#  Observational evidence points at AGB stars as possible progenitors of
CEMP-s & r/s stars 111 Based on data collected at Himalayan Chandra Telescope,
IAO, Hanle, India and at Subaru Telescope, which is operated by the National
Astronomical Observatory of Japan. The subaru data are retrieved from the JVO
portal (http://jvo.nao.ac.jp/portal/v2/) operated by the NAOJ

Meenakshi Purandardas Indian Institute of Astrophysics, Koramangala, Bangalore
560034, India;  Department of physics, Bangalore university, Jnana Bharathi
Campus, Karnataka 560056, India  Aruna Goswami Indian Institute of
Astrophysics, Koramangala, Bangalore 560034, India;

(Received December xx, 2020; Revised XXXX; Accepted xxx)

###### Abstract

Origin of enhanced abundance of heavy elements observed in the surface
chemical composition of carbon-enhanced metal-poor (CEMP) stars still remain
poorly understood. Here, we present detailed abundance analysis of seven CEMP
stars based on high resolution (R${\sim}$ 50 000) spectra that reveal enough
evidence of Asymptotic Giant Branch (AGB) stars being possible progenitors for
these objects. For the objects HE0110$-$0406, HE1425$-$2052 and HE1428$-$1950,
we present for the first time a detailed abundance analysis. Our sample is
found to consists of one metal-poor ([Fe/H]$<$$-1.0$) and six very metal-poor
([Fe/H]$<$$-2.0$) stars with enhanced carbon and neutron-capture elements. We
have critically analysed the observed abundance ratios of [O/Fe], [Sr/Ba] and
[hs/ls] and examined the possibility of AGB stars being possible progenitors.
The abundance of oxygen estimated in the programme stars are characteristics
of AGB progenitors except for HE1429$-$0551 and HE1447$+$0102\. The estimated
values of [Sr/Ba] and [hs/ls] ratios also support AGB stars as possible
progenitors. The locations of the programme stars in the absolute carbon
abundance A(C) vs. [Fe/H] diagram along with the Group I objects hint at
binary nature of the object. We have studied the chemical enrichment histories
of the programme stars based on abundance ratios [Mg/C], [Sc/Mn] and [C/Cr].
Using [C/N] and 12C/13C ratios we have examined if any internal mixing had
modified their surface chemical compositions. Kinematic analysis shows that
the objects HE 0110$-$0406 and HE1447$+$0102 are thick disk objects and the
remaining five objects belong to the halo population of the Galaxy.

stars: abundances—stars: chemically peculiar—stars: carbon

††journal: ApJ

##  1 Introduction

A substantial fraction of iron-deficient stars show enhancement of carbon
(Rossi et al., 1999). Studies of these group of stars are of special
significance as their surface compositions are believed to hold the fossil
records of the nucleosynthetic products of first stars. Such metal-poor
([Fe/H] $<$ $-1.0$) stars with enhanced carbon ([C/Fe] $>$ 0.7, Aoki et al.
(2007)) are known as carbon-enhanced metal-poor (CEMP) stars. Majority of the
CEMP stars also exhibit enhancement of neutron-capture elements (Norris et
al., 1997; Barbuy et al., 1997; Sneden et al., 2003). CEMP stars are
classified into different subgroups based on the abundance patterns of heavy
elements (Beers & Christlieb, 2005).

Initial classification of CEMP stars was based on the abundance patterns of
heavy elements Ba and Eu (Beers & Christlieb, 2005). Subsequently, different
authors (Jonsell et al., 2006; Masseron et al., 2010; Abate et al., 2016;
Frebel, 2018; Hansen et al., 2019; Goswami et al., 2021) have used different
criteria to classify the CEMP stars into various subclasses. A comprehensive
and comparative study of the classification schemes can be found in Goswami et
al. (2021). CEMP stars that are found to be enhanced in slow neutron-capture
process (s-process) elements are called CEMP-s stars. The majority of CEMP
stars belong to this group. The observed enhancement of carbon and s-process
elements in CEMP-s stars is thought to be originated due to mass transfer from
a binary companion that has once gone through the AGB phase and synthesised
the carbon and heavy elements. The radial velocity variations exhibited by
these objects provide strong evidence in favour of the binary mass transfer
scenario (McClure & Woodsworth, 1990; Preston & Sneden, 2001; Hansen et al.,
2016a; Jorissen et al., 2016). Hansen et al. (2016a) report that 82% of CEMP-s
stars in their sample exhibit radial velocity variations and are confirmed
binaries. These authors could not find any signatures of orbital motion for
four of the CEMP-s stars in their sample and these stars appeared to be single
or in a wide binary system. Choplin et al. (2017) showed that the abundance
patterns of three of these four CEMP-s stars match well with the yields of
fast rotating massive stars.

CEMP stars that show over abundances of rapid neutron-capture process
(r-process) elements are known as CEMP-r stars. CEMP-r stars are the rarest
sub-group of CEMP stars and are found to be unlikely in a binary system
(Hansen et al., 2011). Hansen et al. (2015) suggest that the enhancement of
r-process elements in CEMP-r stars is not attributed to any binary companion,
but it might be inherited from the interstellar medium from which the star was
formed.

CEMP stars that are enhanced in both s- and r- process elements are called
CEMP-r/s stars. Many studies (Hill et al., 2000; Cohen et al., 2003; Qian &
Wasserburg, 2003; Jonsell et al., 2006) have discussed various formation
scenarios to explain the peculiar abundance patterns observed in CEMP-r/s
stars. However none of these processes could re-produce the observed
frequencies of CEMP-r/s stars except the intermediate neutron-capture process
or i-process (Jonsell et al., 2006; Abate et al., 2016). The i-process occurs
at a neutron density which is intermediate to the neutron densities required
for s- and r-processes (Dardelet et al., 2015; Roederer et al., 2016; Hampel
et al., 2016, 2019). Hampel et al. (2016) could reproduce the observed
frequencies of CEMP-r/s stars based on their model calculations of i-process
nucleosynthesis. Many authors have showed that i-process can occur at
different evolutionary stages of an AGB star (Cowan & Rose, 1977; Campbell &
Lattanzio, 2008; Cristallo et al., 2009a; Campbell et al., 2010; Stancliffe et
al., 2011; Banerjee et al., 2018; Clarkson et al., 2018) as well as in rapidly
accreting white dwarf (Denissenkov et al., 2017, 2019) . According to this
postulate, CEMP-r/s star is expected to be in a binary system and accreted the
materials produced by its companion through i-process which is supported by
the discovery of many of the CEMP-r/s stars in binary systems (Barbuy et al.,
2005; Hansen et al., 2016a).

Another group of CEMP stars that do not show any enhancement in heavy elements
are called CEMP-no stars. As shown by many studies, the peculiar abundance
patterns observed in CEMP-no stars is attributed to the interstellar medium
from which they are formed, is polluted by Spin stars (Meynet et al., 2010;
Chiappini, 2013) or Faint supernovae that underwent mixing and fall back
(Umeda & Nomoto, 2005; Nomoto et al., 2013; Tominaga et al., 2014) or metal-
free massive stars (Heger & Woosley, 2010). Although most of the CEMP-no stars
are found to be single objects (Norris et al., 2013), a few of them are found
to be confirmed binaries (Hansen et al., 2016b). Arentsen et al. (2019)
suggest that the CEMP-no stars may be in a binary system with an extremely
metal-poor companion that once passed through the AGB phase that had not
produced any significant amount of s-process elements.

In this paper, we present the results from a high-resolution spectroscopic
analysis of seven CEMP stars. The paper is organized as follows: section 2
presents a brief summary of the previous works on the programme stars. The
selection of the objects and the source of spectra are described in section 3.
In section 4, we present the photometric temperature estimates of the
programme stars. The methodology used for the estimation of atmospheric
parameters is presented in section 5. The results of the abundance analysis
are presented in section 6. Uncertainties in our abundance estimates are
discussed in section 7. The kinematic analysis is presented in section 8.
Interpretations of the results are discussed in section 9 and the conclusions
are drawn in section 10.

##  2 Previous studies on the program stars: a brief summary and novelty of
this work

Some aspects of the programme stars were addressed by different authors. Here
we present a summary of their main results and the novelty of this work.
HE 0110$-$0406, and HE 1425$-$2052
Goswami (2005) and Goswami et al. (2010a) classified these objects as
potential CH star candidates based on low resolution (R$\sim$ 1300)
spectroscopic analysis. Beers et al. (2007) presented the results of broadband
BVRCICCCD photometry for HE 1425$-$2052\. We present for the first time an
abundance analysis for the object HE 1425$-$2052 using high-resolution
spectra. For HE 0110$-$0406, we have performed for the first time a high-
resolution spectroscopic analysis and presented some results in our earlier
two papers, Purandardas et al. (2019a) and Purandardas & Goswami (2020). In
this work, we have examined the possible progenitor of this object based on
the estimated elemental abundance ratios.
HE 1428$-$1950
Beers et al. (2007) presented photometric analysis of this object. From low-
resolution spectroscopic analysis Goswami et al. (2010a) showed that this
object exhibits the spectral properties of a C-R star. The spectrum of this
object was found to match well with the spectrum of the well known C-R star RV
Sct, with the CH band around 4300 Å, slightly stronger than that in the RV Sct
star. The object also showed stronger Ca I line at 4226 Å, and a stronger
feature due to Na D1 and D2 than usually seen in the spectra of CH stars.
Placco et al. (2011) have reported the atmospheric parameters for this object
based on the analysis of low-resolution (R $\sim$ 2000) spectra. Kennedy et
al. (2011) have presented the atmospheric parameters, metallicity, carbon and
oxygen abundances of this star using the low- resolution optical (R$\sim$
2000) and medium-resolution (R $\sim$ 3000) near IR spectra. Polarimetric
studies of Goswami & Karinkuzhi (2013) reported significant polarization for
this object over different colours of BVRI with percentage of polarization in
the range 0.2 to 0.9 that indicates the presence of a circumstellar envelope
for this object. Munari et al. (2014) reported the effective temperature and
reddening of this star using the AAVSO Photometric All-Sky Survey (APASS)
data. Although some properties of this object have been addressed in many
studies, results based on high-resolution spectroscopic analysis are missing.
We present for the first time a detailed abundance analysis for this object
using high-resolution spectra.
HE 1429$-$0551 and HE 1447+0102
Low-resolution spectroscopic analysis of HE 1429$-$0551 and HE 1447+0102 by
Goswami (2005) and Goswami et al. (2010a) showed that these objects are
potential CH star candidate. Beers et al. (2007) reported the B, V and R
magnitudes for these objects. For both the stars, Aoki et al. (2007) presented
the abundance analysis results for C, N, Na, Mg, Ca, Ti, and Ba using the
high-resolution (R$\sim$ 50000) spectra obtained using the High Dispersion
Spectrograph (HDS) of the Subaru telescope. Polarimetric studies of Goswami et
al. (2010b) in the V-band, reported percentage of polarization ${\sim}$ 1.3
$\pm$ 0.06 for HE 1429$-$0551 that indicates the existence of a circumstellar
envelop for this object. Bisterzo et al. (2011) grouped these objects as a
CEMP-s II stars that fall in the category of CEMP-s star with [hs/ls] $\geq$
1.5. They used the value for [hs/ls] from Aoki et al. (2007) for their
classification. Based on their model calculations, they claimed that the
observed abundance patterns in HE 1429$-$0551 and HE 1447+0102 may be due to
mass transfer from an AGB companion with a mass in the range 1.4 - 2 M⊙ and
1.5 - 2 M⊙ respectively. Jorissen et al. (2016) studied the radial velocity
variations exhibited by HE 1429$-$0551 and showed that the object does not
exhibit any signs of orbital motion and it requires more observations to
confirm its binarity. Karinkuzhi et al. (2021) report the abundances of a few
light and heavy elements for HE 1429$-$0551 based on an analysis of high-
resolution (R $\sim$ 86 000) spectrum obtained with HERMES spectrograph
mounted on the Mercator telescope. In this work, we have reported elemental
abundances of two other elements Sc and V for HE 1429$-$0551 and elemental
abundances of O, Sc, V, Cr, Mn, Ni, Sr, La, Ce, Pr, Nd, Sm, and Eu for HE
1447+0102 that are not available in literature. We have attempted to find the
possible progenitors of these objects based on the observed abundance
patterns. We have also checked the possibility of any internal mixing based on
[C/N] and 12C/13C ratios which was not addressed in any of the previous works.
HE 1523$-$1155, and HE 1528$-$0409
Goswami (2005) classified these objects as potential CH star candidates based
on the low-resolution spectroscopic analysis. Beers et al. (2007) reported B,
V and R magnitudes for these object. Aoki et al. (2007) reported the
abundances of C, N, Na, Mg, Ca, Ti, and Ba in these objects based on high-
resolution spectroscopic analysis. Polarimetric studies by Goswami et al.
(2010b) showed a percentage of polarization ${\sim}$ 0.85 $\pm$ 0.08 in V-band
for these objects that indicates presence of a circumstellar envelope.
Bisterzo et al. (2011) classified these objects as CEMP-s II stars with the
observed abundance patterns attributed to an AGB companion with a mass in the
range 1.5 - 2.0 M⊙. Hansen et al. (2016a) showed that the object HE
1523$-$1155 is a binary star based on the radial velocity variations exhibited
by this object. They claimed that more radial velocity monitoring of the
object is required to determine the final orbital parameters with confidence.
In this work, we have presented the estimates and analyses of a few more
elements such as O, Sc, V, Cr, Ni, Zn, Sr, Y, La, Ce, Pr, Nd, Sm, and Eu in
these objects, that are not available in literature.

##  3 Object selection and source of spectra

Objects are selected from the lists of potential CH star candidates of Goswami
(2005) and Goswami et al. (2010a) for a follow-up detailed high resolution
spectroscopic analysis to confirm their classification as well as to
understand the characteristic properties of their surface chemical
composition. In Figure 1 (Lower panel) a few sample spectra at low-resolution
obtained using the Himalayan Faint Object Spectrograph Camera (HFOSC) and
grism 7 covering the wavelength regions from 3860 - 7000 Å are shown. The
spectral resolution ($\lambda/\delta\lambda$) is $\sim$ 1330.

