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# Chemical abundances of 20 barium stars from the OHP spectra

Guochao Yang, Jingkun Zhao, Yanchun Liang, Monique Spite, Francois Spite, Jianrong Shi, Shuai Liu, Nian Liu, Wenyuan Cui, Gang Zhao

► To cite this version:

Guochao Yang, Jingkun Zhao, Yanchun Liang, Monique Spite, Francois Spite, et al. Chemical abundances of 20 barium stars from the OHP spectra. Monthly Notices of the Royal Astronomical Society, 2024, 534, pp.3104-3119. 10.1093/mnras/stae2283. insu-04821649

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# Chemical abundances of 20 barium stars from the OHP spectra

Guochao Yang,<sup>1,2</sup> Jingkun Zhao,<sup>2,*</sup> Yanchun Liang,<sup>2,*</sup> Monique Spite,<sup>3</sup> Francois Spite,<sup>3</sup> Jianrong Shi,<sup>2,4</sup> Shuai Liu,<sup>2</sup> Nian Liu,<sup>1</sup> Wenyuan Cui<sup>5</sup> and Gang Zhao<sup>1,2</sup>

Accepted 2024 September 30. Received 2024 September 29; in original form 2024 April 18

## ABSTRACT

Based on the high resolution and high signal-to-noise spectra, we derived the chemical abundances of 20 elements for 20 barium (Ba-) stars. For the first time, the detailed abundances of four sample stars, namely HD 92482, HD 150430, HD 151101, and HD 177304 have been analysed. Additionally, Ba element abundance has been measured using high-resolution spectra for the first time in six of the other 16 sample stars. Based on the [s/Fe] ratios, the Ba-unknown star HD 115927 can be classified as a strong Ba-star, while the Ba-likely star HD 160538 can be categorized into a mild Ba-star. Consequently, our sample comprises three strong and 17 mild Ba-stars. The light odd-Z metal elements and Fe-peak elements exhibit near-solar abundances. The [α/Fe] ratios demonstrate decreasing trends with increasing metallicity. Moreover, the abundances of neutron-capture (n-capture) elements show significant enhancements in different degrees. Using a threshold of the signed distances to the solar rapid-process (r-process) abundance pattern $d_s = 0.6$, we find that all of our sample stars are normal Ba-stars, indicating that the enhancements of slow-process (s-process) elements should be attributed to material transfer from their companions. We compare the observed n-capture patterns of sample stars with the FRUITY models, and estimate the mass of the Thermally-Pulsing Asymptotic Giant Branch stars that previously contaminated the Ba-stars. The models with low masses can successfully explain the observations. From a kinematic point of view, we note that most of our sample stars are linked with the thin disc, while HD 130255 may be associated with the thick disc.

**Key words:** stars: abundances – stars: atmospheres – stars: chemically peculiar.

## 1 INTRODUCTION

Various nucleosynthesis processes operating in different mass stars are responsible for generating diverse chemical elements (Sneden, Cowan & Gallino 2008). The light metal elements from C to Zn can be produced through charged-particle fusion reactions, whereas the elements heavier than Fe-peak are predominantly synthesized by neutron-capture (n-capture) process (Burbidge et al. 1957; Busso et al. 2001; Ji & Frebel 2018; Cowan et al. 2021). Based on the time-scales of n-capture and radioactive decay of unstable nuclei, the n-capture nucleosynthesis is generally categorized into slow (s-) process, primarily occurring in asymptotic giant branch (AGB) stars and rapid (r-) process, commonly associated with violent astrophysical environments such as Type II supernovae (SNe II) explosions or neutron star mergers (NSMs) (Qian & Wasserburg 2007; Abbott et al. 2017; Watson et al. 2019; Stancliffe 2021). Abundance analysis of chemically peculiar stars can help us investigating the nucleosynthesis processes of elements.

Barium (Ba-) stars are a significant class that shows anomalously strong lines of s-process elements (e.g. Sr, Y, Ba, and La) in their spectra and were first defined by Bidelman & Keenan (1951). Ba-stars are giants or dwarfs with spectral types from ~G to K and effective temperatures from ~4000 to 6000 K. The variations of radial velocity, together with the enhancement of s-process elements, imply that the observed Ba-star should have an unseen white dwarf (WD) companion (McClure, Fletcher & Nemec 1980). In other words, maybe all of the Ba-stars are in binary systems, where the Ba-stars obtained s-process material from the AGB companions, which now evolved into WDs (McClure 1983; Jorissen et al. 1998, 2019). Furthermore, some Ba-stars (e.g. HD 48565 and HD 114520) are found in triple systems (North, Jorissen & Mayor 2000; Escorza et al. 2019).

The chemical abundance features of Ba-stars can provide important information for understanding the different nucleosynthesis processes and the mass accretion in binaries or triples (Cseh et al. 2018; Stancliffe 2021). Based on the strength of Ba II line at $\lambda = 4554$ Å and the scale of Ba intensities proposed by Warner (1965), Lu (1991) compiled a catalogue containing 389 Ba-star candidates. In the past decades, numerous studies have analysed chemical abundances of Ba-stars using high resolution and high

<footnote n="1"><sup>1</sup>School of Physics and Astronomy, China West Normal University, 637002 Nanchong, PR China</footnote>

<footnote n="2"><sup>2</sup>Key Laboratory of Optical Astronomy, National Astronomical Observatories, Chinese Academy of Sciences, 100012 Beijing, PR China</footnote>

<footnote n="3"><sup>3</sup>GEPI, Observatoire de Paris-Meudon, F-92195 Meudon, France</footnote>

<footnote n="4"><sup>4</sup>School of Astronomy and Space Science, University of Chinese Academy of Sciences, 100049 Beijing, PR China</footnote>

<footnote n="5"><sup>5</sup>Department of Physics, Hebei Normal University, 050024 Shijiazhuang, PR China</footnote>

<footnote>*E-mail: zjk@bao.ac.cn (JZ); ycliang@nao.cas.cn (YL)</footnote>
signal-to-noise (S/N) spectra (e.g. Začs 1994; Liang, Zhao & Zhang 2000; Boyarchuk et al. 2002; Liang et al. 2003; Antipova et al. 2004; Yushchenko et al. 2004; Allen & Barbuy 2006a, b; Smiljanic, Porto de Mello & da Silva 2007; Liu, Liang & Deng L. 2009; Pereira & Drake 2009; Pereira et al. 2011; Katime Santrich, Pereira & de Castro 2013).

In recent years, de Castro et al. (2016) conducted a study on a large sample of 182 Ba-stars and determined the abundances of Na, Al, α-elements, Fe-peak elements, and the s-process elements Y, Zr, La, Ce, and Nd. They found that the [s/Fe] and [hs/ls] (where hs refers to the heavier s-process elements, such as Ba, La, Ce, and Nd, and ls for the lighter s-process elements, such as Sr, Y, and Zr) show strong anticorrelations with metallicity (see also Goriely & Mowlavi 2000; Busso et al. 2001). Based on the data set of chemical abundances derived by de Castro et al. (2016), Cseh et al. (2018) compared [hs/ls] ratios of 169 Ba-stars to those from the theoretical AGB models, and found a good consistency between the observed [hs/ls] ratios and theoretical low-mass (≤3–4 M⊙) AGB models, where $^{13}$C is the main neutron source. Kong et al. (2018a) studied three Ba-dwarfs, and provided stellar mass and abundances of light metal, α-, Fe-peak, and s-process elements. They suggested that the formation scenarios of s-process enhancement in Ba-dwarfs are similar to those of Ba-giants, where the s-process elements originate from the WD companions during the AGB phase through Roche-lobe overflow or wind accretion.

