HD2796 | 2367454697327877504 | 1.4859 | 0.0626 |  |  $-$1.375 | 0.164 |  |  $-$51.052 | 0.084 |  | 0.661 | 0.027 | 0.028 |
J1054$+$0528 | 3864140775805950208 | 0.2149 | 0.0393 |  |  $-$8.678 | 0.073 |  |  $-$ 4.183 | 0.055 |  | 3.548 | 0.404 | 0.502 |
J1529$+$0804 | 1164484488577137792 | 0.2365 | 0.0362 |  |  $-$2.401 | 0.056 |  |  $-$9.169 | 0.05 |  | 3.628 | 0.425 | 0.544 |
J1645$+$4357 | 1357725650023190784 | 0.0454 | 0.0207 |  | 0.327 | 0.033 |  |  $-$5.663 | 0.042 |  | 8.355 | 0.996 | 1.215 |
J2114$-$0616 | 6910940758263238912 | 0.4496 | 0.0373 |  |  $-$20.444 | 0.067 |  |  $-$6.376 | 0.068 |  | 2.082 | 0.153 | 0.178 |
J2216$+$2232 | 1878089211702170880 | 0.3721 | 0.0384 |  |  $-$4.454 | 0.056 |  |  $-$1.477 | 0.058 |  | 2.440 | 0.211 | 0.253 |

Note. — The d1 and d2 columns indicate the 16th percentile and 84th percentile
confidence intervals.

Table 8: Kinematics parameters  Star ID | VR | VΦ | VZ | V⟂ | Zmax | Rapo | Rperi | e | E | Lz
---|---|---|---|---|---|---|---|---|---|---
HD2796 | -95.03 | 89.62 | 37.81 | 102.28 | 1.02 | 9.59 | 2.28 | 0.62 | -1.45 | 763.40
J1054$+$0528 | 82.98 | 108.74 | -15.88 | 84.49 | 3.59 | 10.81 | 3.60 | 0.50 | -1.35 | 1036.00
J1529$+$0804 | -86.75 | 101.22 | -8.18 | 87.13 | 3.01 | 7.43 | 2.52 | 0.50 | -1.52 | 624.00
J1645$+$4357 | -74.40 | 210.51 | -58.18 | 94.45 | 7.67 | 14.50 | 7.82 | 0.30 | -1.15 | 1807.00
J2114$-$0616 | -220.45 | 314.95 | 37.41 | 223.60 | 4.70 | 33.16 | 5.64 | 0.71 | -0.85 | 2323.00
J2216$+$2232 | -35.03 | -56.42 | 182.30 | 185.64 | 8.70 | 9.22 | 4.08 | 0.41 | -1.35 | -478.40