For the object HE 0110$-$0406, the high resolution ($\lambda/\delta\lambda\sim
60,000$) spectrum was obtained using the high resolution fiber fed Hanle
Echelle SPectrograph (HESP) attached to the Himalayan Chandra Telescope (HCT).
The wavelength coverage of the spectra spans from 3530-9970 Å. The
spectrograph operates at a resolution of 60 000 with a slicer, and a
resolution of 30 000 without slicer. A CCD with 4096$\times$4096 pixels of 15
micron size is used to record the spectrum. The high-resolution spectra of the
remaining programme stars were retrieved from the JVO portal
http://jvo.nao.ac.jp/portal/v2/ operated by the National Astronomical
Observatory of Japan (NAOJ) where the reduced and wavelength calibrated
spectra are publicly available. These spectra were obtained using the High
Dispersion Spectrograph (HDS) of the 8.2 m Subaru Telescope with a resolution
of $\sim$ 50 000 and a wavelength coverage from 4100 to 6850Å, with a small
gap between 5440 and 5520Å due to the physical separation between the two EEV
CCDs with 2048$\times$4096 pixels with two by two on-chip binning. The signal
to noise ratio of the spectra of the programme stars range from $\sim$ 53 - 67
per resolution element measured at 6000 Å. The basic data of the programme
stars are given in Table 1. Figure 1 (Upper panel) shows a few examples of
sample spectra at high resolution.

Table 1: Basic data for the programme stars

Star |  RA$(2000)$ |  Dec.$(2000)$ | B | V | J | H | K | Exposure | S/N | Date of obs. | Source
---|---|---|---|---|---|---|---|---|---|---|---
|  |  |  |  |  |  |  | (seconds) | (at 6000Å) |  | of spectra
HE 0110-0406 | 01 12 37.16 |  $-03$ 50 29.15 | 13.77 | 12.48 | 10.52 | 9.98 | 9.86 | 2700 | 66 | 24-12-2017 | HCT/HESP
HE1425-2052 | 14 28 39.54 |  $-21$ 06 04.78 | 13.97 | 12.70 | 10.04 | 9.45 | 9.27 | 600 | 64 | 25-05-2003 | Subaru/HDS
HE1428-1950 | 14 30 59.39 |  $-20$ 03 41.91 | 13.28 | 12.04 | 9.98 | 9.47 | 9.32 | 600 | 58 | 25-05-2003 | Subaru/HDS
HE1429-0551 | 14 32 31.29 |  $-06$ 05 00.20 | 14.01 | 12.61 | 10.73 | 10.27 | 10.07 | 600 | 65 | 25-05-2003 | Subaru/HDS
HE1447+0102 | 14 50 15.04 | 00 50 15.09 | 15.63 | 14.61 | 13.21 | 12.76 | 12.68 | 600 | 57 | 26-05-2003 | Subaru/HDS
HE1523-1155 | 15 26 41.04 |  $-12$ 05 42.66 | 14.57 | 13.22 | 11.37 | 10.85 | 10.75 | 600 | 67 | 26-05-2003 | Subaru/HDS
HE1528-0409 | 15 30 54.30 |  $-04$ 19 40.36 | 16.00 | 14.76 | 12.94 | 12.45 | 12.36 | 600 | 53 | 26-05-2003 | Subaru/HDS

Figure 1:  Upper panel: Examples of high resolution spectra of the programme
stars in the wavelength region 5160 to 5190 Å. Lower panel: Examples of low
resolution spectra of the programme stars in the wavelength region 3860 to
7000 Å

.

##  4 Photometric temperatures

We have determined the photometric temperatures for our programme stars using
the calibration equations of Alonso et al. (1999) following the procedures as
described in Goswami et al. (2006, 2016). Interstellar extinctions for objects
with b $<$ 50 are determined using the formulae from Chen et al. (1998). The
estimated values of reddening for our programme stars range from 0 - 0.04 and
are negligible. J, H, and K magnitudes of the objects are taken from SIMBAD
data base which came from 2MASS (Cutri et al., 2003). Photometric temperatures
are used as the initial guess for the determination of spectroscopic
temperatures of the programme stars. The photometric temperatures obtained for
our programme stars are presented in Table 2.

Table 2: Temperatures from photometry

Star name | $T_{e}ff$ | $T_{e}ff$ | $T_{e}ff$ | $T_{e}ff$ | $T_{e}ff$ | $T_{e}ff$ | $T_{e}ff$ | $T_{e}ff$ | $T_{e}ff$ | $T_{e}ff$ | $T_{e}ff$ | $T_{e}ff$ | Spectroscopic
---|---|---|---|---|---|---|---|---|---|---|---|---|---
|  |  ($-1.5$) |  ($-2.0)$ |  ($-2.5$) | [Fe/H] |  ($-1.5$) |  ($-2.0$) |  ($-2.5$) | [Fe/H] |  ($-1.5$) |  ($-2.0$) | [Fe/H] | estimates
|  (J$-$K) |  (J$-$H) |  (J$-$H) |  (J$-$H) | (J-H) |  (V$-$K) |  (V$-$K) |  (V$-$K) |  (V$-$K) |  (B$-$V) |  (B$-$V) |  (B$-$V) |
HE 0110$-$0406 | 4657.90 | 4756.88 | 4733.47 | 4688.11 | 4759.92 | 4447.21 | 4446.24 | 4451.10 | 4449.23 | 4162.19 | 4167.47 | 4166.76 | 4670
HE 1425$-$2052 | 4358.16 | 4526.41 | 4503.99 | 4486.04 | 4497.10 | - | - | - | - | 4189.81 | 4193.06 | 4196.62 | 4300
HE 1428$-$1950 | 4620.68 | 4820.74 | 4797.06 | 4750.81 | 4788.78 | 4365.58 | 4364.03 | 4368.09 | 4364.44 | 4231.34 | 4231.51 | 4234.86 | 4500
HE 1429$-$0551 | 4626.35 | 5064.78 | 5040.04 | 4990.39 | 4994.05 | 4507.99 | 4507.46 | 4512.92 | 4512.42 | - | - | - | 4940
HE 1447+0102 | 5084.56 | 5133.68 | 5108.64 | 5057.99 | 5082.30 | 5136.28 | 5144.59 | 5160.14 | 5152.73 | 4519.41 | 4496.74 | 4496.38 | 5220
HE 1523$-$1155 | 4755.96 | 4792.56 | 4769.00 | 4723.15 | 4730.87 | 4561.47 | 4561.32 | 4567.31 | 4566.09 | 4099.98 | 4109.71 | 4136.83 | 4780
HE 1528$-$0409 | 4872.38 | 4941.53 | 4917.33 | 4869.41 | 4884.08 | 4624.51 | 4621.44 | 4624.16 | 4622.89 | 4222.72 | 4223.54 | 4236.76 | 5250

Note. The numbers in the parenthesis below $T_{eff}$ indicate the metallicity
values at which the temperatures are calculated. The temperatures calculated
using the adopted metallicity of the stars are presented in columns 6, 10 and
13. Temperatures are given in Kelvin.

##  5 Spectroscopic stellar parameters

Radial velocities of the programme stars are measured using the shift in the
wavelength for a large number of clean and unblended lines in their spectra.
The estimated radial velocities range from $-151.5$ to 121.2 Km s-1 (Table 3).
Except for HE 1447+0102 and HE 1528$-$0409 for which the radial velocity
information are not available in literature, our estimates of radial
velocities of the remaining programme stars show a difference of $\sim$ 1.6 –
11.8 km s-1 from the corresponding literature values which may be a clear
indication of the stars being in binary systems.

Table 3: Radial velocities of the programme stars.

Star |  Vr | Reference
---|---|---
|  (Kms-1) |
HE 0110-0406 |  $-44.4$$\pm$3.80 | 1
|  $-32.8$$\pm$2.30 | 2
HE 1425$-$2052 |  121.2$\pm$0.22 | 1
|  133$\pm$1.59 | 2
HE 1428$-$1950 |  58.8$\pm$0.44 | 1
|  62.1$\pm$0.82 | 2
|  61.9$\pm$3.58 | 3
HE 1429$-$0551 |  $-48.6$$\pm$0.11 | 1
|  $-43.1$$\pm$0.88 | 2
|  $-44.9\pm$0.11 | 4
HE 1447+0102 |  $-61.8$$\pm$0.10 | 1
|  $-61.2$$\pm$0.21 | 4
HE 1523$-$1155 |  $-40.5$$\pm$0.08 | 1
|  $-38.9$$\pm$3.50 | 2
|  $-45.0$$\pm$0.12 | 4
HE 1528$-$0409 |  $-151.5$$\pm$0.13 | 1
|  $-157.5$$\pm$0.25 | 4

1\. Our work, 2. Gaia Collaboration et al. (2018), 3. Kunder et al. (2017), 4.
Aoki et al. (2007)

Atmospheric parameters the effective temperature Teff, surface gravity, log g
and microturbulent velocity ${\zeta}$ of the programme stars are determined
from the measured equivalent widths of a set of clean lines due to Fe I and Fe
II (Table 4). Effective temperature is taken as the value obtained when the
trend between the abundances of Fe I and Fe II and the corresponding
excitation potentials gives a zero slope. At this temperature, microturbulent
velocity is fixed to be that value for which there is no trend between the
abundances derived from the Fe I and Fe II lines and the reduced equivalent
widths. Under this temperature and microturbulent velocity, the surface
gravity, log g is taken to be that value corresponding to which the abundances
of Fe I and Fe II are nearly the same. The abundances of Fe I and Fe II for
the programme stars as a function of excitation potential and as a function of
equivalent widths are shown in Figure 2.

Figure 2:  The iron abundances of programme stars as a function of excitation
potential (upper panel) and equivalent widths (lower panel). The open circles
indicate Fe I lines and solid triangles represent Fe II lines. Table 4:
Equivalent widths (in mÅ) of Fe lines used for deriving atmospheric
parameters.

Wavelength(Å) | Element |  $E_{low}$(eV) | log gf | HE 0110-0406 | HE 1425-2052 | HE 1428-1950 | HE 1429-0551 | HE 1447+0102 | HE 1523-1155 | HE 1528-0409
---|---|---|---|---|---|---|---|---|---|---
\- 4139.927 | Fe I | 0.990 | -3.629 | - | - | - | - | - | 41.2(5.21) | 20.0(5.42)
4337.046 |  | 1.557 | -1.695 | - | - | - | 78.1(4.86) | - | - | -
4422.567 |  | 2.845 | -1.110 | - | 96.8(5.10) | - | - | - | - | -
4445.471 |  | 0.087 | -5.441 | - | 60.0(5.30) | 65.3(5.47) | - | - | - | -
4446.833 |  | 3.686 | -1.330 | - | - | - | - | - | 78.7(5.08) | -
4466.551 |  | 2.832 | -0.590 | - | - | 135.7(5.15) | 70.7(5.07) | 50.2(4.97) | - | -
4484.219 |  | 3.603 | -0.720 | - | 87.9(5.56) | 55.5(5.19) | - | - | 24.9(5.03) | -
4531.147 |  | 1.485 | -2.155 | - | - | - | 62.2(4.87) | - | - | -
4566.514 |  | 3.301 | -2.250 | 24.6(6.07) | - | - | - | - | - | -
4595.358 |  | 3.302 | -1.720 | 60.9(6.16) | - | 34.2(5.50) | - | - | - | -
4630.121 |  | 2.279 | -2.600 | - | - | 60.6(5.44) | - | - | - | -
4632.911 |  | 1.608 | -2.913 | - | - | 77.7(5.08) | - | - | 38.5(5.12) | -
4643.463 |  | 3.654 | -1.290 | - | - | 27.0(5.37) | - | - | - | -
4690.136 |  | 3.686 | -1.680 | 31.7(6.09) | - | - | - | - | - | -
4745.800 |  | 3.654 | -0.790 | - | - | 77.3(5.55) | - | - | - | -
4789.655 |  | 3.547 | -0.840 | - | - | - | 25.4(5.21) | - | 20.4(4.95) | -
4871.317 |  | 2.868 | -0.410 | - | 166.7(5.36) | 160.8(5.27) | 89.0(5.23) | 69.2(5.18) | - | -
4882.144 |  | 3.417 | -1.640 | 62.6(6.21) | - | - | - | - | - | -
4903.308 |  | 2.882 | -1.080 | - | 104.4(5.14) | 131.1(5.51) | 43.2(5.04) | 39.2(5.25) | 46.8(4.93) | -
4924.770 |  | 2.279 | -2.220 | - | - | 120.2(5.69) | - | - | 43.7(5.29) | 26.4(5.23)
4939.686 |  | 0.859 | -3.340 | - | - | 164.4(5.50) | - | - | - | -
4950.104 |  | 3.417 | -1.670 | 58.1(6.17) | - | - | - | - | - | -
4967.890 |  | 4.191 | -0.622 | 65.0(6.15) | - | - | - | - | - | -
4985.251 |  | 3.928 | -0.485 | - | 83.6(5.65) | - | - | - | - | -
4994.129 |  | 0.915 | -3.080 | - | - | - | - | 43.5(5.15) | - | -
5001.862 |  | 3.881 | 0.010 | - | 104.6(5.35) | - | 37.2(4.96) | 26.6(4.96) | 59.7(5.25) | -
5006.117 |  | 2.832 | -0.767 | - | - | - | - | 57.3(5.23) | - | 44.4(5.02)
5028.126 |  | 3.573 | -1.474 | - | - | 23.8(5.35) | - | - | - | -
5051.634 |  | 0.915 | -2.795 | - | - | - | 75.2(4.95) | - | - | 51.6(5.07)
5079.224 |  | 2.198 | -2.067 | - | - | - | 38.8(5.13) | 29.9(5.30) | - | -
5109.650 |  | 4.301 | -0.980 | 33.7(6.12) | - | - | - | - | - | -
5166.281 |  | 0.000 | -4.195 | - | - | - | 68.3(5.07) | 47.3(5.26) | 82.8(5.03) | 29.3(4.99)
5171.595 |  | 1.485 | -1.793 | - | - | - | 96.7(5.03) | - | 100.3(4.89) | -
5192.343 |  | 2.998 | -0.521 | - | - | 156.1(5.37) | 64.9(4.96) | - | - | 55.1(5.15)
5194.941 |  | 1.557 | -2.090 | - | - | 176.1(5.29) | 88.1(5.24) | - | - | 57.4(5.17)
5195.468 |  | 4.221 | -0.002 | - | - | 83.3(5.49) | - | - | 28.1(5.04) | -
5198.711 |  | 2.222 | -2.135 | - | 130.5(5.60) | - | 39.1(5.23) | 21.1(5.17) | 41.5(5.07) | 27.5(5.09)
5215.179 |  | 3.266 | -0.933 | 106.5(5.99) | 76.6(5.14) | - | 32.9(5.11) | 32.0(5.35) | 36.9(5.03) | -
5217.389 |  | 3.211 | -1.097 | 113.8(6.22) | 109.7(5.61) | - | - | - | 41.9(5.23) | -
5226.862 |  | 3.038 | -0.667 | - | - | - | 42.9(4.76) | 49.0(5.17) | - | 41.9(5.07)
5232.939 |  | 2.940 | -0.190 | - | 161.7(5.05) | 178.4(5.28) | 86.6(4.97) | - | 92.8(4.94) | 74.4(5.19)
5242.491 |  | 3.634 | -0.840 | - | - | 82.4(5.58) | - | - | 19.2(4.98) | -
5247.049 |  | 0.087 | -4.946 | - | 153.5(5.70) | 112.2(5.30) | 26.6(5.23) | - | 32.6(5.06) | -
5250.645 |  | 2.198 | -2.050 | - | - | - | 27.0(4.87) | - | - | 18.4(5.00)
5253.461 |  | 3.283 | -1.670 | - | 19.9(5.06) | - | - | - | - | -
5263.305 |  | 3.266 | -0.970 | 121.2(6.29) | - | - | - | - | - | -
5266.555 |  | 2.998 | -0.490 | - | 156.5(5.35) | 160.9(5.39) | 74.0(5.08) | 61.2(5.18) | 78.4(5.02) | -
5281.790 |  | 3.038 | -1.020 | - | 122.7(5.46) | - | 50.8(5.25) | 41.2(5.37) | - | 30.6(5.19)
5283.621 |  | 3.241 | -0.630 | - | - | - | - | - | 59.7(5.10) | 45.6(5.33)
5288.528 |  | 3.695 | -1.670 | 49.2(6.33) | - | - | - | - | - | -
5307.360 |  | 1.608 | -2.987 | - | 95.0(5.18) | - | 27.2(5.12) | - | 36.4(5.08) | -
5324.178 |  | 3.211 | -0.240 | - | - | 175.6(5.63) | 63.4(4.88) | 67.0(5.28) | 83.0(5.10) | 50.5(5.00)
5339.928 |  | 3.626 | -0.680 | - | - | - | 30.6(5.21) | - | 33.3(5.13) | -
5373.698 |  | 4.473 | -0.860 | 45.4(6.39) | - | - | - | - | - | -
5383.369 |  | 4.313 | 0.500 | 123.6(6.15) | - | - | - | - | - | -
5389.479 |  | 4.415 | -0.410 | 67.9(6.21) | - | - | - | - | - | -
5424.069 |  | 4.320 | 0.520 | 123.0(6.12) | - | - | - | - | - | -
5506.778 |  | 0.990 | -2.797 | - | - | - | - | 51.9(5.04) | - | -
5554.882 |  | 4.548 | -0.440 | - | 24.6(5.52) | - | - | - | - | -
5569.618 |  | 3.417 | -0.540 | - | - | - | - | - | - | 30.0(5.08)
5576.090 |  | 3.430 | -1.000 | - | 97.9(5.63) | - | - | - | - | -
5586.756 |  | 3.368 | -0.210 | - | 142.1(5.32) | 130.4(5.12) | 64.8(5.04) | - | 71.6(5.01) | -
5701.545 |  | 2.559 | -2.216 | 100.6(6.19) | - | - | - | - | - | -
6024.049 |  | 4.548 | -0.120 | 78.7(6.20) | 54.9(5.66) | - | - | - | - | -
6136.615 |  | 2.453 | -1.400 | - | 161.0(5.40) | 166.2(5.43) | 69.2(5.17) | 58.2(5.35) | 66.5(4.94) | 34.0(4.94)
6136.993 |  | 2.198 | -2.950 | 81.2(6.12) | - | - | - | - | - | -
6137.694 |  | 2.588 | -1.403 | 145.9(6.06) | 161.9(5.60) | 147.5(5.40) | - | 53.2(5.41) | 67.9(5.14) | 41.0(5.23)
6180.203 |  | 2.727 | -2.780 | 57.4(6.30) | - | - | - | - | - | -