Analysing chemical abundances of 18 Ba-stars using high resolution and high S/N ratio spectra from Keck/HIRES, Liu et al. (2021) discovered two new Ba-giants (HD 16178 and HD 22233) and one Ba-subgiant (HD 2946). Roriz et al. (2021a, 2021b) investigated a large sample of 180 Ba-stars, presenting Rb and n-capture elements, e.g. Sr, Nb, Mo, Ru, La, Sm, and Eu, and compared the observed [hs/ls] and [Eu/La] ratios with those from s-process models. They found that the [Nb, Mo, Ru/Sr] and [Ce, Nd, Sm/La] ratios of most sample stars are higher than those of the model predictions, suggesting the need of new theoretical and observational work to explain these features. Considering the mass transfer in Ba binaries, Cseh et al. (2022) compared observed abundance patterns of 28 Ba-stars to the AGB models of Monash (Lugaro et al. 2012; Fishlock et al. 2014; Karakas & Lugaro 2016; Karakas et al. 2018) and FRUITY (Cristallo et al. 2009, 2011, 2015, 2016). They concluded that the abundance patterns of n-capture elements in most of the Ba-stars are well represented by the ejecta of low-mass non-rotating AGB stars, where $^{13}$C is the main source of neutrons. Additionally, they found that some Ba-stars represent abundance features of not only s-process but also other nucleosynthetic processes. By utilizing FRUITY models, Goswami P. & Goswami (2023) conducted mass estimations for the primary companions of 158 Ba-stars, comprising 52 mild and 106 strong Ba-stars. Their analysis revealed that the masses of primary companions in strong Ba-stars reach a peak at 2.5 M⊙ with a standard deviation of 0.51 M⊙, whereas those in mild Ba-stars peak at 3.7 M⊙ with a standard deviation of 1.03 M⊙. Through a comparison of the observed n-capture patterns of four Ba-stars with Monash and FRUITY nucleosynthesis models, Roriz et al. (2024) determined that the low-mass AGB stars, responsible for the pollution, effectively replicate the observed data.

Han et al. (1995) estimated that there may exist 800–22 000 Ba-stars with magnitude $V \le 10.0$ in our Galaxy. The anomalies in elemental abundances of Ba-stars reflect the nature of the nucleosynthetic processes occurring within the interior of AGB stars. Given that the different nucleosynthetic processes inside various AGB stars are not yet fully understood, studying more Ba-stars is

of particular interest. To explore the abundance characteristics of Ba-stars in detail, we obtained high resolution and high S/N spectra for 31 Ba-stars using the 2.16-m telescope at Xinglong station of the National Astronomical Observatories, Chinese Academy of Sciences (NAOC), and 20 Ba-stars using the 1.93-m telescope at the Haute-Provence Observatory (OHP, France) during 2001 and 2005. The chemical abundances of the former 31 Ba-stars have been analysed in our previous papers (see details in Liang et al. 2003; Liu et al. 2009; Yang et al. 2016).

In this work, we observed the spectra of the latter 20 Ba-stars at OHP and derived the abundances of light metal, α-, Fe-peak and n-capture elements. In Section 2, we present the observations and data reductions of the sample stars. The stellar atmospheric parameters are provided in Section 3. The results are shown in Section 4. Discussions are presented in Section 5. In Section 6, we conduct a kinematic investigation of the sample stars. The conclusions are drawn in Section 7.

## 2 OBSERVATIONS AND DATA REDUCTION

The sample stars investigated in this paper were selected from Lu (1991) and Jorissen et al. (1998). The basic parameters are listed in Table 1. The spectra of the sample stars were obtained using the ELODIE spectrograph installed on the 1.93-m telescope at OHP. The high resolution and high S/N ratio spectra were acquired during a single observing run in 2003 (see details in Table 1). The spectra cover a wavelength range of 4000–7500 Å with a resolving power of $R \sim 48\,000$. For most of the sample stars, the S/N exceeds 100 in the entire wavelength region. A portion of the spectrum of the sample star HD 119650 is shown in Fig. 1.

The 1D wavelength-calibrated spectra were retrieved from the ELODIE<fnref n="1" /> archive (Moultaka et al. 2004). We normalized the spectra by fitting the continuum using cubic spline functions with smooth parameters and computed stellar radial velocities by fitting the profiles of approximately 80 pre-selected lines (Wang et al. 2011). The measurements of equivalent width (EW) are performed using two methods: Gaussian function fitting and direct integration. A Gaussian fitting is adopted for weak lines, while direct integration is used for unblended strong lines (Zhao et al. 2000). Lines with $20\,\text{mÅ} < \text{EW} < 120\,\text{mÅ}$ have been chosen for the following abundance calculation of Fe. For other elements, the lines with $20\,\text{mÅ} < \text{EW} < 250\,\text{mÅ}$ have been adopted due to limited availability. Furthermore, close attention is paid to ensuring the consistency of the continuum level.

For assessing the reliability of the EW measurements in this study, we take HD 119650 and HD 125079 as sample stars and compared the EWs of common lines measured in our work with those reported by Mahanta et al. (2016) and Karinkuzhi & Goswami (2014). The comparison results are shown in Fig. 2, revealing a mean deviation of about $3.5\,\text{mÅ}$ for the EWs of the common lines between our study and the previous investigations.

## 3 STELLAR ATMOSPHERIC PARAMETERS

The determination of the stellar atmospheric parameters, i.e. the effective temperature $T_{\text{eff}}$, surface gravity $\log g$, metallicity [Fe/H], and micro-turbulent velocity $\xi_t$, is essential for establishing elemental abundances. The effective temperatures were determined using two methods. Initially, the effective temperatures were obtained from

<footnote n="1">¹http://atlas.obs-hp.fr/elodie/</footnote>
<table><caption>**Table 1.** Log of observations and basic data of the sample stars.</caption><thead><tr><th>Star name</th><th>Date</th><th>Exp. time (s)</th><th>Ba</th><th>Sp.</th><th>V</th><th>K</th><th>Source</th></tr></thead><tbody><tr><td>HD 77912</td><td>2003 May 23</td><td>300</td><td>0.2m</td><td>G7II</td><td>4.54</td><td>2.40</td><td>1</td></tr><tr><td>HD 85503</td><td>2003 May 23</td><td>100</td><td>1.0:</td><td>K2III</td><td>3.88</td><td>1.22</td><td>1</td></tr><tr><td>HD 92482</td><td>2003 May 23</td><td>3600</td><td>3.0</td><td>G8III</td><td>8.62</td><td>6.32</td><td>1</td></tr><tr><td>HD 119650</td><td>2003 May 23</td><td>1800</td><td>0.3m</td><td>K1II</td><td>7.59</td><td>4.91</td><td>1</td></tr><tr><td>HD 125079</td><td>2003 May 23</td><td>3600</td><td>2.0m</td><td>G2III</td><td>8.63</td><td>6.61</td><td>1</td></tr><tr><td>HD 150430</td><td>2003 May 23</td><td>3500</td><td>0.5:</td><td>K0</td><td>8.19</td><td>5.88</td><td>1</td></tr><tr><td>HD 160507</td><td>2003 May 23</td><td>1100</td><td>1.0:</td><td>G5III</td><td>6.56</td><td>4.39</td><td>1</td></tr><tr><td>HD 160538</td><td>2003 May 23</td><td>1200</td><td>bl</td><td>K0</td><td>6.63</td><td>4.09</td><td>2</td></tr><tr><td>HD 177304</td><td>2003 May 23</td><td>3600</td><td>0.5:</td><td>K</td><td>8.61</td><td>5.81</td><td>1</td></tr><tr><td>HD 107950</td><td>2003 May 24</td><td>920</td><td>0.2</td><td>G7III</td><td>4.76</td><td>2.82</td><td>1</td></tr><tr><td>HD 133208</td><td>2003 May 24</td><td>300</td><td>0.5</td><td>G8III</td><td>3.52</td><td>1.34</td><td>1</td></tr><tr><td>HD 115927</td><td>2003 May 24</td><td>3600</td><td>?</td><td>K0II</td><td>7.70</td><td>5.18</td><td>1</td></tr><tr><td>HD 130255</td><td>2003 May 24</td><td>3600</td><td>1.0c</td><td>G9V</td><td>8.64</td><td>5.83</td><td>2</td></tr><tr><td>HD 168532</td><td>2003 May 24</td><td>900</td><td>0.5</td><td>K3III</td><td>5.28</td><td>1.85</td><td>1</td></tr><tr><td>HD 151101</td><td>2003 May 24</td><td>600</td><td>0.3</td><td>K0III</td><td>4.84</td><td>2.11</td><td>1</td></tr><tr><td>HD 176524</td><td>2003 May 24</td><td>600</td><td>0.2</td><td>K0III</td><td>4.81</td><td>2.32</td><td>1</td></tr><tr><td>HD 176670</td><td>2003 May 24</td><td>600</td><td>0.5</td><td>K2.5III</td><td>4.93</td><td>1.57</td><td>1</td></tr><tr><td>HD 101079</td><td>2003 May 27</td><td>3600</td><td>1.0</td><td>K1</td><td>8.18</td><td>6.02</td><td>2</td></tr><tr><td>HD 131873</td><td>2003 May 27</td><td>120</td><td>0.3</td><td>K4III</td><td>2.08</td><td>-1.39</td><td>1</td></tr><tr><td>HD 136138</td><td>2003 May 27</td><td>900</td><td>0.3</td><td>G8III</td><td>5.69</td><td>3.61</td><td>1</td></tr></tbody></table>