Wavelength(Å) | Element |  $E_{low}$(eV) | log gf | HE 0110-0406 | HE 1425-2052 | HE 1428-1950 | HE 1429-0551 | HE 1447+0102 | HE 1523-1155 | HE 1528-0409
---|---|---|---|---|---|---|---|---|---|---
6219.279 |  | 2.198 | -2.433 | 126.1(6.24) | - | - | - | - | - | -
6230.726 |  | 2.559 | -1.281 | 167.3(6.24) | 150.6(5.29) | - | 52.8(4.90) | - | - | 37.9(5.01)
6246.317 |  | 3.603 | -0.960 | 108.1(6.29) | - | - | - | - | - | -
6252.554 |  | 2.404 | -1.687 | - | 154.8(5.53) | 163.0(5.59) | 40.3(4.92) | 39.0(5.23) | 54.3(4.96) | -
6301.498 |  | 3.654 | -0.745 | - | 60.4(5.20) | 67.2(5.24) | - | - | 30.1(5.11) | -
6335.328 |  | 2.198 | -2.230 | - | 109.9(5.28) | 123.3(5.41) | 21.1(4.83) | 22.5(5.20) | 39.7(5.01) | -
6411.647 |  | 3.654 | -0.820 | 99.1(6.07) | - | - | - | - | - | -
6421.349 |  | 2.278 | -2.027 | 146.6(6.22) | - | - | 39.8(5.09) | - | 55.5(5.15) | -
6430.844 |  | 2.176 | -2.006 | 155.1(6.19) | - | - | - | - | 73.9(5.29) | -
6494.980 |  | 2.404 | -1.273 | - | - | - | 59.4(4.79) | 61.0(5.19) | - | -
6575.019 |  | 2.588 | -2.820 | - | 39.9(5.61) | - | - | - | - | -
6593.871 |  | 2.433 | -2.422 | 112.3(6.27) | 93.3(5.60) | - | - | - | - | -
6677.989 |  | 2.692 | -1.470 | - | - | - | - | - | - | 33.0(5.22)
4491.405 | Fe II | 2.856 | -2.700 | - | - | 65.0(5.53) | - | - | - | -
4508.288 |  | 2.856 | -2.210 | - | 45.2(5.30) | 76.5(5.21) | 41.0(4.97) | - | - | -
4515.339 |  | 2.844 | -2.480 | - | 43.5(5.53) | - | - | - | 28.4(5.06) | 26.4(5.08)
4731.453 |  | 2.891 | -3.360 | 47.1(6.12) | - | - | - | - | - | -
4923.927 |  | 2.891 | -1.320 | - | - | - | - | 86.1(5.29) | - | 79.2(5.25)
5197.577 |  | 3.231 | -2.100 | - | - | 62.8(5.31) | 36.7(5.16) | 38.5(5.28) | - | -
5991.376 |  | 3.153 | -3.557 | 31.4(6.29) | - | - | - | - | - | -

The numbers in the parenthesis in columns 5-11 give the derived abundances
from the respective line. log gf values are taken from kurucz atomic line list
(https://www.cfa.harvard.edu).

Only those lines with excitation potential from 0 - 5 eV and equivalent widths
from 20 - 180 mÅ are considered for the analysis. We have made use of MOOG
(Sneden 1970, updated version 2013) for the analysis under the assumption of
local thermodynamic equilibrium. We have used alpha normal model atmospheres
([$\alpha$/Fe] = 0) for our analysis and the model atmospheres are adopted
from the Kurucz grid of model atmospheres with no convective overshooting
(http://cfaku5.cfa.hardvard.edu/). Solar abundances are taken from Asplund et
al. (2009). The derived atmospheric parameters of our programme stars along
with the literature values are presented in Table 5.

Table 5: Derived atmospheric parameters of our programme stars and literature
values.

Star |  Teff | log g | $\zeta$ | [Fe I/H] | [Fe II/H] | [Fe/H] | Reference
---|---|---|---|---|---|---|---
| (K) | cgs |  (km s-1) |  |  |  |
|  ($\pm$100) |  ($\pm$0.2) |  ($\pm$0.2) |  |  |  |
HE 0110$-$0406 | 4670 | 1.00 | 1.92 |  $-1.31$$\pm$0.09 (27) |  $-1.29$$\pm$0.12 (2) | $-1.30$ | 1
HE 1425$-$2052 | 4300 | 1.50 | 2.70 |  $-2.10$$\pm$0.20 (28) |  $-2.09$$\pm$0.16 (2) | $-2.09$ | 1
HE 1428$-$1950 | 4500 | 0.50 | 2.75 |  $-2.11$$\pm$0.16 (27) |  $-2.15$$\pm$0.16 (3) | $-2.11$ | 1
| 4562 | 3.75 | - | - | - | $-2.07$ | 2
| 4531 | 0.85 | - | - | - | $-2.07$ | 3
HE 1429$-$0551 | 4940 | 1.50 | 1.54 |  $-2.47$$\pm$0.15 (31) |  $-2.44$$\pm$0.13 (2) | $-2.45$ | 1
| 4700 | 1.50 | 2.00 |  $-2.47$$\pm$0.20 (17) |  $-2.48$$\pm$0.16 (3) | $-2.47$ | 4
| 4832 | 1.14 | 2.01 | - | - | $-2.70$ | 5
HE 1447+0102 | 5220 | 1.90 | 1.36 |  $-2.29$$\pm$0.12 (20) |  $-2.21$$\pm$0.00 (2) | $-2.25$ | 1
| 5100 | 1.70 | 1.80 |  $-2.47$$\pm$0.27 (12) |  $-2.45$$\pm$0.14 (3) | $-2.46$ | 4
HE 1523$-$1155 | 4780 | 1.60 | 1.52 |  $-2.43$$\pm$0.11 (30) |  $-2.44$ (1) | $-2.42$ | 1
| 4800 | 1.60 | 1.80 |  $-2.15$$\pm$0.19 (16) |  $-2.17$$\pm$0.14 (3) | $-2.16$ | 4
HE 1528$-$0409 | 5250 | 2.00 | 1.29 |  $-2.37$$\pm$0.13 (19) |  $-2.34$$\pm$0.12 (2) | $-2.35$ | 1
| 5000 | 1.80 | 1.80 |  $-2.61$$\pm$0.19 (15) |  $-2.59$$\pm$0.14 (1) | $-2.60$ | 4

1\. Our work, 2. Placco et al. (2011), 3. Kennedy et al. (2011), 4. Aoki et
al. (2007), 5. Karinkuzhi et al. (2021)
Note : The number of lines used for analysis is given inside the parenthesis.

##  6 Abundance analysis

Abundances of various elements are determined using the equivalent width
measurements as well as using the spectrum synthesis calculation of clean and
unblended lines. Lines are identified by over plotting the Arcturus spectrum
on the spectra of the programme stars. A master line list is prepared using
the measured equivalent widths of the lines and the other line information
such as excitation potential and log gf values that are taken from the Kurucz
database. Abundances of light elements, C and N, odd-Z element Na, $\alpha$
elements such as Mg, Ca, Sc and Ti and Fe-peak elements such as Cr, Mn, Co and
Ni are estimated whenever possible. Abundances of neutron-capture elements Sr,
Y, Ba, La, Ce, Pr, Nd, Sm and Eu are also determined whenever the lines are
available. For abundance determination of the elements Sc, V, Mn, Ba, La and
Eu, we have also used spectrum synthesis calculation considering their
hyperfine structures taken from various sources. Estimated abundances are
presented in Tables 6 \- 8.

###  6.1 Carbon, Nitrogen, and Oxygen

The abundance of oxygen is determined from the spectrum synthesis of the
oxygen forbidden line [OI] 6300.3 Å (Figure 3). We could determine O in all
the programme stars except HE 0110$-$0406\. In HE 0110$-$0406, [OI] 6300.3 Å
and O I 6363.7 Å lines are found to be blended and not suitable for abundance
determination. O is enhanced in all the stars with [O/Fe] ranging from 1.12 to
1.88 , except for HE 1523$-$1155 and HE 1528$-$0409\. The line list used for
the synthesis of O is taken from Kurucz database.

Figure 3:  Synthesis of [OI] line around 6300 Å. Dotted line represents
synthesized spectra and the solid line indicates the observed spectra. Short
dashed line represents the synthetic spectra corresponding to $\Delta$[O/Fe] =
-0.2 and long dashed line corresponds to $\Delta$[O/Fe] = +0.2

The abundance of carbon is determined using the spectrum synthesis calculation
of the C2 band at 5165 Å (Figure 4). We could also determine the C abundance
using the spectrum synthesis calculation of CH band around 4315 Å only for HE
1428$-$1950 and HE 1425$-$2052\. For the remaining stars CH band is found to
be saturated except HE 0110$-$0406, and HE 1528$-$0409 for which this region
is noisy and not suitable for abundance determination. Since we could
determine the abundance of C from the synthesis of C2 band at 5165 Å for all
the stars, we have used this value throughout our analysis. C is enhanced in
all the programme stars with [C/Fe] ranging from 0.73 to 2.29. We could
estimate the carbon isotopic ratio in all the programme stars. Carbon isotopic
ratio is determined from the spectrum synthesis calculation of the C2 swan
system around 4740Å (Figure 5). The estimated values are presented in Table 9.

Figure 4: Synthesis of C2 band around 5165 Å. Dotted line represents
synthesized spectra and the solid line indicates the observed spectra. Short
dashed line represents the synthetic spectra corresponding to $\Delta$ [C/Fe]
= -0.2 and long dashed line corresponds to $\Delta$[C/Fe] = +0.2 Figure 5:
Spectral synthesis fits of the C2 features around 4740 Å. Solid line indicates
the observed spectra. Short and long dashed lines are shown to illustrate the
sensitivity of the line strengths to the isotopic carbon abundance ratios

Nitrogen abundance is determined from the spectrum synthesis calculation of CN
band at 4215Å (Figure 6). N is enhanced in all the programme stars except HE
1425$-$2052 with [N/Fe] $\sim$ 0.39. The molecular line lists for C and N are
adopted from Sneden et al. (2014) and Ram et al. (2014).

Figure 6:  Synthesis of CN band around 4215 Å. Dotted line represents
synthesized spectra and the solid line indicates the observed spectra. Short
dashed line represents the synthetic spectra corresponding to $\Delta$[N/Fe] =
-0.2 and long dashed line corresponds to $\Delta$[N/Fe] = +0.2

###  6.2 Na, Mg, Ca, Sc, Ti, and V

Abundance of sodium is determined from the equivalent width measurements of
three Na I lines (Table 14). HE 1428$-$1950 shows enhancement of Na with
[Na/Fe] $\sim$ 1.59. In the remaining objects, Na is moderately enhanced with
[Na/Fe] ranging from 0.25 to 0.86, except HE 1528$-$0409 for which [Na/Fe]
$\sim$ 0.25. Abundance of magnesium is estimated using the equivalent width
measurement of three Mg I lines (Table 14). The estimated values of Mg
abundance lie in the range [Mg/Fe] $\sim$ 0.27 to 0.96.

We have used the equivalent width measurements of 11 Calcium I lines (Table
14) to estimate the abundance of Ca in the programme stars. Ca is near solar
in HE 1429$-$0551 and HE 0110$-$0406\. In the remaining objects, [Ca/Fe]
values range from 0.23 to 0.47. The abundance of Scandium is determined using
the spectrum synthesis calculation of the lines Sc II 4415.56, 4431.350 and
6245.63 using the hyperfine structures from Prochaska & McWilliam (2000). Sc
abundance is found to be near solar in all the programme stars except for HE
1425$-$2052 and HE 1523$-$1155 for which [Sc/Fe] values are 0.44 and 0.35
respectively.

The abundance of titanium is determined using the equivalent width
measurements of Ti I and Ti II lines (Table 14). Ti is found to be under
abundant in HE 1447+0102 with [Ti/Fe] $\sim$ $-0.28$. In the remaining
objects, [Ti/Fe] values are in the range 0.12 to 0.43. The abundance of
Vanadium is determined from the spectrum synthesis calculation of the lines V
I 4864.73 Å and 5727.05 Å using the hyperfine structures from Prochaska &
McWilliam (2000). HE 1428$-$1950 and HE 1528$-$0409 show enhancement of V with
[V/Fe] $\sim$ 0.90. While the abundance of V is moderately enhanced in HE
1425$-$2052 and HE 1447+0102, V is marginally enhanced in HE 1429$-$0551 and
HE 1523$-$1155\. V is found to be under abundant in HE 0110$-$0406.

###  6.3 Cr, Mn, Co, Ni, and Zn

Abundance of cromium is estimated using the equivalent width measurements of a
number of Cr I and Cr II lines (Table 14). While Cr is enhanced in HE
1428$-$1950 with [Cr/Fe] $\sim$ 1.29, it is found to be under abundant in HE
1425$-$2052 with [Cr/Fe] $\sim$ $-0.26$. HE 1523$-$1155 shows moderate
enhancement of Cr and near solar values for HE 1429$-$0551 and HE
0110$-$0406\. We could not measure Cr in HE 1447+0102 as the lines are found
to be blended and not suitable for abundance determination. The abundance of
Manganese is determined using the spectrum synthesis calculation of the lines
Mn I 4470.14 Å and 6013.51 Å considering the hyperfine structure from
Prochaska & McWilliam (2000). [Mn/Fe] values in the programme stars range from
$-0.25$ to 0.46.

The abundance of cobalt could be determined only in HE 0110$-$0406, HE
1425$-$2052 and HE 1428$-$1950\. The Co I lines are found to be blended in the
remaining objects. Co is enhanced in HE 1428$-$1950 with [Co/Fe] $\sim$ 0.93
and near solar in HE 0110$-$0406 and HE 1425$-$2052\. The abundance of Nickel
could be determined in all the programme stars except for HE 1528$-$0409\. In
all the stars, Ni is found to be slightly enhanced with [Ni/Fe] in the range
0.26 to 0.45. The abundance of Zinc could be determined only in HE
0110$-$0406, HE 1425$-$2052, HE 1428$-$1950 and HE 1523$-$1155 with [Zn/Fe] in
the range $-0.41$ to 0.50. In the remaining stars, the Zn I lines are found to
be highly blended and could not be used for abundance calculation.

###  6.4 Sr, Y, and Zr

The abundance of strontium is determined from the spectrum synthesis
calculation of the line Sr I 4607.33 Å. Sr is enhanced in all the programme
stars. We could not estimate the abundance of Sr in HE 1429$-$0551 and HE
1528$-$0409 as this region is very noisy and not suitable for abundance
calculation.

We have used the equivalent width measurements of three Yttrium II lines
(Table 14) for the determination of Y abundance in the programme stars. While
Y is enhanced in HE 0110$-$0406, HE 1425$-$2052 and HE 1428$-$1950 with [Y/Fe]
$\sim$ 1.28, 1.37, and 1.25 respectively, it is only moderately enhanced in HE
1523$-$1155 with [Y/Fe] $\sim$ 0.68. In the remaining stars, the abundance of
Y could not be determined as the lines are blended. We could not determine the
abundance of zirconium in any of our programme stars as the lines due to Zr
are either blended or the region is found to be noisy and not suitable for
abundance determination.

###  6.5 Ba, La, Ce, Pr, Nd, Sm, and Eu

Abundance of barium is determined from the spectrum synthesis calculations of
the lines Ba II 5853.66 and 6141.71 Å using the hyperfine structures of Ba
from McWilliam (1998). We could not estimate the Ba abundance in HE
1428$-$1950 as the lines Ba II 5853.66 and 6141.71Å are saturated. We have
also checked other lines due to Ba such as 4554.03 and 4934.08 Å in this
object; these lines are found to be highly blended and not suitable for
abundance calculation. Ba is enhanced in all the remaining programme stars
with [Ba/Fe] ranging from 1.61 to 3.21. In HE 1425$-$2052, Ba is found to be
highly enhanced with [Ba/Fe] $\sim$ 3.21. We have used the spectrum synthesis
calculation of the line Lanthanum II 4921.78 using the hyperfine structure
from Jonsell et al. (2006). La is enhanced in all the programme stars with
[La/Fe] ranging from $\sim$1.04 to 2.49. The abundance of Cerium is estimated
from the equivalent width measurements of ten Ce II lines (Table 14). All the
programme stars exhibit enhancement of Ce with [Ce/Fe] ranging from 1.14 to
2.04.

We could determine the abundance of praseodymium only in HE 0110$-$0406, HE
1429$-$0551, HE 1447+0102 and HE 1528$-$0409\. In all these stars Pr is found
to be enhanced with [Pr/Fe] ranging from 1.09 to 2.05. The abundance of
neodymium could be measured in all the programme stars and all of them exhibit
enhancement of Nd with [Nd/Fe] in the range 1.03 to 2.34. We could estimate
the abundance of samarium in all the programme stars with [Sm/Fe] in the range
1.15 to 2.66. We have used the spectrum synthesis calculation of the lines
Europium II 4129.71, 6437.63 and 6645.11 Å to determine the abundance of Eu in
the programme stars. The hyperfine structures of Eu are taken from Worley et
al. (2013). Eu is enhanced in all the programme stars except HE 0110$-$0406
and HE 1528$-$0409.

We have checked for the lines due to the heavy elements such as Ru I 4869.15
Å, several Nb lines Nb I 4152.50, 5642.09, 5729.18, and 5983.2 Å, Mo I lines
at 4904.39, and 5196.04 Å, and Dy I 4923.16 Å. In our programme stars, all
these lines are found to either absent or very weak and blended with
equivalent widths less than 10mÅ. We have also checked the line Pb I 4057.7Å
in the spectrum of HE 0110-0406 and this region is found to be noisy.

We have estimated the values of [ls/Fe], [hs/Fe], [hs/ls] where ls represents
light s-process elements (Sr, Y, Zr) and hs represents the heavy s-process
elements (Ba, La, Ce, Nd). The estimated values of [hs/ls] and the carbon
isotopic ratios are presented in Table 9. Table 10 presents a comparison of
the abundances of our programme stars with the literature values.

Table 6: Elemental abundances in HE 0110$-$0406

| Z |  solar log$\epsilon^{\ast}$ |  log$\epsilon$ | [X/H] | [X/Fe]
---|---|---|---|---|---
C (C2, 5165 Å) | 6 | 8.43 |  7.85$\pm$0.20(syn) | $-0.58$ | 0.73
N (CN, 4215 Å) | 7 | 7.83 |  7.15$\pm$0.20(syn) | $-0.68$ | 0.63
Na I | 11 | 6.24 | 5.53(1) | $-0.71$ | 0.60
Mg I | 12 | 7.60 | 6.56(1) | $-1.04$ | 0.27
Ca I | 20 | 6.34 |  5.02$\pm$0.13(4) | $-1.32$ | $-0.01$
Sc II | 21 | 3.15 |  1.70$\pm$0.20(1,syn) | $-1.45$ | $-0.14$
Ti I | 22 | 4.95 |  3.76$\pm$0.14(4) | $-1.19$ | 0.12
V I | 23 | 3.93 |  2.17$\pm$0.20(1,syn) | $-1.76$ | $-0.45$
Cr I | 24 | 5.64 |  4.29$\pm$0.20(3) | $-1.35$ | $-0.04$
Mn I | 25 | 5.43 |  4.30$\pm$0.20(1,syn) | $-1.13$ | 0.18
Fe I | 26 | 7.50 |  6.19$\pm$0.09(27) | $-1.31$ | -
Fe II | 26 | 7.50 |  6.21$\pm$0.12(2) | $-1.29$ | -
Co I | 27 | 4.99 |  3.73$\pm$0.04(3) | $-1.26$ | 0.05
Ni I | 28 | 6.22 | 5.36(1) | $-0.86$ | 0.45
Zn I | 30 | 4.56 | 2.84(1) | $-1.72$ | $-0.41$
Sr I | 38 | 2.87 |  2.34$\pm$0.20(1,syn) | $-0.53$ | 0.78
Y II | 39 | 2.21 | 2.20(1) | $-0.01$ | 1.28
Ba II | 56 | 2.18 |  2.50$\pm$0.20(1,syn) | 0.32 | 1.61
La II | 57 | 1.10 |  1.07$\pm$0.20(1,syn) | $-0.03$ | 1.26
Ce II | 58 | 1.58 |  1.49$\pm$0.21(3) | $-0.09$ | 1.20
Pr II | 59 | 0.72 |  0.82$\pm$0.09(2) | 0.10 | 1.39
Nd II | 60 | 1.42 |  1.49$\pm$ 0.18(4) | 0.07 | 1.36
Sm II | 62 | 0.96 |  0.90$\pm$0.14(4) | $-0.06$ | 1.23
Eu II | 63 | 0.52 |  $-0.25$$\pm$0.20(1,syn) | $-0.77$ | 0.52

∗ Asplund et al. (2009), The number inside the parenthesis shows the number of
lines used for the abundance determination.

Table 7: Elemental abundances in HE 1425$-$2052, HE 1428$-$1950 and HE
1429$-$0551

|  |  |  HE 1425$-$2052 |  HE 1428$-$1950 |  HE 1429$-$0551
---|---|---|---|---|---
| Z |  Solar log$\epsilon^{a}$ |  log$\epsilon$ | [X/H] | [X/Fe] |  log$\epsilon$ | [X/H] | [X/Fe] |  log$\epsilon$ | [X/H] | [X/Fe]
C (C2, 5165 Å) | 6 | 8.43 |  7.76$\pm$0.20(syn) | $-0.67$ | 1.43 |  7.92$\pm$0.20(syn) | $-0.51$ | 1.60 |  8.25$\pm$0.20(syn) | $-0.18$ | 2.29
C (CH, 4315 Å) | 6 | 8.43 |  7.72$\pm$0.20(syn) | $-0.71$ | 1.39 |  7.95$\pm$0.20(syn) | $-0.48$ | 1.63 | - | - | -
N (CN, 4215 Å) | 7 | 7.83 |  6.12$\pm$0.20(syn) |  $-$1.71 | 0.39 |  6.75$\pm$0.20(syn) | $-1.08$ | 1.03 |  6.44$\pm$0.20(syn) | $-1.39$ | 1.08
O | 8 | 8.69 |  7.71$\pm$0.20(syn) | $-0.98$ | 1.12 |  7.84$\pm$0.20(syn) | $-0.85$ | 1.26 |  8.10$\pm$0.20(syn) | $-0.59$ | 1.88
Na I | 11 | 6.24 | 4.75(1) |  $-$1.49 | 0.61 |  5.72$\pm$0.20(1,syn) | $-0.52$ | 1.59 | 4.63(1) | $-1.61$ | 0.86
Mg I | 12 | 7.60 | 5.83(1) |  $-$1.77 | 0.33 |  5.92$\pm$0.20(1,syn) | $-1.68$ | 0.43 | 5.55(1) | $-2.05$ | 0.42
Ca I | 20 | 6.34 |  4.59$\pm$0.20 (6) |  $-$1.75 | 0.35 |  4.70$\pm$0.20(2) | $-1.64$ | 0.47 |  3.92$\pm$0.10(3) | $-2.42$ | 0.05
Sc II | 21 | 3.15 |  1.50$\pm$0.20(1,syn) |  $-$1.65 | 0.44 |  1.09$\pm$0.20(1,syn) | $-2.06$ | 0.09 |  0.62$\pm$0.20(1,syn) | $-2.53$ | $-0.09$
Ti I | 22 | 4.95 |  3.20$\pm$0.13 (2) |  $-$1.75 | 0.35 | - | - | - | 2.75(1) | $-2.20$ | 0.27
Ti II | 22 | 4.95 |  3.15$\pm$0.20(2) |  $-$1.80 | 0.29 |  3.14$\pm$0.20(2) | $-1.81$ | 0.34 |  2.79$\pm$0.18(3) | $-2.16$ | 0.28
V I | 23 | 3.93 |  2.33$\pm$0.20(1,syn) |  $-$1.60 | 0.50 |  2.73$\pm$0.20(1,syn) | $-1.20$ | 0.91 |  1.81$\pm$0.20(1,syn) | $-2.12$ | 0.35
Cr I | 24 | 5.64 | 3.28(1) | $-2.36$ | $-0.26$ |  4.81$\pm$0.16(2) | $-0.83$ | 1.28 | 3.18(1) | $-2.46$ | 0.01
Cr II | 24 | 5.64 | - | - | - | 4.79(1) | $-0.85$ | 1.30 | - | - | -
Mn I | 25 | 5.43 |  3.08$\pm$0.20(1,syn) |  $-$2.35 | $-0.25$ |  3.66$\pm$0.20(1,syn) | $-1.77$ | 0.34 |  3.42$\pm$0.20(1,syn) | $-2.01$ | 0.46
Fe I | 26 | 7.50 |  5.40$\pm$0.20(28) |  $-$2.10 | - |  5.39$\pm$0.16(27) | $-2.11$ | - |  5.03$\pm$0.15(31) | $-2.47$ | -
Fe II | 26 | 7.50 |  5.41$\pm$0.16(2) |  $-$2.09 | - |  5.35$\pm$0.16(3) | $-2.15$ | - |  5.06$\pm$0.13(2) | $-2.44$ | -
Co I | 27 | 4.99 | 2.97(1) |  $-$2.02 | 0.08 | 3.81(1) | $-1.18$ | 0.93 | - | - | -
Ni I | 28 | 6.22 |  4.52$\pm$0.16(2) |  $-$1.70 | 0.40 |  4.38$\pm$0.05(2) | $-1.84$ | 0.27 | 4.01(1) | $-2.21$ | 0.26
Zn I | 30 | 4.56 | 2.88(1) |  $-$1.68 | 0.42 | 2.95(1) | $-1.61$ | 0.50 | - | - | -
Sr I | 38 | 2.87 |  2.31$\pm$0.20(1,syn) |  $-$0.56 | 1.53 |  1.80$\pm$0.20(1,syn) | $-1.07$ | 1.04 | - | - | -
Y II | 39 | 2.21 |  1.49$\pm$0.04(2) |  $-$0.72 | 1.37 | 1.30(1) | $-0.91$ | 1.25 | - | - | -
Ba II | 56 | 2.18 |  3.30$\pm$0.20(1,syn) | 1.12 | 3.21 | - | - | - |  1.62$\pm$0.20(1,syn) | $-0.56$ | 1.88
La II | 57 | 1.10 |  1.49$\pm$0.20(1,syn) | 0.39 | 2.49 |  $-0.01$$\pm$0.20(1,syn) | $-1.11$ | 1.04 |  $-0.20$$\pm$0.20(1,syn) | $-1.30$ | 1.14
Ce II | 58 | 1.58 | 1.70(1) | $-0.14$ | 1.95 |  0.57$\pm$0.17(3) | $-1.01$ | 1.14 |  1.04$\pm$0.20(4) | $-0.54$ | 1.90
Pr II | 59 | 0.72 | - | - | - | - | - | - |  $-0.63$(1) | $-1.35$ | 1.09
Nd II | 60 | 1.42 |  1.67$\pm$0.20(6) | 0.25 | 2.34 |  0.30$\pm$0.06(3) | $-1.12$ | 1.03 |  0.72$\pm$0.15(7) | $-0.70$ | 1.74
Sm II | 62 | 0.96 |  1.53$\pm$0.20(6) | 0.57 | 2.66 |  $-0.04$(1) | $1.00$ | 1.15 | 0.29(1) | $-0.67$ | 1.77
Eu II | 63 | 0.52 |  0.41$\pm$0.20(1,syn) | $-0.11$ | 1.98 |  $-0.56$$\pm$0.20(1,syn) | $-1.08$ | 1.07 |  $-0.67$$\pm$0.20(1,syn) | $-1.19$ | 1.25

a Asplund et al. (2009), The number inside the parenthesis shows the number of
lines used for the abundance determination.