Notes. Ba intensity in Col. 4. c: certain Ba-stars, m: marginal Ba-stars, bl: Ba-likely stars, ':' refers to the case of Ba intensity being uncertain, and '?' means the Ba intensity being unidentified.

References: 1. Lu (1991); 2. Jorissen et al. (1998).

<figure><img src="image_4.png" /><figcaption>**Figure 1.** A portion of spectrum of the sample star HD 119650.</figcaption></figure>

<figure><img src="image_5.png" /><figcaption>**Figure 2.** Comparisons of EWs of common lines measured in this work with those of Mahanta et al. (2016) and Karinkuzhi & Goswami (2014) for the sample star HD 119650 (the left-hand panel) and HD 125079 (the right-hand panel), respectively. The solid lines are the linear fits to the measured points. The dashed lines are one-to-one correlations.</figcaption></figure>

the photometric colour V-K utilizing the calibration provided by Alonso, Arribas & Martínez-Roger (1999). The V-K colour data were taken from the SIMBAD<fnref n="2" /> data base, and the colour excess E(V-K) was computed from the 3D dust maps of Green

et al. (2019). The derived effective temperatures and corresponding uncertainties are listed in Table 2 (Columns 2 and 3). The large uncertainties in the effective temperatures of HD 77912, HD 107950, HD 136138, HD 151101, HD 160538, and HD 176524 can be attributed to the large uncertainties in their V-K values (≥0.30 dex). The effective temperatures can also be spectroscopically determined by ensuring that the iron abundances derived from Fe I lines are independent of the excitation potentials, and we get the final results when the slope of [Fe I/H] versus excitation potential is less than 0.005. We estimated the uncertainties from this on temperature value, and the derived results are presented in Columns 4 and 5 of Table 2. In our analysis, the spectroscopic temperatures are derived from a vast array of clean Fe lines to calculate the elemental abundances of the sample stars.

The surface gravities were also derived using two methods. The first method is based on the formulas as follows:

$$ \log \frac{g}{g_{\odot}} = \log \frac{M}{M_{\odot}} + 4 \log \frac{T_{\text{eff}}}{T_{\text{eff}\odot}} + 0.4(M_{\text{bol}} - M_{\text{bol}\odot}) \quad (1) $$

$$ M_{\text{bol}} = V - A_{V} + 5 + 5 \log \pi + \text{BC}, \quad (2) $$

here, M, M<sub>bol</sub>, V, A<sub>V</sub>, π, and BC represent the stellar mass, absolute bolometric magnitude, apparent magnitude, interstellar extinction, parallax, and bolometric correction, respectively. The stellar mass M was derived from the Y<fnref n="2" /> (Yonsei-Yale) Stellar Models (Yi, Kim & Demarque 2003). The parallaxes of HD 131873 and the other sample stars were adopted from Hipparcos and Gaia DR3 (Gaia Collaboration et al. 2023), respectively. The bolometric correction BC was calculated using the relation provided by Alonso et al. (1999). For the Sun, T<sub>eff</sub>, log g<sub>☉</sub>, and M<sub>bol</sub>☉ were adopted as 5770 K, 4.44 dex, and 4.75 mag, respectively. The second method for determining log g involves using the ionization equilibrium. We obtained the final log g with this method by requiring that the difference of iron abundances from Fe I and Fe II lines is less than 0.05 dex. The corresponding uncertainties in surface gravities were estimated using the two formulas mentioned above and by adjusting

<footnote n="2">²http://simbad.u-strasbg.fr/Simbad</footnote>
<table><caption>**Table 2.** Physical parameters of the sample stars.</caption><thead><tr><th>Star name</th><th><i>T</i><sub>eff</sub> (K)<br>(V − K)</th><th>Δ<i>T</i><sub>eff</sub> (K)<br>(V − K)</th><th><i>T</i><sub>eff</sub> (K)<br>(Spec)</th><th>Δ<i>T</i><sub>eff</sub> (K)<br>(Spec)</th><th>log <i>g</i><br>(Parallax)</th><th>Δlog <i>g</i><br>(Parallax)</th><th>log <i>g</i><br>(Spec)</th><th>Δlog <i>g</i><br>(Spec)</th><th>[Fe I/H]</th><th>Δ[Fe/H]</th><th>[Fe II/H]</th><th>ξ<sub>t</sub> (km s<sup>−1</sup>)</th><th>Δξ<sub>t</sub> (km s<sup>−1</sup>)</th></tr></thead><tbody><tr><td>HD 77912</td><td>4995</td><td>304</td><td>5021</td><td>59</td><td>1.99</td><td>0.11</td><td>2.06</td><td>0.10</td><td>0.05</td><td>0.13</td><td>0.05</td><td>2.05</td><td>0.16</td></tr><tr><td>HD 85503</td><td>4493</td><td>124</td><td>4635</td><td>70</td><td>2.62</td><td>0.07</td><td>2.65</td><td>0.13</td><td>0.38</td><td>0.17</td><td>0.36</td><td>1.50</td><td>0.15</td></tr><tr><td>HD 92482</td><td>4778</td><td>33</td><td>4900</td><td>46</td><td>2.34</td><td>0.03</td><td>2.36</td><td>0.08</td><td>−0.24</td><td>0.15</td><td>−0.22</td><td>1.50</td><td>0.10</td></tr><tr><td>HD 101079</td><td>4982</td><td>40</td><td>5100</td><td>67</td><td>2.70</td><td>0.02</td><td>2.90</td><td>0.12</td><td>0.25</td><td>0.15</td><td>0.25</td><td>1.40</td><td>0.15</td></tr><tr><td>HD 107950</td><td>5184</td><td>359</td><td>5247</td><td>63</td><td>2.62</td><td>0.13</td><td>2.55</td><td>0.10</td><td>0.23</td><td>0.13</td><td>0.23</td><td>1.60</td><td>0.13</td></tr><tr><td>HD 115927</td><td>4581</td><td>26</td><td>4881</td><td>40</td><td>2.29</td><td>0.05</td><td>2.68</td><td>0.08</td><td>−0.21</td><td>0.15</td><td>−0.19</td><td>1.25</td><td>0.18</td></tr><tr><td>HD 119650</td><td>4450</td><td>16</td><td>4510</td><td>38</td><td>1.83</td><td>0.03</td><td>2.05</td><td>0.12</td><td>−0.22</td><td>0.15</td><td>−0.22</td><td>1.52</td><td>0.15</td></tr><tr><td>HD 125079</td><td>5351</td><td>36</td><td>5334</td><td>50</td><td>3.15</td><td>0.02</td><td>3.20</td><td>0.15</td><td>−0.16</td><td>0.14</td><td>−0.16</td><td>1.20</td><td>0.15</td></tr><tr><td>HD 130255</td><td>4436</td><td>24</td><td>4509</td><td>72</td><td>1.71</td><td>0.06</td><td>1.90</td><td>0.15</td><td>−0.86</td><td>0.14</td><td>−0.88</td><td>1.30</td><td>0.16</td></tr><tr><td>HD 131873</td><td>4007</td><td>75</td><td>4150</td><td>80</td><td>1.40</td><td>0.06</td><td>1.40</td><td>0.15</td><td>−0.08</td><td>0.14</td><td>−0.09</td><td>1.65</td><td>0.10</td></tr><tr><td>HD 133208</td><td>4931</td><td>166</td><td>5174</td><td>40</td><td>2.52</td><td>0.06</td><td>2.54</td><td>0.11</td><td>0.23</td><td>0.13</td><td>0.23</td><td>1.70</td><td>0.20</td></tr><tr><td>HD 136138</td><td>5123</td><td>375</td><td>5100</td><td>33</td><td>2.76</td><td>0.14</td><td>2.92</td><td>0.05</td><td>0.07</td><td>0.15</td><td>0.10</td><td>1.45</td><td>0.17</td></tr><tr><td>HD 150430</td><td>4891</td><td>31</td><td>4850</td><td>40</td><td>2.60</td><td>0.03</td><td>2.57</td><td>0.05</td><td>−0.26</td><td>0.13</td><td>−0.21</td><td>1.35</td><td>0.14</td></tr><tr><td>HD 151101</td><td>4482</td><td>244</td><td>4504</td><td>34</td><td>1.78</td><td>0.12</td><td>1.81</td><td>0.06</td><td>0.13</td><td>0.14</td><td>0.16</td><td>1.70</td><td>0.05</td></tr><tr><td>HD 160507</td><td>4936</td><td>48</td><td>5044</td><td>47</td><td>2.85</td><td>0.03</td><td>2.91</td><td>0.07</td><td>0.09</td><td>0.14</td><td>0.13</td><td>1.40</td><td>0.14</td></tr><tr><td>HD 160538</td><td>4572</td><td>340</td><td>4812</td><td>89</td><td>2.75</td><td>0.16</td><td>3.00</td><td>0.15</td><td>−0.02</td><td>0.13</td><td>−0.05</td><td>1.70</td><td>0.20</td></tr><tr><td>HD 168532</td><td>4293</td><td>190</td><td>4307</td><td>39</td><td>1.30</td><td>0.10</td><td>1.90</td><td>0.08</td><td>0.33</td><td>0.16</td><td>0.34</td><td>1.65</td><td>0.08</td></tr><tr><td>HD 176524</td><td>4621</td><td>270</td><td>4700</td><td>24</td><td>2.19</td><td>0.12</td><td>2.20</td><td>0.07</td><td>0.13</td><td>0.15</td><td>0.15</td><td>1.55</td><td>0.12</td></tr><tr><td>HD 176670</td><td>4067</td><td>118</td><td>4300</td><td>52</td><td>1.05</td><td>0.08</td><td>1.35</td><td>0.18</td><td>0.09</td><td>0.14</td><td>0.06</td><td>1.80</td><td>0.05</td></tr><tr><td>HD 177304</td><td>5093</td><td>32</td><td>5200</td><td>41</td><td>2.34</td><td>0.02</td><td>2.60</td><td>0.10</td><td>0.23</td><td>0.14</td><td>0.24</td><td>1.65</td><td>0.15</td></tr></tbody></table>

<figure><img src="image_6.png" /><figcaption>**Figure 3.** The determination of atmospheric parameters of the sample stars HD 119650 (the upper two panels) and HD 125079 (the lower two panels) based on the spectra. The filled and open circles refer to the Fe abundances computed from Fe I and Fe II lines, respectively. The straight lines are the linear-fitting results of Fe I abundances. Panels (a) and (c) show the excitation equilibrium of Fe abundances derived from Fe I lines. Panels (b) and (d) show the ionization equilibrium of Fe abundances from Fe I and Fe II lines. Furthermore, the EW-independence of Fe I abundances is also shown in panels (b) and (d).</figcaption></figure>

log *g* to force the abundance difference between Fe I and Fe II lines to reach ±0.05 dex.

The micro-turbulent velocities were determined by adjusting them until the Fe I abundances exhibited no dependence on EW. The uncertainties were estimated by adjusting the ξ_t to force the slope of the linear fit to be 0.001. Metallicities were initially either adopted from previous papers or set to be 0.0 dex. The final results were obtained by iteratively adjusting all the four atmospheric parameters mentioned above until they reached a consistent solution. The uncertainties in [Fe/H] were derived from the scatter of iron abundances measured from different Fe I lines.

Taking HD 119650 and HD 125079 as examples, Fig. 3 shows the excitation equilibrium, ionization equilibrium, and EW-independent of the derived iron abundances from Fe I and Fe II lines. The stellar atmospheric parameters derived from different methods for the 20

<table><caption>**Table 3.** Physical parameters of the sample stars (continued).</caption><thead><tr><th>Star name</th><th><i>M</i>/<i>M</i><sub>☉</sub></th><th>log <i>L</i>/<i>L</i><sub>☉</sub></th></tr></thead><tbody><tr><td>HD 77912</td><td>4.50</td><td>2.79</td></tr><tr><td>HD 85503</td><td>1.30</td><td>1.52</td></tr><tr><td>HD 92482</td><td>1.90</td><td>2.07</td></tr><tr><td>HD 101079</td><td>3.00</td><td>1.80</td></tr><tr><td>HD 107950</td><td>3.00</td><td>2.20</td></tr><tr><td>HD 115927</td><td>0.90</td><td>1.42</td></tr><tr><td>HD 119650</td><td>1.60</td><td>2.16</td></tr><tr><td>HD 125079</td><td>2.20</td><td>1.44</td></tr><tr><td>HD 130255</td><td>0.70</td><td>1.95</td></tr><tr><td>HD 131873</td><td>1.80</td><td>2.72</td></tr><tr><td>HD 133208</td><td>3.50</td><td>2.25</td></tr><tr><td>HD 136138</td><td>3.00</td><td>1.78</td></tr><tr><td>HD 150430</td><td>1.85</td><td>1.83</td></tr><tr><td>HD 151101</td><td>4.50</td><td>2.85</td></tr><tr><td>HD 160507</td><td>2.20</td><td>1.63</td></tr><tr><td>HD 160538</td><td>1.00</td><td>1.12</td></tr><tr><td>HD 168532</td><td>6.00</td><td>2.81</td></tr><tr><td>HD 176524</td><td>2.50</td><td>2.28</td></tr><tr><td>HD 176670</td><td>4.50</td><td>3.23</td></tr><tr><td>HD 177304</td><td>3.50</td><td>2.20</td></tr></tbody></table>