Table 8: Elemental abundances in HE 1447+0102, HE 1523$-$1155 and HE
1528$-$0409

|  |  | HE 1447+0102 |  HE 1523$-$1155 |  HE 1528$-$0409
---|---|---|---|---|---
| Z |  Solar log$\epsilon^{a}$ |  log$\epsilon$ | [X/H] | [X/Fe] |  log$\epsilon$ | [X/H] | [X/Fe] |  log$\epsilon$ | [X/H] | [X/Fe]
C (C2, 5165 Å) | 6 | 8.43 |  8.19$\pm$0.20(syn) | $-0.24$ | 2.05 |  7.83$\pm$0.20(syn) | $-0.60$ | 1.83 |  8.13$\pm$0.20(syn) | $-0.30$ | 2.07
N (CN, 4215 Å) | 7 | 7.83 |  7.13$\pm$0.20(syn) | $-0.70$ | 1.59 |  6.52$\pm$0.20(syn) | $-1.31$ | 1.12 |  7.13$\pm$0.20(syn) | $-0.7$ | 1.67
O | 8 | 8.69 |  8.09$\pm$0.20(syn) | $-0.60$ | 1.69 |  7.08$\pm$0.20(syn) | $-1.61$ | 0.82 |  6.91$\pm$0.20(syn) | $-1.78$ | 0.59
Na I | 11 | 6.24 |  4.60$\pm$0.05(2) | $-1.64$ | 0.65 | 4.53(1) | $-1.71$ | 0.72 |  4.12$\pm$0.08(2) | $-2.12$ | 0.25
Mg I | 12 | 7.60 |  5.88$\pm$0.19(2) | $-1.72$ | 0.57 | 6.13(1) | $-1.47$ | 0.96 | 5.80(1) | $-1.80$ | 0.57
Ca I | 20 | 6.34 |  4.28$\pm$0.14(5) | $-2.06$ | 0.23 |  4.26$\pm$0.10(5) | $-2.08$ | 0.35 |  4.28$\pm$0.12(5) | $-2.06$ | 0.31
Sc II | 21 | 3.15 |  1.18$\pm$0.20(1,syn) | $-1.96$ | 0.24 |  1.05$\pm$0.20(1,syn) | $-2.10$ | 0.35 |  0.90$\pm$0.20(1,syn) | $-2.25$ | 0.09
Ti I | 22 | 4.95 | - | - | - | 2.85(1) | $-2.10$ | 0.33 | - | - | -
Ti II | 22 | 4.95 |  2.46$\pm$0.13(4) | $-2.49$ | $-0.28$ |  2.93$\pm$0.10(4) | $-2.02$ | 0.43 |  2.81$\pm$0.13(3) | $-2.14$ | 0.20
V I | 23 | 3.93 |  2.27$\pm$0.20(1,syn) | $-1.66$ | 0.63 |  1.81$\pm$0.20(1,syn) | $-2.12$ | 0.31 |  2.46$\pm$0.20(1,syn) | $-1.47$ | 0.90
Cr I | 24 | 5.64 | - | - | - |  3.73$\pm$0.16(2) | $-1.91$ | 0.52 | - | - | -
Cr II | 24 | 5.64 | 3.72 | $-1.92$ | 0.29 | 3.64(1) | $-2.00$ | 0.45 | - | - | -
Mn I | 25 | 5.43 |  3.60$\pm$0.20(1,syn) | $-1.83$ | 0.46 | - | - | - | - | - | -
Fe I | 26 | 7.50 |  5.21$\pm$0.12(20) | $-2.29$ | - |  5.07$\pm$0.11(30) | $-2.43$ | - |  5.13$\pm$0.13(19) | $-2.37$ | -
Fe II | 26 | 7.50 |  5.29$\pm$0.0(2) | $-2.21$ | - | 5.06(1) | $-2.44$ | - |  5.16$\pm$0.12(2) | $-2.34$ | -
Ni I | 28 | 6.22 | 4.37(1) | $-1.85$ | 0.44 |  4.20$\pm$0.12(3) | $-2.02$ | 0.41 | - | - | -
Zn I | 30 | 4.56 | - | - | - | 2.54(1) | $-2.02$ | 0.41 | - | - | -
Sr I | 38 | 2.87 |  2.15$\pm$0.20(1,syn) | $-0.72$ | 1.57 |  1.37$\pm$0.20(1,syn) | $-1.50$ | 0.93 | - | - | -
Y II | 39 | 2.21 | - | - | - | 0.44(1) | $-1.77$ | 0.68 | - | - | -
Ba II | 56 | 2.18 |  1.90$\pm$0.20(1,syn) | $-0.28$ | 1.93 |  1.57$\pm$0.20(1,syn) | $-0.61$ | 1.83 |  1.55$\pm$0.20(1,syn) | $-0.63$ | 1.71
La II | 57 | 1.10 |  1.01$\pm$0.20(1,syn) | $-0.09$ | 2.12 |  0.24$\pm$0.20$\pm$0.20(1,syn) | $-0.86$ | 1.59 |  0.09$\pm$0.20(1,syn) | $-1.01$ | 1.33
Ce II | 58 | 1.58 |  1.41$\pm$0.12(3) | $-0.17$ | 2.04 |  0.70$\pm$0.14(4) | $-0.88$ | 1.57 |  1.13$\pm$0.20(2) | $-0.45$ | 1.89
Pr II | 59 | 0.72 |  0.46$\pm$0.06(2) | $-0.26$ | 1.95 | - | - | - |  0.43$\pm$0.07(2) | $-0.29$ | 2.05
Nd II | 60 | 1.42 |  1.26$\pm$0.17(11) | $-0.16$ | 2.05 |  0.59$\pm$0.12(8) | $-0.83$ | 1.62 |  0.90$\pm$0.20(5) | $-0.52$ | 1.82
Sm II | 62 | 0.96 |  1.05$\pm$0.15(3) | 0.09 | 2.30 |  0.40$\pm$0.18(5) | $-0.56$ | 1.89 | 0.99(1) | 0.03 | 2.37
Eu II | 63 | 0.52 |  0.26$\pm$0.20(1,syn) | $-0.26$ | 1.95 |  $-0.59$$\pm$0.20 (1,syn) | $-1.11$ | 1.34 |  $-1.36$$\pm$0.20 (1,syn) | $-1.88$ | 0.46

a Asplund et al. (2009), The number inside the parenthesis shows the number of
lines used for the abundance determination.

Table 9: Estimates of [Fe/H], [ls/Fe], [hs/Fe], [hs/ls] and 12C/13C Star name | [Fe/H] | [ls/Fe] | [hs/Fe] | [hs/ls] |  12C/13C
---|---|---|---|---|---
HE 0110$-$0406 | $-1.30$ | 1.03 | 1.36 | 0.33 | 45
HE 1425$-$2052 |  $-$2.10 | 1.45 | 2.49 | 1.04 | 8.0
HE 1428$-$1950 | $-2.13$ | 1.14 | 1.07 | 0.07 | 23.0
HE 1429$-$0551 | $-2.45$ | - | 1.66 | - | 20.0
HE 1447+0101 | $-2.25$ | - | 2.03 | - | 18.0
HE 1523$-$1155 | $-2.44$ | 0.80 | 1.65 | 0.85 | 5.5
HE 1528$-$0409 | $-2.35$ | - | 1.69 | - | 7.0
Table 10: Comparison of the abundances of our programme stars with the
literature values.

Star name | [Fe I/H] | [Fe II/H] | [Fe/H] | [Sr/Fe] | [Y/Fe] | [Ba II/Fe] | Ref
---|---|---|---|---|---|---|---
HE 1429$-$0551 | $-2.47$ | $-2.44$ |  | - | - | 1.88 | 1
| - | - | $-2.70$ | - | 0.89 | 1.52 | 2
| $-2.47$ | $-2.48$ | $-2.47$ | - | - | 1.57 | 3
HE 1447+0102 | $-2.29$ | $-2.21$ | $-2.25$ | 1.57 | - | 1.93 | 1
| $-2.47$ | $-2.45$ | $-2.46$ | - | - | 2.70 | 3
HE 1523$-$1155 | $-2.43$ | $-2.44$ | $-2.43$ | 0.93 | 0.68 | 1.83 | 1
| $-2.15$ | $-2.17$ | $-2.16$ | - | - | 1.72 | 3
HE 1528$-$0409 | $-2.37$ | $-2.34$ | $-2.35$ | - | - | 1.71 | 1
| $-2.61$ | $-2.59$ | $-2.60$ | - | - | 2.30 | 3
Star name | [La II/Fe] | [Ce II/Fe] | [Pr II/Fe] | [NdII/Fe] | [Sm II/Fe] | [Eu II/Fe] | Ref
HE 1429$-$0551 | 1.14 | 1.90 | 1.09 | 1.74 | 1.77 | 1.25 | 1
| 1.30 | 1.42 | 1.38 | 1.43 | 1.27 | 1.08 | 2

1\. This work, 2. Karinkuzhi et al. (2021), 3. Aoki et al. (2007)

##  7 Abundance uncertainties

The total uncertainties in the derived elemental abundances are contributed by
random errors and systematic errors. While the uncertainties in the line
parameters such as equivalent widths, line blending, and oscillator strength
cause random errors, the uncertainties in the stellar atmospheric parameters
causes systematic errors. We have followed the procedures as described in
Shejeelammal et al. (2021) to calculate the uncertainties on the elemental
abundances derived from the equivalent width measurement as well as the
spectrum synthesis calculation. The total uncertainties in the estimated
elemental abundances, log$\epsilon$ :

$\sigma^{2}_{log\epsilon}=\sigma^{2}_{ran}+(\frac{\partial
log\epsilon}{\partial T})^{2}\sigma_{T_{eff}}^{2}+(\frac{\partial
log\epsilon}{\partial logg})^{2}\sigma^{2}_{logg}+(\frac{\partial
log\epsilon}{\partial\zeta})^{2}\sigma^{2}_{\zeta}+(\frac{\partial
log\epsilon}{\partial[Fe/H]})^{2}\sigma^{2}_{[Fe/H])}~{}~{}~{}~{}~{}~{}~{}~{}~{}~{}(1)$

where $\sigma^{2}_{ran}$ = $\sigma_{s}/\sqrt{N}$ . $\sigma_{s}$ represents the
standard deviation in the elemental abundances derived using N number of lines
due to that element. The $\sigma$’s indicate the typical uncertainties in the
adopted stellar atmospheric parameters, and are Teff $\sim$ $\pm$ 100 K, log g
$\sim$ $\pm$ 0.2 dex, $\zeta$ $\sim$ $\pm$ 0.2 km s-1 , and [Fe/H] $\sim$
$\pm$ 0.1 dex. The uncertainty in [X/Fe] is determined using the relation :

$\sigma^{2}_{[X/Fe]}=\sigma^{2}_{X}+\sigma^{2}_{[Fe/H]}~{}~{}~{}~{}~{}~{}~{}~{}~{}~{}(2)$

As an example, the differential elemental abundances derived using the
equation (2) for the object HE 1523-1155 are given in Table 11. The error for
the abundances derived from the spectrum synthesis calculation is taken to be
0.2 dex (indicated by “syn” in Tables 6, 7, 8). This is the minimum change in
the abundance value producing well distinguished synthetic spectra with
respect to the best fits. For abundances derived using the equivalent width
measurements, the errors listed represent the standard deviation (Tables 6, 7,
8).

Table 11: Differential elemental abundances (log$\epsilon$) derived for the object HE 1523$-$1155\.  element |  $\Delta$ Teff |  $\Delta$ log g | $\Delta\zeta$ |  $\delta$ [Fe/H] |  ($\sum$ $\sigma^{2}_{i}$)1/2 |  $\sigma$[X/Fe]
---|---|---|---|---|---|---
|  ($\pm$100K) |  ($\pm$0.2dex) |  ($\pm$0.2kms-1) |  ($\pm$0.1 dex) |  |
C |  $\pm$0.21 |  $\mp$0.03 | 0.00 | 0.00 | 0.21 | 0.31
N |  $\pm$0.31 |  $\mp$0.07 |  $\pm$0.02 |  $\pm$0.04 | 0.32 | 0.39
O |  $\pm$0.05 |  $\pm$0.06 | 0.00 |  $\pm$0.02 | 0.08 | 0.24
Na I |  $\pm$0.15 |  $\mp$0.08 |  $\mp$0.08 |  $\mp$0.02 | 0.19 | 0.22
Mg I |  $\pm$0.07 |  $\mp$0.05 |  $\mp$0.09 |  $\mp$0.01 | 0.12 | 0.17
Ca I |  $\pm$0.08 |  $\mp$0.02 |  $\mp$0.05 | 0.00 | 0.10 | 0.15
Sc II |  $\pm$0.06 |  $\pm$0.08 | 0.00 |  $\pm$0.02 | 0.10 | 0.25
Ti I |  $\pm$0.14 |  $\mp$0.03 |  $\mp$0.04 |  $\mp$0.01 | 0.15 | 0.19
Ti II |  $\pm$0.05 |  $\pm$0.05 |  $\mp$0.08 | 0.00 | 0.11 | 0.16
V I |  $\pm$0.12 |  $\mp$0.03 |  $\mp$0.05 |  $\mp$0.01 | 0.13 | 0.26
Cr I |  $\pm$0.13 |  $\mp$0.02 |  $\mp$0.08 |  $\mp$0.01 | 0.15 | 0.22
Cr II |  $\mp$0.01 |  $\pm$0.07 |  $\mp$0.01 |  $\pm$0.01 | 0.07 | 0.13
Fe I |  $\pm$0.11 |  $\pm$0.01 |  $\pm$0.04 |  $\mp$0.04 | 0.12 |
Fe II |  $\pm$0.04 |  $\pm$0.09 |  $\mp$0.03 |  $\pm$0.01 | 0.10 |
Ni I |  $\pm$0.11 |  $\mp$0.01 |  $\mp$0.03 |  $\mp$0.01 | 0.11 | 0.17
Zn I |  $\pm$0.04 |  $\pm$0.03 |  $\mp$0.02 |  $\pm$0.01 | 0.05 | 0.12
Sr I |  $\pm$0.30 |  $\pm$0.08 |  $\pm$0.08 |  $\pm$0.05 | 0.32 | 0.40
Y II |  $\pm$0.06 |  $\pm$0.06 |  $\mp$0.12 |  $\pm$0.01 | 0.15 | 0.18
Ba II |  $\pm$0.08 |  $\mp$0.04 |  $\mp$0.25 |  $\mp$0.09 | 0.28 | 0.36
La II |  $\pm$0.07 |  $\pm$0.05 |  $\mp$0.06 | 0.00 | 0.10 | 0.25
Ce II |  $\pm$0.07 |  $\pm$0.06 |  $\mp$0.06 |  $\pm$0.01 | 0.11 | 0.17
Nd II |  $\pm$0.08 |  $\pm$0.06 |  $\mp$0.05 |  $\pm$0.01 | 0.11 | 0.16
Sm II |  $\pm$0.07 |  $\pm$0.06 |  $\mp$0.04 |  $\pm$0.01 | 0.10 | 0.17
Eu II |  $\pm$0.04 |  $\pm$0.06 |  $\pm$0.02 |  $\pm$0.02 | 0.08 | 0.24

##  8 Kinematic analysis

Kinematic analysis of the programme stars are performed following a detailed
procedure as explained in Purandardas & Goswami (2021). We have used right-
handed coordinate system for the determination of space velocities of the
programme stars. That is, the three components U, V and W of the space
velocities, U is positive in the direction of Galactic center, V is positive
in the direction Galactic rotation and W is positive in the direction of the
North Galactic Pole (Johnson & Soderblom, 1987). The values of proper motions
and parallax required for the calculations of space velocities are taken from
SIMBAD and Gaia (Gaia Collaboration et al., 2016, 2018) respectively. We have
used our radial velocity estimates for the calculations.

We have determined the probability for a star’s membership into the thin disk,
the thick disk or the halo population. Our estimates show that the objects HE
0110$-$0406 and HE 1447+0102 belong to thick disk population with a
probability of $\sim$ 94% and 84% respectively. The remaining objects belong
to the Halo population with a probability $>$ 81%. The spatial velocities and
the probability estimates for the programme stars are presented in Table 12.

Table 12: Spatial velocity and probability estimates for the programme stars

Star name |  ULSR |  VLSR |  WLSR |  Vspa |  pthin |  pthick |  phalo | Population
---|---|---|---|---|---|---|---|---
|  (kms-1) |  (kms-1) |  (kms-1) |  (kms-1) |  |  |  |
HE 0110$-$0406 |  66.27$\pm$4.91 |  $-13.36$$\pm$1.88 |  29.28$\pm$3.96 |  73.67$\pm$5.64 | 0.94 | 0.06 | 0.00 | Thick
HE 1425$-$2052 |  289.78$\pm$100.54 |  $-156.13$$\pm$63.38 |  $-230.36$$\pm$161.09 |  401.77$\pm$45.33 | 0.00 | 0.00 | 1.00 | Halo
HE 1428$-$1950 |  135.01$\pm$17.25 |  $-131.24$$\pm$25.52 |  $-128.65$$\pm$35.87 |  228.04$\pm$24.33 | 0.00 | 0.18 | 0.81 | Halo
HE 1429$-$0551 |  7.65$\pm$5.19 |  $-501.65$$\pm$97.37 |  $-186.93$$\pm$29.35 |  535.40$\pm$101.39 | 0.00 | 0.00 | 1.00 | Halo
HE 1447+0102 |  $-26.39$$\pm$0.74 |  $-159.84$$\pm$16.69 |  $-53.82$$\pm$1.36 |  170.71$\pm$16.17 | 0.00 | 0.84 | 0.16 | Thick
HE 1523$-$1155 |  $-128.62$$\pm$14.75 |  $-307.01$$\pm$44.54 |  66.61$\pm$11.44 |  339.46$\pm$43.54 | 0.00 | 0.00 | 0.99 | Halo
HE 1528$-$0409 |  117.61$\pm$73.27 |  $-290.93$$\pm$100.23 |  $-353.23$$\pm$86.61 |  472.48$\pm$105.49 | 0.00 | 0.00 | 1.00 | Halo

##  9 Discussion

We have carried out a detailed abundance analysis for a sample of seven
objects, one object is found to be metal-poor and the remaining six are found
to be very metal-poor stars. Our analysis shows that the objects HE 1447+0102
and HE 1523$-$1155 satisfy the criteria for CEMP-r/s stars ([C/Fe] $>$ 0.7
(Aoki et al., 2007), [Ba/Fe] $\geq$ 1.0, [Eu/Fe] $\geq$ 1.0, 0.0 $\leq$
[Ba/Eu] $\leq$ 1.0 and/or 0.0 $\leq$ [La/Eu] $\leq$ 0.7 (Goswami et al.,
2021)) and all the remaining objects satisfy the criteria for CEMP-s stars
([C/Fe] $>$ 0.7 (Aoki et al., 2007), [Ba/Fe] $\geq$ 1.0, [Ba/Eu] $>$ 0.0
(Goswami et al., 2021)). We could not classify HE 1428-1950 based on criteria
involving Ba abundance as we could not determine the abundance of Ba in this
object.

All the programme stars exhibit enhancement of carbon and nitrogen except HE
1425$-$2052 for which nitrogen is only slightly enhanced. Oxygen is also found
to be enhanced in the stars except HE 1523$-$1155 and HE 1528$-0409$. The
$\alpha$-elements are found to be slightly enhanced in all the programme stars
except HE 0110$-$0406 for which the abundance of $\alpha$-elements are found
to be near solar. While HE 1425$-$2052 shows near solar abundance of Fe-peak
elements, HE 1428$-$1950 exhibits enhancement of Fe-peak elements except for
Ti and Sc. The remaining objects show slight enhancement in Fe-peak elements.

A comparison of the abundances of various elements estimated in our programme
stars with their counterparts observed in CEMP-s, CEMP-r/s and CEMP-no stars
taken from various sources is shown in the Figures 7 and 8. Although CEMP-s
and r/s stars fall at slightly higher metallicity range than CEMP-no stars,
all of them exhibit similar $\alpha$ and Fe peak elemental abundance patterns
except for Mn and V. A substantial number of CEMP-no stars show [Mn/Fe] and
[V/Fe] $<$ 0 unlike CEMP-s and CEMP-r/s stars. Majority of the CEMP-no stars
exhibit Na abundance in the range $-$0.50 $<$ [Na/Fe] $<$ 0.50, and CEMP-s and
r/s stars show [Na/Fe] $>$ 0\. Unlike CEMP-s and CEMP-r/s stars, the heavy
elements such as Sr and Ba are found to be underabundant in CEMP-no stars.

Figure 7: Abundance ratios of light elements with respect to metallicity
estimated in our programme stars are represented using red color symbols: HE
0110$-$0406 is represented by open triangle. Filled circle represents HE
1425$-$2052\. Filled triangle indicates HE 1428$-$1950\. HE 1429$-$0551 is
represented by filled square. Filled pentagon represents HE 1447+0102. Filled
hexagon and asterisk symbol indicate HE 1523$-$1155 and HE 1528$-$0409
respectively. CEMP-s and CEMP-r/s stars are represented by filled squares
(magenta) and filled triangles respectively. Filled circles(blue) represent
CEMP-no stars. The abundance values for CEMP-no stars used for the comparison
are taken from Christlieb et al. (2004), Plez & Cohen (2005), Yong et al.
(2013), Hansen et al. (2014), Bonifacio et al. (2015), Bessell et al. (2015)
and Frebel (2018). The abundance values for CEMP-s and CEMP-r/s stars are
taken from Lucatello et al. (2003), Barklem et al. (2005), Cohen et al.
(2006), Goswami et al. (2006), Aoki et al. (2007), Karinkuzhi & Goswami
(2015), Purandardas et al. (2019b), Purandardas et al. (2019a) and Goswami et
al. (2021). Representative error bar is shown at the left bottom corner of the
figure. Figure 8: Same as Figure 7, but for heavy elements. Representative
error bar is shown at the right top corner of the figure.

The location of a star in the absolute carbon abundance A(C) vs. metallicity
[Fe/H] diagram, and various elemental abundance ratios can give important
clues about the possible origin of the observed enhancement of carbon as well
as neutron-capture elements. A discussion along this line is presented in
detail in the following subsections, in the context of the programme stars.

###  9.1 Locations of the programme stars in the A(C) vs. [Fe/H] diagram:

Many authors have shown that the CEMP stars exhibit bimodal distribution in
the A(C) vs. [Fe/H] diagram (Spite et al., 2013; Bonifacio et al., 2015;
Hansen et al., 2015; Yoon et al., 2016). While Spite et al. (2013) claimed a
higher carbon band around A(C) $\sim$ 8.25 and a lower carbon band around A(C)
$\sim$ 6.5, Yoon et al. (2016) suggested a higher carbon band which peaks
around A(C) $\sim$ 7.96 and a lower band peaking at A(C) $\sim$ 6.28. Yoon et
al. (2016) further classified the objects in the A(C) vs. [Fe/H] diagram into
three groups. In this diagram, the objects that show very weak dependence of
A(C) on [Fe/H] are called as Group I objects. This group is distributed around
the higher carbon band. They are mostly composed of CEMP-s and CEMP-r/s stars
and a large fraction of them are confirmed binaries. Group II and Group III
objects are extremely metal-poor objects, mostly composed of CEMP-no stars,
that are clustered around the lower carbon band. While for Group III objects,
A(C) shows no clear dependence on [Fe/H], the A(C) values for Group II objects
exhibit very clear dependence on [Fe/H]. A large fraction of Group II and
Group III objects are found to be single and hence the observed abundance
anomalies are believed to be intrinsic in origin.

The locations of our programme stars in the A(C) vs. [Fe/H] diagram is shown
in Figure 9. In order to locate the programme stars in the A(C) vs. [Fe/H]
diagram of Yoon et al. (2016), corrections have to be applied to the estimated
abundance of carbon. The necessary corrections are applied using the public
online tool by Placco et al. (2014) 222http://vplacco.pythonanywhere.com/.
Table 13 presents the carbon correction and corrected carbon abundances
obtained for our programme stars. All of our programme stars are found to be
Group I objects in the A(C) vs. [Fe/H] diagram. We therefore expect that the
observed enhancement of carbon and heavy elements in these stars may be
attributed to a binary companion. But none of the programme stars are
confirmed binaries except HE 1523$-$1155.

Table 13: Carbon correction and the corrected carbon abundances obtained for
our programme stars

star name | Carbon correction | log$\epsilon$(C)c
---|---|---
HE 0110$-$0406 | 0.12 | 7.97
HE 1425$-$2052 | 0.13 | 7.89
HE 1428$-$1950 | 0.12 | 8.04
HE 1429$-$0551 | 0.08 | 8.33
HE 1447+0102 | 0.06 | 8.25
HE 1523$-$1155 | 0.11 | 7.94
HE 1528$-$0409 | 0.05 | 8.18

Note : log$\epsilon$(C)c represents the carbon abundance after applying the
correction.

Figure 9: Corrected A(C) vs. [Fe/H] diagram for the compilation of CEMP stars
taken from Yoon et al. (2016). Cyan symbols indicate CEMP-r/s stars: binary
stars are represented by eight sided star and stars with no information about
the binary status are indicated by cross symbols. CEMP-s stars are represented
using magenta symbols: Open and filled circles indicates binary and single
stars respectively. Stars with no information about the binarity are
represented using open triangle. Blue symbols represent Group II CEMP-no
stars: Open and filled hexagons represent binary and single stars
respectively. Stars with no information about the binarity are presented using
open square. Group III CEMP-no stars are represented using green symbols:
binary and single stars are represented by open and filled pentagons
respectively. Skeleton triangle represents stars with no information about
binary status. The symbols used for our programme stars are same as in Figure
7. Representative error bar is shown at the right top corner of the figure.

###  9.2 Possible progenitors of the programme stars

The enhancement of heavy elements in CEMP-s and CEMP-r/s stars are extrinsic
in origin as the evolutionary status of these objects do not support such
enhancements. A brief discussion on CEMP-s and CEMP-r/s stars is presented in
section 1. As discussed earlier, i-process can well explain the observed over
abundances of s- and r-process elements in CEMP-r/s stars. It is found that
the neutron densities required for i-process can be achieved in AGB stars
during proton ingestion episodes (Cowan & Rose, 1977; Campbell & Lattanzio,
2008; Cristallo et al., 2009a; Campbell et al., 2010; Stancliffe et al., 2011;
Banerjee et al., 2018; Clarkson et al., 2018). According to this scenario, a
CEMP-r/s star accrets material synthesized by its binary companion during its
AGB phase of evolution by i-process nucleosynthesis that can enrich its
surface compositions by s- and r- process elements.

Abundance patterns in 20 CEMP-r/s stars could be reproduced by Hampel et al.
(2016) based on their i-process model which supports the i-process as the
probable mechanism for the formation of CEMP-r/s stars. Hampel et al. (2016)
have calculated the yields of heavy elements from i-process nucleosynthesis
occurring at different neutron densities based on single-zone nuclear network
calculations. We have compared the abundance patterns estimated in the
CEMP-r/s stars, HE 1447+0102 and HE 1523$-$1155 with the model yields of
Hampel et al. (2016) as shown in Figure 10. The abundance of neutron-capture
elements in HE 1447+0102 and HE 1523$-$1155 match well with the expected
yields from the i-process nucleosynthesis at a neutron density n $\sim$ 1013
with a dilution factor $\sim$ 0.98.