sample stars are provided in Table 2. Furthermore, the stellar masses derived from the $Y^2$ models and the luminosities of our sample stars calculated based on masses are listed in Table 3. Based on the analysis of stellar atmosphere parameters above, we estimated the typical uncertainties of the stellar atmosphere parameters for our sample stars as Δ*T*<sub>eff</sub> = 100 K, Δlog *g* = 0.1 dex, Δ[Fe/H] = 0.1 dex, and Δξ<sub>t</sub> = 0.2 km s<sup>-1</sup>. These uncertainties are later used in the abundance uncertainty estimations (see Section 4.1). From the table, we can see that *T*<sub>eff</sub> of HD 115927 estimated by *V* − *K* is 300 K lower than that derived from the spectroscopic method. Based on the photometric colours, Casagrande et al. (2011) computed the *T*<sub>eff</sub> of this star as 4566 K, which is similar to the value derived from *V* − *K* in this work. By fitting synthetic spectra, Forsberg et al. (2019) computed the *T*<sub>eff</sub> of this star as 4768 K with an uncertainty of 50 K. Using
<table><caption>**Table 4.** Comparisons of atmospheric parameters with literature values.</caption><thead><tr><th>Star name</th><th>T<sub>eff</sub></th><th>log g</th><th>[Fe/H]</th><th>ξ<sub>t</sub></th><th>References</th></tr></thead><tbody><tr><td rowspan="3">HD 77912</td><td>5021</td><td>2.06</td><td>0.05</td><td>2.05</td><td>1</td></tr><tr><td>4899</td><td>1.75</td><td>-0.14</td><td>2.13</td><td>2</td></tr><tr><td>5001</td><td>2.03</td><td>0.12</td><td>2.16</td><td>3</td></tr><tr><td rowspan="2">HD 85503</td><td>4635</td><td>2.65</td><td>0.38</td><td>1.50</td><td>1</td></tr><tr><td>4480</td><td>2.61</td><td>0.17</td><td>2.20</td><td>4</td></tr><tr><td rowspan="2">HD 101079</td><td>5100</td><td>2.90</td><td>0.25</td><td>1.40</td><td>1</td></tr><tr><td>5000</td><td>2.70</td><td>0.10</td><td>1.30</td><td>5</td></tr><tr><td rowspan="2">HD 107950</td><td>5247</td><td>2.55</td><td>0.23</td><td>1.60</td><td>1</td></tr><tr><td>5171</td><td>2.60</td><td>0.01</td><td>1.63</td><td>2</td></tr><tr><td rowspan="2">HD 115927</td><td>4881</td><td>2.68</td><td>-0.21</td><td>1.25</td><td>1</td></tr><tr><td>4768</td><td>2.64</td><td>-0.50</td><td>1.29</td><td>6</td></tr><tr><td rowspan="2">HD 119650</td><td>4510</td><td>2.05</td><td>-0.22</td><td>1.52</td><td>1</td></tr><tr><td>4825</td><td>2.85</td><td>0.04</td><td>1.62</td><td>7</td></tr><tr><td rowspan="2">HD 125079</td><td>4500</td><td>1.60</td><td>-0.10</td><td>1.40</td><td>8</td></tr><tr><td>5334</td><td>3.20</td><td>-0.16</td><td>1.20</td><td>1</td></tr><tr><td rowspan="2">HD 130255</td><td>5520</td><td>3.30</td><td>-0.18</td><td>1.25</td><td>9</td></tr><tr><td>4509</td><td>1.90</td><td>-0.88</td><td>1.30</td><td>1</td></tr><tr><td rowspan="2">HD 131873</td><td>4400</td><td>1.50</td><td>-1.11</td><td>1.30</td><td>8</td></tr><tr><td>4150</td><td>1.40</td><td>-0.08</td><td>1.65</td><td>1</td></tr><tr><td rowspan="2">HD 133208</td><td>4030</td><td>1.83</td><td>-0.29</td><td>2.40</td><td>4</td></tr><tr><td>5174</td><td>2.54</td><td>0.23</td><td>1.70</td><td>1</td></tr><tr><td rowspan="2">HD 136138</td><td>5001</td><td>2.35</td><td>-0.07</td><td>1.61</td><td>2</td></tr><tr><td>5100</td><td>2.92</td><td>0.07</td><td>1.45</td><td>1</td></tr><tr><td rowspan="2">HD 160507</td><td>4995</td><td>2.80</td><td>-0.19</td><td>1.50</td><td>10</td></tr><tr><td>5044</td><td>2.91</td><td>0.09</td><td>1.40</td><td>1</td></tr><tr><td rowspan="2">HD 160538</td><td>5029</td><td>2.96</td><td>-0.04</td><td>1.57</td><td>11</td></tr><tr><td>4812</td><td>3.00</td><td>-0.02</td><td>1.70</td><td>1</td></tr><tr><td rowspan="2">HD 168532</td><td>4851</td><td>3.00</td><td>-0.14</td><td>1.76</td><td>12</td></tr><tr><td>4307</td><td>1.90</td><td>0.33</td><td>1.65</td><td>1</td></tr><tr><td rowspan="2">HD 176524</td><td>4130</td><td>2.05</td><td>0.00</td><td>2.40</td><td>4</td></tr><tr><td>4700</td><td>2.20</td><td>0.13</td><td>1.55</td><td>1</td></tr><tr><td rowspan="2">HD 176670</td><td>4520</td><td>2.55</td><td>-0.12</td><td>2.10</td><td>4</td></tr><tr><td>4300</td><td>1.35</td><td>0.09</td><td>1.80</td><td>1</td></tr><tr><td></td><td>4220</td><td>2.21</td><td>-0.03</td><td>2.80</td><td>4</td></tr></tbody></table>

Note. References: 1. This work; 2. Takeda, Sato & Murata (2008); 3. Luck (2014); 4. McWilliam (1990); 5. Pereira et al. (2011); 6. Forsberg et al. (2019); 7. Mahanta et al. (2016); 8. de Castro et al. (2016); 9. Karinkuzhi & Goswami (2014); 10. Mishenina et al. (2006); 11. Wang et al. (2011); 12. Merle et al. (2016).

the spectroscopic method, we derived the $T_{\text{eff}}$ of this star as 4881 K with an uncertainty of 40 K, which is close to the value of Forsberg et al. (2019) when considering the uncertainties. Additionally, the log g of HD 168532 obtained through the parallax of Gaia DR3 is 0.60 dex smaller than that obtained by the spectroscopic method. The reasons for this discrepancy may include: First, the extinction of the V magnitude of this star is significant, reaching 0.56. Secondly, the relative error in the parallax of this star is large, reaching 5.9 per cent.

For checking the reliability of our derived stellar atmospheric parameters, we present a comparison of our obtained stellar atmospheric parameters with those from the literature, whenever available, in Table 4. The mean differences of the four stellar atmospheric parameters for the common stars between this work and the literature are $\Delta T_{\text{eff}} \lesssim 123$ K, $\Delta \log g \lesssim 0.24$ dex, $\Delta[\text{Fe/H}] \lesssim 0.20$ dex, and $\Delta \xi_t \lesssim 0.28$ km s$^{-1}$. For HD 119650, a significant discrepancy between our values and those of Mahanta et al. (2016) can be seen. Both our study and Mahanta et al. (2016) employed the spectroscopic method to determine these parameters, and the difference in [Fe/H] may be attributed to the use of different Fe lines and atomic data (e.g. log gf). Mahanta et al. (2016) measured the Fe abundance of this

<table><caption>**Table 5.** Absorption lines used for abundance determination.</caption><thead><tr><th>Species</th><th>λ(Å)</th><th>χ(eV)</th><th>log gf</th><th>Ref.</th></tr></thead><tbody><tr><td rowspan="3">Na I</td><td>5682.63</td><td>2.10</td><td>-0.71</td><td>1</td></tr><tr><td>5688.21</td><td>2.10</td><td>-0.45</td><td>1</td></tr><tr><td>6154.23</td><td>2.10</td><td>-1.55</td><td>1</td></tr><tr><td rowspan="5">Mg I</td><td>6160.75</td><td>2.10</td><td>-1.25</td><td>1</td></tr><tr><td>4571.10</td><td>0.00</td><td>-5.47</td><td>2</td></tr><tr><td>4703.00</td><td>4.34</td><td>-0.37</td><td>2</td></tr><tr><td>5528.42</td><td>4.34</td><td>-0.47</td><td>2</td></tr><tr><td>5711.10</td><td>4.34</td><td>-1.75</td><td>2</td></tr><tr><td rowspan="3">Al I</td><td>6799.00</td><td>5.75</td><td>-1.56</td><td>2</td></tr><tr><td>5557.07</td><td>3.14</td><td>-2.23</td><td>3</td></tr><tr><td>6696.03</td><td>3.14</td><td>-1.35</td><td>4</td></tr><tr><td></td><td>6698.67</td><td>3.14</td><td>-1.65</td><td>4</td></tr></tbody></table>

Notes. References: 1. NIST; 2. Zhao et al. (2016); 3. Bensby, Feltzing & Oey (2014); 4. De Smedt et al. (2015); 5. Lawler et al. (2019); 6. Roederer et al. (2018); 7. Den Hartog et al. (2011); 8. O'Brian et al. (1991); 9. Fuhr, Martin & Wiese (1988); 10. Lambert et al. (1996); 11. Bard & Kock (1994); 12. Chen et al. (2000); 13. Nissen & Schuster (1997); 14. Hannaford et al. (1982); 15. Biemont et al. (1981); 16. McWilliam (1998); 17. Mashonkina & Gehren (2000); 18. Lawler et al. (2001).

(The complete version of this table is available in machine-readable form in the online edition of the journal.)

star using 101 Fe I and 7 Fe II lines, while we utilized 142 Fe I and 18 Fe II lines. The $T_{\text{eff}}$ of this star in our work is 315 K lower than that of Mahanta et al. (2016), and the log g is 0.80 dex lower. We notice that the $T_{\text{eff}}$ and log g values of this star determined through colour indices V-K and parallax are 4450 K and 1.83 dex, respectively, which are very close to our results. There is also a significant disparity in the log g of HD 176670 between this study and McWilliam (1990), which may be attributed to our use of the reliable Gaia parallax data, which were not available in 1990 when neither Hipparcos nor Gaia data existed.