Figure 10: Best-fitting i-process model (solid blue curve) for HE 1447+0102
and HE 1523$-$1155\. The points with error bars indicate the observed
abundances.

We have compared the abundance patterns of CEMP-s stars in our sample with the
model predictions of AGB stars provided by FRANEC Repository of Updated
Isotopic Tables & Yields (FRUITY) models (Cristallo et al., 2009b, 2011,
2015); the data are publicly available at http://fruity.oa-teramo.inaf.it/
following the detailed procedure presented in Shejeelammal et al. (2021) . The
best fits obtained for the observed abundances are shown in Figures 11 and 12.
For the object HE 0110-0406, the best fit is obtained for the model with
metallicity, z = 0.001, and mass, M = 1.5M⊙. For HE 1429-0551, and HE
1528-0409, the model with z = 0.00002 and M = 2M⊙ gives the best fit. For HE
1425-2052, the best fit is obtained for the model with z = 0.00002, and mass,
M = 1.5M⊙. The observed abundance patterns in HE 1428-1950 could not be
matched with any of the available yield values from the FRUITY models.

Figure 11: Best-fitting FRUITY model (solid black curve) for HE 0110$-$0406,
and HE 1425$-$2052\. The points with error bars indicate the observed
abundances.

Figure 12: Best-fitting FRUITY model (solid black curve) for HE 1429$-$0551,
and HE 1528$-$0409\. The points with error bars indicate the observed
abundances.

The abundance of oxygen and the [Sr/Ba] ratio can be used as indicators of the
possible formation sites of a star (Choplin et al., 2017). Choplin et al.
(2017) analysed the abundance patterns of four single CEMP-s stars based on
the models of rotating and non-rotating massive stars with masses 10-150 M⊙.
They find that the models of fast rotating massive stars could reproduce the
observed abundance patterns in three of the CEMP-s stars in their sample.
Their studies show that CEMP-s stars with fast rotating massive star as the
progenitor show oxygen abundance in the range between 1.5 to 2. Karakas (2010)
predict $-0.2$ $<$ [O/Fe] $<$ 1.2 for CEMP-s stars with AGB progenitors.
Studies by Frischknecht et al. (2012) show that [Sr/Ba] $>$ 0 for massive
rotating stars. Choplin et al. (2017) also support this idea and they claim
that [Sr/Ba] ratio is higher in fast rotating massive stars than in AGB stars.

We could determine [Sr/Ba] ratio in HE 0110$-$0406, HE 1425$-$2052, HE
1447+0102 and HE 1523$-$1155\. For all of them the [Sr/Ba] ratio is found to
be $<$ 0\. This indicates an AGB star as the source of the peculiar abundance
patterns in these stars. While the [Sr/Ba] ratio in HE 1447+0102 indicates an
AGB star as the progenitor, the [O/Fe] measured in this object is the
characteristic of a fast rotating massive star progenitor. This shows that
this object may be polluted by multiple events. The [O/Fe] estimated in the
remaining stars represent an AGB progenitor except HE 1429$-$0551 and HE
1447+0102. We could measure both [Sr/Ba] ratio and [O/Fe] only in three of the
programme stars, HE 1425$-$2052, HE 1447+0102 and HE 1523$-$1155 and their
locations in the [Sr/Ba] vs. [O/Fe] diagram are shown in Figure 13.

Estimation of [ls/hs] ratio is another way to find the possible progenitor of
a star. While the models of fast rotating massive stars predict [ls/hs] $\geq$
0 (Choplin et al., 2017; Chiappini, 2013; Cescutti et al., 2013), the AGB
models predict [ls/hs] $<$ 0 (Abate et al., 2015). We could measure the
[ls/hs] ratio in HE 0110$-$0406, HE 1425$-$2052, HE 1428$-$1950 and HE
1523$-$1155\. For all of them, [ls/hs] ratio is found to be $<$ 0\. This
indicates an AGB progenitor for these objects.

Figure 13: Locations of the programme stars HE 1425$-$2052, HE 1447+0102 and
HE 1523$-$1155 in [Sr/Ba] vs. [O/Fe] diagram. Symbols used for the programme
stars are same as in Figure 7. Filled circles (black) represent CEMP stars
from Hansen et al. (2019)

Elemental abundance ratios such as [Mg/C], [Sc/Mn], [C/Cr] and [Ca/Fe] can
also be used to examine whether the star is polluted by a single event or
several pollution events. Hartwig et al. (2018) presented a novel diagnostic
to identify stars that are formed from the gas enriched by only one previous
supernova based on their cosmological models. They showed that [Mg/C] $<$
$-1.0$, [Sc/Mn] $<$ 0.50, [C/Cr] $>$ 0.50 and [Ca/Fe] $>$ 2 represent mono-
enrichment of a star. As an example we have shown the locations of our
programme stars in the [Mg/C] vs. [Fe/H] diagram (Figure 14). All the
programme stars except HE 1523$-$1155 fall in the region representing mono-
enrichment. Hansen et al. (2019) show that some of the CEMP-s stars in their
sample are also mono-enriched. The estimated value of [Ca/Fe] in all our
programme stars show that they are all multi-enriched as expected for CEMP-s
and CEMP-r/s stars. The value of [Sc/Mn] measured in all the programme stars
except HE 1425$-$2052 show that they are all mono-enriched. The estimated
value of [C/Cr] ratio in the programme stars also show that they are mono-
enriched except HE 1428$-$1950\. Various diagnostic elemental abundance ratios
estimated in each of the programme star show that their surface composition is
both mono- and multi- enriched which does not fit with the results from
Hartwig et al. (2018). This indicates that the different diagnostic elements
in our programme stars are not originated from the previous supernova and it
may be attributed to their binary companion.

Figure 14: Locations of programme stars in [Mg/C] vs. [Fe/H] plot. Symbols
used for the programme stars are same as in Figure 7. Filled circles represent
the stars from literature (Roederer et al., 2014; Plez & Cohen, 2005;
Christlieb et al., 2004; Hansen et al., 2014; Frebel, 2018; Yong et al., 2013;
Bonifacio et al., 2015; Bessell et al., 2015; Lucatello et al., 2003; Barklem
et al., 2005; Cohen et al., 2006; Aoki et al., 2007; Goswami et al., 2006;
Karinkuzhi & Goswami, 2015; Purandardas et al., 2019b, a; Goswami et al.,
2021; Shejeelammal et al., 2021)

###  9.3 Mixing diagnostic

As all the program stars are luminous objects with log(L/L⊙) $>$ 2.50, an
extra mixing can happen and modify the chemical compositions of the stars
(Gratton et al., 2000; Spite et al., 2005). We have therefore examined any
possibility of internal mixing processes that could have modified the observed
surface compositions of our programme stars using estimates of [C/N] ratios
and carbon isotopic ratios. None of the programme stars show any signatures of
internal mixing based on [C/N] ratios (Figure 15). Spite et al. (2006) noted
that [C/N] ratio is not a clean indicator of mixing as the abundances of
carbon and nitrogen in the interstellar medium show large variations. In such
cases, carbon isotopic ratio can be used as a good indicator of mixing as it
is high in primordial matter and it is insensitive to the choice of
atmospheric parameters for the stars (Spite et al., 2006). We could determine
${}^{12}C$/13C ratio in all our programme stars and the estimated values of
12C/13C show that HE 1425$-$2052, HE 1523$-$1155 and HE 1528$-$0409 are well
mixed objects (Figure 16).

Figure 15: Positions of the programme stars in the [C/N] vs. Teff diagram.
Symbols used for the programme stars are same as in Figure 7. Filled circles
represent the stars from literature (Spite et al., 2006; Aoki et al., 2007;
Goswami et al., 2016; Hansen et al., 2016b, 2019; Purandardas et al., 2019b,
a; Goswami et al., 2021; Shejeelammal et al., 2021) Figure 16: Positions of
the programme stars in 12C/13C vs. Teff diagram. Symbols used for the
programme stars are same as in Figure 7. Open circles represent the stars from
Spite et al. (2006) and Aoki et al. (2007)

##  10 Conclusion

Results from follow-up high resolution spectral analysis of seven potential CH
star candidates from Goswami (2005) and Goswami et al. (2010a) are presented.
Our analysis shows one object to be metal-poor and the remaining six are very
metal-poor stars. For the objects HE 0110$-$0406, HE 1425$-$2052 and HE
1428$-$1950, this work presents for the first time a detailed abundance
analysis study. Our analysis shows that the objects HE 1447+0102 and HE
1523$-$1155 to be CEMP-r/s stars. The abundances of neutron-capture elements
in HE 1447+0102 and HE 1523$-$1155 show good match with the yields from the
i-process nucleosynthesis at a neutron density n $\sim$ 1013. Based on our
analysis, the objects HE 1425$-$2052, HE 1429$-$0551 and HE 1528$-$0409 are
found to be CEMP-s stars. The object HE 1428$-$1950 could not be classified
based on CEMP stars classification scheme as we could not estimate the
abundance of Ba in this object.

The locations of the programme stars in the absolute carbon abundance, A(C)
vs. [Fe/H] diagram show that, they all belong to Group I objects. As most of
the Group I objects are known as binaries, the locations of the stars among
the Group I objects may indicate binarity. We note that, the object HE
1523$-$1155 in our sample is a confirmed binary (Hansen et al., 2016a).

The estimated [O/Fe] values indicate AGB stars as possible progenitors of the
programme stars except for HE 1429$-$0551 and HE 1447+0102. For these two
objects, the estimated oxygen abundance characterises fast rotating massive
star progenitors.

Other useful indicators of the formation sources of the stars are [Sr/Ba] and
[hs/ls] ratios. The estimated [Sr/Ba] ratio in HE 0110$-$0406, HE 1425$-$2052,
HE 1447+0102 and HE 1523$-$1155 points at AGB stars as possible progenitors.
The estimated [hs/ls] ratio in HE 0110$-$0406, HE 1425$-$2052, HE 1428$-$1950
and HE 1523$-$1155 also indicate AGB stars as possible progenitors.

We have also studied the chemical enrichment histories of the programme stars
based on various other elemental abundance ratios such as [Mg/C], [Sc/Mn] and
[C/Cr]. While HE 1425$-$2052 is found to be multi-enriched, remaining stars
are all found to be mono-enriched based on [Sc/Mn] ratio. The estimated [C/Cr]
ratio in the programme stars also shows that they are mono-enriched except for
HE 1428$-$1950\. All the stars are also found to be mono-enriched based on
[Mg/C] ratio, except for HE 1523$-$1155\. The estimated Ca abundance of the
stars show that their surface composition is enriched by multiple pollution
events.

The programme stars are luminous with log(L/L⊙) $>$ 2\. In order to understand
whether any internal mixing processes have modified their surface
compositions, we have looked for signatures of internal mixing in our
programme stars based on the estimates of the [C/N] and 12C/13C ratios. While
the [C/N] ratio estimated in the programme stars show that they are all
unmixed, the estimated carbon isotopic ratios show that HE 1425$-$2052, HE
1523$-$1155 and HE 1528$-$0409 are well mixed.

Kinematic analysis shows that the objects HE 0110$-$0406 and HE 1447+0102 to
be thick disk objects and the remaining objects belong to the halo population
of the Galaxy.

The estimates of abundance ratios presented in this paper can be used for
theoretical studies to constrain the physics and the nucleosynthesis occurring
in low-mass AGB stars. Abundance results of CEMP-r/s stars will also provide
observational constrains for i-process nucleosynthesis which is believed to be
the possible source of the origin of the peculiar chemical abundance patterns
observed in CEMP-r/s stars.

##  11 ACKNOWLEDGEMENT

Funding from the DST SERB project EMR/2016/005283 is gratefully acknowledged.
We are thankful to Melanie Hampel for providing us with the i-process yields
in the form of number fractions, and Partha Pratim Goswami for helping us to
generate the model fits. We thank the referee for many constructive
suggestions and useful comments on the manuscript which improved this
paper.This work made use of the SIMBAD astronomical database, operated at CDS,
Strasbourg, France, the NASA ADS, USA and data from the European Space Agency
(ESA) mission Gaia (https://www.cosmos.esa.int/gaia), processed by the Gaia
Data Processing and Analysis Consortium (DPAC,
https://www.cosmos.esa.int/web/gaia/dpac/consortium).

##  12 Data Availability

The data underlying this article will be shared on reasonable request to the
authors.

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##  Appendix A Lines used for deriving elemental abundances