# 4 ANALYSIS OF CHEMICAL ABUNDANCES

## 4.1 Abundances and their uncertainties

Based on the EWs and stellar atmospheric parameters derived above, we determined the abundances of light odd-Z elements (Na and Al), α-elements (Mg, Si, Ca, and Ti), Fe-peak elements (Cr, Fe, and Ni), and n-capture elements (Sr, Y, Zr, Ce, Nd, and Sm). Our methodology involved using the model atmospheres, which were interpolated from a grid of plane-parallel, local thermodynamic equilibrium (LTE) models provided by Kurucz (1993). Additionally, we employed the ABONTEST8 program, developed by Dr. P. Magain, to calculate the theoretical line EWs. The final determination of abundances was made by ensuring consistency between the theoretical EWs and observed values. The entire set of abundance-analysis procedures used in this work has been employed in other studies (e.g. Liang et al. 2003; Liu et al. 2009; Yang et al. 2016; Kong et al. 2018a). For the elements Sc, Mn, Ba, La, and Eu, we determined the abundances using the spectral synthesis method. We utilized the IDL code Spectrum Investigation Utility (SIU) developed by Reetz (1999) to fit the observed spectra. The final abundances of the 20 sample stars are listed in Tables 6 and 7. The chemical abundances of the Sun were adopted from Asplund et al. (2009).
<table><caption>**Table 6.** Element abundances of the sample stars.</caption><thead><tr><th rowspan="2">Species</th><th colspan="2">HD 77912</th><th colspan="2">HD 85503</th><th colspan="2">HD 92482</th><th colspan="2">HD 101079</th><th colspan="2">HD 107950</th><th colspan="2">HD 115927</th><th colspan="2">HD 119650</th></tr><tr><th>log ε(X)</th><th>[X/Fe]</th><th>log ε(X)</th><th>[X/Fe]</th><th>log ε(X)</th><th>[X/Fe]</th><th>log ε(X)</th><th>[X/Fe]</th><th>log ε(X)</th><th>[X/Fe]</th><th>log ε(X)</th><th>[X/Fe]</th><th>log ε(X)</th><th>[X/Fe]</th></tr></thead><tbody><tr><td>Fe I</td><td>7.55</td><td>—</td><td>7.88</td><td>—</td><td>7.26</td><td>—</td><td>7.75</td><td>—</td><td>7.73</td><td>—</td><td>7.29</td><td>—</td><td>7.28</td><td>—</td></tr><tr><td>Fe II</td><td>7.55</td><td>—</td><td>7.86</td><td>—</td><td>7.28</td><td>—</td><td>7.75</td><td>—</td><td>7.73</td><td>—</td><td>7.31</td><td>—</td><td>7.28</td><td>—</td></tr><tr><td>Na I</td><td>6.78</td><td>0.49</td><td>—</td><td>—</td><td>6.27</td><td>0.27</td><td>6.75</td><td>0.26</td><td>6.78</td><td>0.31</td><td>6.14</td><td>0.11</td><td>6.25</td><td>0.23</td></tr><tr><td>Mg I</td><td>7.59</td><td>-0.06</td><td>7.92</td><td>-0.06</td><td>7.43</td><td>0.07</td><td>7.83</td><td>-0.02</td><td>7.75</td><td>-0.08</td><td>7.40</td><td>0.01</td><td>7.30</td><td>-0.08</td></tr><tr><td>Al I</td><td>6.57</td><td>0.07</td><td>7.06</td><td>0.23</td><td>6.19</td><td>-0.02</td><td>6.66</td><td>-0.04</td><td>6.60</td><td>-0.08</td><td>6.14</td><td>-0.10</td><td>6.15</td><td>-0.08</td></tr><tr><td>Si I</td><td>7.64</td><td>0.08</td><td>8.17</td><td>0.28</td><td>7.56</td><td>0.29</td><td>7.83</td><td>0.07</td><td>7.78</td><td>0.04</td><td>7.44</td><td>0.14</td><td>7.51</td><td>0.22</td></tr><tr><td>Ca I</td><td>6.41</td><td>0.02</td><td>6.73</td><td>0.01</td><td>6.13</td><td>0.03</td><td>6.55</td><td>-0.04</td><td>6.61</td><td>0.04</td><td>6.13</td><td>0.00</td><td>6.04</td><td>-0.08</td></tr><tr><td>Sc II</td><td>3.31</td><td>0.11</td><td>3.69</td><td>0.16</td><td>3.06</td><td>0.15</td><td>3.56</td><td>0.16</td><td>3.45</td><td>0.07</td><td>3.07</td><td>0.13</td><td>3.11</td><td>0.18</td></tr><tr><td>Ti I</td><td>4.96</td><td>-0.04</td><td>5.43</td><td>0.10</td><td>4.69</td><td>-0.02</td><td>5.13</td><td>-0.07</td><td>5.11</td><td>-0.07</td><td>4.78</td><td>0.04</td><td>4.60</td><td>-0.13</td></tr><tr><td>Ti II</td><td>5.35</td><td>0.35</td><td>5.47</td><td>0.14</td><td>4.95</td><td>0.24</td><td>5.36</td><td>0.16</td><td>5.44</td><td>0.26</td><td>5.09</td><td>0.35</td><td>4.82</td><td>0.09</td></tr><tr><td>Cr I</td><td>5.84</td><td>0.15</td><td>6.08</td><td>0.06</td><td>5.43</td><td>0.03</td><td>5.86</td><td>-0.03</td><td>5.97</td><td>0.10</td><td>5.38</td><td>-0.05</td><td>5.28</td><td>-0.14</td></tr><tr><td>Cr II</td><td>6.16</td><td>0.47</td><td>6.42</td><td>0.40</td><td>5.78</td><td>0.38</td><td>6.09</td><td>0.20</td><td>6.22</td><td>0.35</td><td>5.71</td><td>0.28</td><td>5.59</td><td>0.17</td></tr><tr><td>Mn I</td><td>5.43</td><td>-0.06</td><td>5.80</td><td>-0.02</td><td>5.30</td><td>0.11</td><td>5.73</td><td>0.05</td><td>5.62</td><td>-0.05</td><td>5.25</td><td>0.03</td><td>5.23</td><td>0.02</td></tr><tr><td>Ni I</td><td>6.26</td><td>-0.01</td><td>6.81</td><td>0.21</td><td>6.00</td><td>0.02</td><td>6.50</td><td>0.03</td><td>6.43</td><td>-0.02</td><td>6.01</td><td>0.00</td><td>5.99</td><td>-0.01</td></tr><tr><td>Sr I</td><td>—</td><td>—</td><td>—</td><td>—</td><td>3.46</td><td>0.83</td><td>3.26</td><td>0.14</td><td>2.94</td><td>-0.16</td><td>3.08</td><td>0.42</td><td>—</td><td>—</td></tr><tr><td>Y I</td><td>2.53</td><td>0.27</td><td>2.55</td><td>-0.04</td><td>2.75</td><td>0.78</td><td>2.82</td><td>0.36</td><td>—</td><td>—</td><td>2.48</td><td>0.48</td><td>2.03</td><td>0.04</td></tr><tr><td>Y II</td><td>2.72</td><td>0.46</td><td>3.02</td><td>0.43</td><td>3.23</td><td>1.26</td><td>3.17</td><td>0.71</td><td>2.73</td><td>0.29</td><td>2.84</td><td>0.84</td><td>2.65</td><td>0.66</td></tr><tr><td>Zr I</td><td>2.82</td><td>0.19</td><td>2.91</td><td>-0.05</td><td>3.20</td><td>0.86</td><td>3.15</td><td>0.32</td><td>—</td><td>—</td><td>2.74</td><td>0.37</td><td>2.40</td><td>0.04</td></tr><tr><td>Zr II</td><td>3.15</td><td>0.52</td><td>3.54</td><td>0.58</td><td>3.61</td><td>1.27</td><td>3.42</td><td>0.59</td><td>3.04</td><td>0.23</td><td>3.14</td><td>0.77</td><td>2.98</td><td>0.62</td></tr><tr><td>Ba II</td><td>2.86</td><td>0.63</td><td>2.87</td><td>0.31</td><td>3.46</td><td>1.52</td><td>3.05</td><td>0.62</td><td>2.72</td><td>0.31</td><td>3.11</td><td>1.14</td><td>2.61</td><td>0.65</td></tr><tr><td>Ba II*</td><td>2.79</td><td>0.56</td><td>2.81</td><td>0.25</td><td>3.38</td><td>1.44</td><td>2.91</td><td>0.48</td><td>2.67</td><td>0.26</td><td>3.04</td><td>1.07</td><td>2.54</td><td>0.58</td></tr><tr><td>La II</td><td>1.40</td><td>0.25</td><td>1.83</td><td>0.35</td><td>2.30</td><td>1.44</td><td>1.79</td><td>0.44</td><td>1.45</td><td>0.12</td><td>2.04</td><td>1.15</td><td>1.48</td><td>0.60</td></tr><tr><td>Ce II</td><td>1.92</td><td>0.29</td><td>2.37</td><td>0.41</td><td>2.83</td><td>1.49</td><td>2.36</td><td>0.53</td><td>2.18</td><td>0.37</td><td>2.70</td><td>1.33</td><td>1.69</td><td>0.33</td></tr><tr><td>Nd II</td><td>1.96</td><td>0.49</td><td>2.26</td><td>0.46</td><td>2.60</td><td>1.42</td><td>2.37</td><td>0.70</td><td>2.17</td><td>0.52</td><td>2.57</td><td>1.36</td><td>1.84</td><td>0.64</td></tr><tr><td>Sm II</td><td>1.52</td><td>0.51</td><td>1.88</td><td>0.54</td><td>2.00</td><td>1.28</td><td>1.57</td><td>0.36</td><td>1.43</td><td>0.24</td><td>2.04</td><td>1.29</td><td>1.09</td><td>0.35</td></tr><tr><td>Eu II</td><td>0.72</td><td>0.15</td><td>1.00</td><td>0.10</td><td>0.85</td><td>0.57</td><td>0.80</td><td>0.03</td><td>0.77</td><td>0.02</td><td>0.63</td><td>0.32</td><td>0.60</td><td>0.30</td></tr></tbody></table><table><thead><tr><th colspan="2">HD 125079</th><th colspan="2">HD 130255</th><th colspan="2">HD 131873</th><th colspan="2">HD 133208</th><th colspan="2">HD 136138</th><th colspan="2">HD 150430</th><th colspan="2">HD 151101</th></tr><tr><th>Species</th><th>log ε(X)</th><th>[X/Fe]</th><th>log ε(X)</th><th>[X/Fe]</th><th>log ε(X)</th><th>[X/Fe]</th><th>log ε(X)</th><th>[X/Fe]</th><th>log ε(X)</th><th>[X/Fe]</th><th>log ε(X)</th><th>[X/Fe]</th><th>log ε(X)</th></tr></thead><tbody><tr><td>Fe I</td><td>7.34</td><td>—</td><td>6.64</td><td>—</td><td>7.42</td><td>—</td><td>7.73</td><td>—</td><td>7.57</td><td>—</td><td>7.24</td><td>—</td><td>7.63</td><td>—</td></tr><tr><td>Fe II</td><td>7.34</td><td>—</td><td>6.62</td><td>—</td><td>7.41</td><td>—</td><td>7.73</td><td>—</td><td>7.60</td><td>—</td><td>7.29</td><td>—</td><td>7.66</td><td>—</td></tr><tr><td>Na I</td><td>6.20</td><td>0.12</td><td>5.41</td><td>0.03</td><td>6.46</td><td>0.30</td><td>6.77</td><td>0.30</td><td>6.47</td><td>0.16</td><td>6.22</td><td>0.24</td><td>6.66</td><td>0.29</td></tr><tr><td>Mg I</td><td>7.44</td><td>0.00</td><td>6.92</td><td>0.18</td><td>7.52</td><td>0.00</td><td>—</td><td>—</td><td>7.62</td><td>-0.05</td><td>7.32</td><td>-0.02</td><td>7.77</td><td>0.04</td></tr><tr><td>Al I</td><td>6.07</td><td>-0.22</td><td>5.83</td><td>0.24</td><td>6.49</td><td>0.12</td><td>6.68</td><td>0.00</td><td>6.60</td><td>0.08</td><td>6.18</td><td>-0.01</td><td>6.54</td><td>-0.04</td></tr><tr><td>Si I</td><td>7.42</td><td>0.07</td><td>7.18</td><td>0.53</td><td>7.52</td><td>0.09</td><td>7.84</td><td>0.10</td><td>7.68</td><td>0.10</td><td>7.43</td><td>0.18</td><td>7.75</td><td>0.11</td></tr><tr><td>Ca I</td><td>6.23</td><td>0.05</td><td>5.70</td><td>0.22</td><td>6.22</td><td>-0.04</td><td>6.61</td><td>0.04</td><td>6.42</td><td>0.01</td><td>5.97</td><td>-0.11</td><td>6.37</td><td>-0.10</td></tr><tr><td>Sc II</td><td>3.18</td><td>0.19</td><td>2.50</td><td>0.21</td><td>3.13</td><td>0.06</td><td>3.51</td><td>0.13</td><td>3.32</td><td>0.10</td><td>3.17</td><td>0.28</td><td>3.43</td><td>0.15</td></tr><tr><td>Ti I</td><td>4.76</td><td>-0.03</td><td>4.31</td><td>0.22</td><td>5.06</td><td>0.19</td><td>5.14</td><td>-0.04</td><td>5.01</td><td>-0.01</td><td>4.55</td><td>-0.14</td><td>4.98</td><td>-0.10</td></tr><tr><td>Ti II</td><td>4.91</td><td>0.12</td><td>4.49</td><td>0.40</td><td>5.11</td><td>0.24</td><td>5.35</td><td>0.17</td><td>5.24</td><td>0.22</td><td>4.97</td><td>0.28</td><td>5.42</td><td>0.34</td></tr><tr><td>Cr I</td><td>5.51</td><td>0.03</td><td>4.78</td><td>0.00</td><td>5.50</td><td>-0.06</td><td>5.92</td><td>0.05</td><td>5.73</td><td>0.02</td><td>5.27</td><td>-0.11</td><td>5.84</td><td>0.07</td></tr><tr><td>Cr II</td><td>5.69</td><td>0.21</td><td>5.14</td><td>0.36</td><td>5.67</td><td>0.11</td><td>6.33</td><td>0.46</td><td>5.89</td><td>0.18</td><td>5.83</td><td>0.45</td><td>6.30</td><td>0.53</td></tr><tr><td>Mn I</td><td>5.24</td><td>-0.03</td><td>4.28</td><td>-0.29</td><td>5.53</td><td>0.18</td><td>5.73</td><td>0.07</td><td>5.49</td><td>-0.01</td><td>5.09</td><td>-0.08</td><td>5.57</td><td>0.01</td></tr><tr><td>Ni I</td><td>6.05</td><td>-0.01</td><td>5.42</td><td>0.06</td><td>6.22</td><td>0.08</td><td>6.45</td><td>0.00</td><td>6.28</td><td>-0.01</td><td>5.95</td><td>-0.01</td><td>6.42</td><td>0.07</td></tr><tr><td>Sr I</td><td>3.54</td><td>0.83</td><td>—</td><td>—</td><td>—</td><td>—</td><td>—</td><td>—</td><td>3.43</td><td>0.49</td><td>—</td><td>—</td><td>—</td><td>—</td></tr><tr><td>Y I</td><td>2.94</td><td>0.89</td><td>—</td><td>—</td><td>2.12</td><td>-0.01</td><td>—</td><td>—</td><td>2.90</td><td>0.62</td><td>2.02</td><td>0.07</td><td>—</td><td>—</td></tr><tr><td>Y II</td><td>3.02</td><td>0.97</td><td>1.79</td><td>0.44</td><td>2.35</td><td>0.22</td><td>2.66</td><td>0.22</td><td>3.04</td><td>0.76</td><td>2.37</td><td>0.42</td><td>2.42</td><td>0.08</td></tr><tr><td>Zr I</td><td>3.32</td><td>0.90</td><td>1.91</td><td>0.19</td><td>2.46</td><td>-0.04</td><td>3.28</td><td>0.47</td><td>3.18</td><td>0.53</td><td>2.40</td><td>0.08</td><td>2.38</td><td>-0.33</td></tr><tr><td>Zr II</td><td>3.51</td><td>1.09</td><td>2.23</td><td>0.51</td><td>—</td><td>—</td><td>3.26</td><td>0.45</td><td>3.53</td><td>0.88</td><td>2.82</td><td>0.50</td><td>—</td><td>—</td></tr><tr><td>Ba II</td><td>3.33</td><td>1.31</td><td>1.96</td><td>0.64</td><td>2.52</td><td>0.42</td><td>2.83</td><td>0.42</td><td>2.84</td><td>0.59</td><td>2.39</td><td>0.47</td><td>2.78</td><td>0.47</td></tr><tr><td>Ba II*</td><td>3.17</td><td>1.15</td><td>1.90</td><td>0.58</td><td>2.47</td><td>0.37</td><td>2.74</td><td>0.33</td><td>2.78</td><td>0.53</td><td>2.33</td><td>0.41</td><td>2.67</td><td>0.36</td></tr><tr><td>La II</td><td>1.97</td><td>1.03</td><td>0.63</td><td>0.39</td><td>1.32</td><td>0.30</td><td>1.56</td><td>0.23</td><td>1.66</td><td>0.49</td><td>1.24</td><td>0.40</td><td>1.32</td><td>0.09</td></tr><tr><td>Ce II</td><td>2.48</td><td>1.06</td><td>1.34</td><td>0.62</td><td>1.77</td><td>0.27</td><td>2.23</td><td>0.42</td><td>2.02</td><td>0.37</td><td>1.79</td><td>0.47</td><td>2.06</td><td>0.35</td></tr><tr><td>Nd II</td><td>2.28</td><td>1.02</td><td>1.29</td><td>0.73</td><td>1.99</td><td>0.65</td><td>2.23</td><td>0.58</td><td>2.20</td><td>0.71</td><td>1.65</td><td>0.49</td><td>2.31</td><td>0.76</td></tr><tr><td>Sm II</td><td>1.41</td><td>0.61</td><td>0.71</td><td>0.61</td><td>—</td><td>—</td><td>1.55</td><td>0.36</td><td>1.51</td><td>0.48</td><td>1.09</td><td>0.39</td><td>1.66</td><td>0.57</td></tr><tr><td>Eu II</td><td>0.70</td><td>0.34</td><td>-0.10</td><td>0.24</td><td>0.54</td><td>0.10</td><td>0.83</td><td>0.08</td><td>0.76</td><td>0.17</td><td>0.58</td><td>0.32</td><td>0.67</td><td>0.02</td></tr></tbody></table>