Table 14: Lines used for deriving elemental abundances in our programme stars

Wavelength(Å) | Element |  $E_{low}$(eV) | log gf | HE 0110-0406 | HE 1425-2052 | HE 1428-1950 | HE 1429-0551 | HE 1447+0102 | HE 1523-1155 | HE 1528-0409
---|---|---|---|---|---|---|---|---|---|---
5682.633 | Na I | 2.102 | -0.700 | 79.3(5.53) | 50.1(4.75) | - | - | - | - | -
5889.951 |  | 0.000 | 0.100 | - | - | - | 192.2(4.63) | 162.5(4.56) | - | 130.6(4.06)
5895.920 |  | 0.000 | -0.200 | - | - | - | - | 148.0(4.64) | 180.2(4.53) | 121.5(4.17)
5172.684 | Mg I | 2.711 | -0.402 | - | - | - | 172.2(5.55) | 167.6(5.75) | - | -
5528.405 |  | 4.346 | -0.620 | - | - | - | - | 81.7(6.02) | 101.5(6.13) | 70.0(5.80)
5711.088 |  | 4.346 | -1.833 | 73.3(6.56) | 47.0(5.83) | - | - | - | - | -
4318.652 | Ca I | 1.899 | -0.208 | - | - | - | 41.0(3.82) | - | 59.4(4.11) | -
5512.980 |  | 2.932 | -0.290 | 47.4(4.85) | - | - | - | - | - | -
5581.965 |  | 2.523 | -0.710 | - | 58.5(4.50) | - | - | - | - | -
5588.749 |  | 2.525 | 0.210 | - | - | 122.4(4.84) | 35.8(3.91) | 42.0(4.21) | 59.6(4.28) | 50.3(4.41)
5590.114 |  | 2.521 | -0.710 | 68.7(5.12) | - | - | - | - | - | -
5857.451 |  | 2.932 | 0.230 | - | 107.1(4.68) | - | - | 22.7(4.18) | 35.5(4.23) | 20.5(4.13)
6102.723 |  | 1.879 | -0.890 | 100.1(4.98) | 119.9(4.52) | 88.2(4.55) | 25.0(4.01) | 35.7(4.47) | 46.0(4.32) | 30.6(4.38)
6162.173 |  | 1.899 | 0.100 | - | - | - | - | 66.4(4.15) | 95.2(4.38) | 70.7(4.29)
6169.563 |  | 2.525 | -0.270 | - | 77.6(4.25) | - | - | - | - | -
6439.075 |  | 2.525 | 0.470 | 142.9(5.14) | 172.1(4.65) | - | - | 64.4(4.40) | - | 53.1(4.18)
6471.662 |  | 2.525 | -0.590 | - | 111.5(4.95) | - | - | - | - | -
4415.557 | Sc II | 0.595 | -0.640 | - | 142.1(1.38) | - | 79.5(1.03) | 78.3(1.38) | 88.6(1.24) | -
4431.352 |  | 0.605 | -1.880 | - | 74.8(1.58) | 52.6(1.09) | - | - | 20.6(0.87) | -
5239.813 |  | 1.455 | -0.770 | - | - | - | 30.6(1.01) | 45.6(1.59) | 42.9(1.23) | 36.7(1.45)
5526.790 |  | 1.768 | 0.130 | 123.7(2.00) | - | - | - | - | - | -
6245.637 |  | 1.507 | -0.980 | 85.3(1.99) | 65.8(1.55) | - | - | - | 26.4(1.09) | -
4555.484 | Ti I | 0.848 | -0.488 | 83.9(3.75) | - | - | - | - | - | -
4759.270 |  | 2.255 | 0.514 | - | - | - | - | - | - | -
4820.410 |  | 1.500 | -0.441 | - | - | - | - | - | - | -
4840.874 |  | 0.899 | -0.509 | 72.8(3.58) | 113.0(3.30) | - | - | - | - | -
5024.844 |  | 0.818 | -0.602 | 88.2(3.80) | - | - | - | - | - | -
5210.385 |  | 0.047 | -0.884 | - | - | - | 35.4(2.75) | - | 52.1(2.85) | -
5282.376 |  | 1.052 | -1.300 | 37.0(3.92) | - | - | - | - | - | -
5922.109 |  | 1.046 | -1.466 | - | 23.1(3.11) | - | - | - | - | -
4443.794 | Ti II | 1.080 | -0.700 | - | - | - | - | 86.6(2.63) | - | -
4468.507 |  | 1.130 | -0.600 | - | - | - | 113.3(2.94) | - | 115.2(2.94) | -
4470.857 |  | 1.164 | -2.280 | - | - | 78.2(2.95) | - | - | - | -
4563.761 |  | 1.221 | -0.960 | - | - | - | - | 62.1(2.33) | - | 71.4(2.66)
4583.409 |  | 1.164 | -2.720 | - | 62.9(3.36) | - | - | - | - | -
4764.526 |  | 1.236 | -2.770 | - | - | - | - | - | 22.0(2.95) | -
5185.913 |  | 1.892 | -1.350 | - | 68.4(2.94) | - | 50.5(2.88) | 20.2(2.40) | 46.8(2.80) | -
5188.680 |  | 1.580 | -1.210 | - | - | - | - | - | - | 59.7(2.91)
5226.543 |  | 1.565 | -1.300 | - | - | 154.5(3.39) | 68.0(2.80) | 38.3(2.47) | 78.8(3.03) | 54.5(2.86)
5418.751 |  | 1.581 | -1.999 | - | - | - | - | - | - | -
4379.230 | V I | 0.300 | 0.580 | - | - | - | - | 48.2(2.19) | - | -
4864.731 |  | 0.017 | -0.960 | - | - | 127.7(3.36) | - | - | - | -
5703.575 |  | 1.050 | -0.211 | - | - | 87.5(3.29) | - | - | - | -
5737.059 |  | 1.063 | -0.740 | 47.8(3.50) | 25.3(2.46) | - | - | - | - | -
6119.523 |  | 1.063 | -0.320 | - | 70.3(2.63) | - | - | - | - | -
6216.354 |  | 0.275 | -1.290 | 66.1(3.27) | 39.2(2.11) | - | - | - | - | -
6251.827 |  | 0.286 | -1.340 | - | - | 88.7(3.34) | - | - | - | -
4652.157 | Cr I | 1.003 | -1.030 | - | - | 178.8(4.92) | - | - | - | -
4829.371 |  | 2.544 | -0.810 | - | - | 52.2(4.70) | - | - | - | -
4870.797 |  | 3.079 | 0.050 | 44.2(4.40) | - | - | - | - | - | -
5296.691 |  | 0.982 | -1.400 | - | - | - | - | - | 59.7(3.84) | -
5345.796 |  | 1.003 | -0.980 | 134.7(4.41) | - | - | - | - | 67.7(3.62) | -
5348.315 |  | 1.003 | -1.290 | 98.5(4.06) | 108.5(3.28) | - | 23.8(3.18) | - | - | -
4848.235 | Cr II | 3.864 | -1.140 | - | - | 112.8(4.79) | - | 20.6(3.72) | 27.4(3.64) | -
4451.586 | Mn I | 2.888 | 0.278 | - | - | 105.3(4.26) | - | - | - | -
4470.144 |  | 2.941 | -0.444 | - | 30.0(3.31) | - | - | - | - | -
4739.087 |  | 2.941 | -0.490 | 64.3(4.40) | - | - | - | - | - | -
4765.846 |  | 2.941 | -0.080 | - | 39.8(3.07) | 60.2(3.82) | - | - | - | -
4783.427 |  | 2.298 | 0.042 | - | - | 121.7(3.82) | - | - | - | -
4823.524 |  | 2.319 | 0.144 | - | - | 131.9(3.90) | - | - | - | -
4792.846 | Co I | 3.252 | -0.067 | 25.9(3.68) | - | - | - | - | 26.9(3.86) | -
4813.467 |  | 3.216 | 0.050 | 38.6(3.77) | 19.7(2.97) | 45.5(3.81) | - | - | - | -
5483.344 |  | 1.710 | -1.490 | 59.9(3.73) | - | - | - | - | - | -
4980.166 | Ni I | 3.606 | 0.070 | - | - | - | - | - | 42.5(4.31) | -
6314.653 |  | 1.935 | -1.770 | - | - | 106.9(4.41) | - | - | - | -
6643.629 |  | 1.676 | -2.300 | 129.9(5.36) | 125.5(4.41) | 88.2(4.35) | 21.3(4.01) | 22.9(4.37) | 33.2(4.08) | -
6767.768 |  | 1.826 | -2.170 | - | 138.7(4.63) | - | - | - | 36.9(4.20) | -
4810.528 | Zn I | 4.077 | -0.137 | 56.0(2.84) | 61.4(2.88) | 78.2(2.95) | - | - | 35.6(2.54) | -
4607.327 | Sr I | 0.000 | -0.570 | 63.0(2.49) | 114.1(2.46) | 37.9(1.97) | - | 27.9(2.51) | - | -

Wavelength(Å) | Element |  $E_{low}$(eV) | log gf | HE 0110-0406 | HE 1425-2052 | HE 1428-1950 | HE 1429-0551 | HE 1447+0102 | HE 1523-1155 | HE 1528-0409
---|---|---|---|---|---|---|---|---|---|---
4883.684 | Y II | 1.084 | 0.070 | - | 171.7(1.51) | 142.3(1.30) | - | - | 67.8(0.44) | -
5087.416 |  | 1.084 | -0.170 | - | - | - | - | - | - | -
6613.733 |  | 1.748 | -1.110 | 80.7(2.20) | 40.9(1.46) | - | - | - | - | -
4130.645 | Ba II | 2.722 | 0.680 | - | - | - | - | - | 71.7(1.63) | -
5853.668 |  | 0.604 | -1.000 | 145.1(1.41) | - | - | 125.0(1.74) | 131.7(2.36) | - | 106.3(1.99)
6141.713 |  | 0.703 | -0.076 | - | - | - | 163.9(1.49) | - | 168.2(1.35) | 180.6(2.16)
6496.897 |  | 0.604 | -0.377 | - | - | - | - | - | - | 154.7(2.04)
4322.503 | La II | 0.173 | -1.120 | - | - | - | 40.0(0.33) | 47.9(0.90) | 36.4(0.17) | 39.8(0.75)
4662.498 |  | 0.000 | -1.240 | - | 149.7(1.34) | - | - | - | - | -
4748.726 |  | 0.926 | -0.860 | - | 101.8(1.43) | - | - | 39.7(1.20) | - | -
4921.776 |  | 0.244 | -0.680 | - | - | - | 69.4(0.60) | 68.9(1.07) | 64.9(0.40) | 53.0(0.67)
5259.379 |  | 0.173 | -1.760 | - | 143.4(1.79) | - | - | - | - | -
5303.528 |  | 0.321 | -1.430 | - | 124.9(1.39) | - | - | - | 20.0(0.13) | -
6320.376 |  | 0.172 | -1.610 | - | - | - | - | 38.0(0.94) | - | -
6390.477 |  | 0.321 | -1.450 | - | - | - | - | - | - | 20.4(0.56)
4336.244 | Ce II | 0.703 | -0.564 | - | - | - | 25.6(0.81) | - | - | -
4418.780 |  | 0.863 | 0.177 | - | - | 68.2(0.44) | - | - | - | -
4427.916 |  | 0.535 | -0.460 | - | - | 59.7(0.52) | - | - | - | 28.2(0.96)
4523.075 |  | 0.516 | -0.304 | - | - | - | - | - | 47.6(0.74) | -
4539.745 |  | 0.327 | -0.459 | - | - | - | - | - | 54.5(0.84) | 50.6(1.31)
4544.953 |  | 0.417 | -0.974 | - | - | - | - | 31.3(1.34) | - | -
4562.359 |  | 0.477 | 0.081 | - | - | - | 82.9(1.37) | - | 55.8(0.51) | -
4725.069 |  | 0.521 | -1.204 | 62.5(1.70) | 86.1(1.70) | - | 20.1(1.04) | 20.7(1.33) | - | -
5330.556 |  | 0.869 | -0.760 | 43.5(1.27) | - | 40.2(0.76) | 20.4(0.95) | 30.7(1.54) | 20.2(0.72) | -
5556.941 |  | 1.813 | -0.405 | 18.1(1.50) | - | - | - | - | - | -
5219.045 | Pr II | 0.795 | -0.240 | - | - | - | 23.6(-0.63) | - | - | -
5259.728 |  | 0.633 | 0.080 | - | - | - | - | - | - | -
5292.619 |  | 0.648 | -0.300 | 71.8(0.88) | - | - | - | 22.4(0.50) | - | 20.5(0.38)
5322.772 |  | 0.482 | -0.315 | 75.4(0.76) | - | - | - | 25.4(0.41) | - | 26.0(0.48)
5892.251 |  | 1.439 | -0.352 | - | - | - | - | - | - | -
6278.676 |  | 1.196 | -0.630 | - | - | - | - | - | - | -
4412.256 | Nd II | 0.063 | -1.420 | - | - | - | - | 31.9(1.22) | - | -
4446.384 |  | 0.204 | -0.590 | - | - | - | 48.3(0.54) | 52.7(1.09) | 50.8(0.52) | -
4451.563 |  | 0.380 | -0.040 | - | - | 105.0(0.37) | 75.4(0.97) | 68.2(1.24) | - | -
4516.346 |  | 0.320 | -0.950 | - | - | - | 40.3(0.84) | 46.5(1.39) | - | 26.2(0.94)
4797.153 |  | 0.559 | -0.950 | - | 99.5(1.45) | - | - | - | 25.0(0.66) | -
4825.478 |  | 0.182 | -0.860 | - | - | 69.3(0.25) | - | 46.7(1.10) | - | -
4859.026 |  | 0.320 | -0.830 | - | 143.1(1.65) | 61.7(0.29) | - | 35.0(0.94) | 50.6(0.82) | -
5212.361 |  | 0.204 | -0.870 | 93.0(1.24) | 169.2(1.79) | - | 46.6(0.67) | 57.9(1.39) | 42.2(0.48) | 34.5(0.87)
5255.506 |  | 0.204 | -0.820 | - | 146.2(1.39) | - | 53.3(0.76) | - | 52.4(0.65) | 26.7(0.62)
5287.133 |  | 0.744 | -1.300 | - | 91.0(1.84) | - | - | - | - | -
5293.163 |  | 0.822 | -0.060 | - | - | - | 47.7(0.60) | 64.8(1.48) | 47.1(0.52) | 39.8(0.87)
5319.815 |  | 0.550 | -0.210 | - | - | - | 56.3(0.62) | 69.5(1.45) | 60.4(0.64) | 57.7(1.19)
5356.967 |  | 1.264 | -0.250 | 76.9(1.62) | - | - | - | 27.9(1.20) | - | -
5442.264 |  | 0.680 | -0.910 | 81.0(1.61) | - | - | - | - | - | -
5702.238 |  | 0.744 | -0.770 | 81.4(1.51) | 138.7(1.88) | - | - | 37.1(1.33) | 19.0(0.44) | -
4220.661 | Sm II | 0.544 | -1.114 | - | - | - | - | 24.0(0.99) | - | -
4424.337 |  | 0.485 | -0.260 | - | - | - | - | - | 45.5(0.16) | 55.7(0.99)
4434.318 |  | 0.378 | -0.576 | - | 158.1(1.70) | 59.5(-0.04) | 40.3(0.29) | - | 51.1(0.48) | -
4499.475 |  | 0.248 | -1.413 | - | 92.7(1.24) | - | - | - | 20.0(0.33) | -
4519.630 |  | 0.543 | -0.751 | - | - | - | - | 38.0(0.94) | 32.2(0.39) | -
4566.202 |  | 0.333 | -1.245 | - | 100.3(1.28) | - | - | 38.7(1.22) | 33.4(0.65) | -
4615.444 |  | 0.543 | -1.262 | - | 98.8(1.55) | - | - | - | - | -
4674.593 |  | 0.184 | -1.055 | 65.7(0.72) | - | - | - | - | - | -
4726.026 |  | 0.333 | -1.849 | 25.0(0.92) | - | - | - | - | - | -
4791.580 |  | 0.104 | -1.846 | 47.8(1.07) | 119.5(1.79) | - | - | - | - | -
4854.368 |  | 0.378 | -1.873 | 20.3(0.87) | 79.7(1.61) | - | - | - | - | -
4129.730 | Eu II | 0.000 | 0.204 | - | - | - | - | 74.9(0.11) | - | 22.3(-1.45)
6437.640 |  | 1.319 | -0.276 | - | 70.1(0.49) | - | - | 22.4(0.25) | - | -
6645.064 |  | 1.379 | 0.204 | - | 89.3(0.31) | 64.6(-0.35) | - | - | - | -

The numbers in the parenthesis in columns 5-11 give the derived abundances
from the respective line. log gf values are taken from kurucz atomic line list
(https://www.cfa.harvard.edu).