Note. *The NLTE-corrected Ba abundances are listed here.

The atomic data for most of the observed lines are adopted from the NIST<fnref n="3" /> data base (Kramida et al. 2023), as well as from Bensby et al. (2014) and De Smedt et al. (2015). Additionally, we supplemented

our data for the 20 elements with information from other sources. Detailed information is provided in Table 5. Furthermore, we considered the hyperfine structure (HFS) that splits the atomic terms into hyperfine components and desaturates the spectral lines, especially for the strong ones. Taking the mild Ba-star HD 119650 and the strong Ba-star HD 125079 as examples, the EWs of the Ba lines

<footnote n="3">³http://physics.nist.gov/asd</footnote>
<table><caption>**Table 7.** Element abundances of the sample stars (continued).</caption><thead><tr><th rowspan="2">Species</th><th colspan="2">HD 160507</th><th colspan="2">HD 160538</th><th colspan="2">HD 168532</th><th colspan="2">HD 176524</th><th colspan="2">HD 176670</th><th colspan="2">HD 177304</th></tr><tr><th>log ε (X)</th><th>[X/Fe]</th><th>log ε (X)</th><th>[X/Fe]</th><th>log ε (X)</th><th>[X/Fe]</th><th>log ε (X)</th><th>[X/Fe]</th><th>log ε (X)</th><th>[X/Fe]</th><th>log ε (X)</th><th>[X/Fe]</th></tr></thead><tbody><tr><td>Fe I</td><td>7.59</td><td>—</td><td>7.48</td><td>—</td><td>7.83</td><td>—</td><td>7.63</td><td>—</td><td>7.59</td><td>—</td><td>7.73</td><td>—</td></tr><tr><td>Fe II</td><td>7.63</td><td>—</td><td>7.45</td><td>—</td><td>7.84</td><td>—</td><td>7.65</td><td>—</td><td>7.56</td><td>—</td><td>7.74</td><td>—</td></tr><tr><td>Na I</td><td>6.53</td><td>0.20</td><td>6.55</td><td>0.33</td><td>6.84</td><td>0.27</td><td>6.58</td><td>0.21</td><td>6.85</td><td>0.52</td><td>6.77</td><td>0.30</td></tr><tr><td>Mg I</td><td>7.62</td><td>-0.07</td><td>7.44</td><td>-0.14</td><td>7.89</td><td>-0.04</td><td>7.57</td><td>-0.16</td><td>7.84</td><td>0.15</td><td>7.68</td><td>-0.15</td></tr><tr><td>Al I</td><td>6.51</td><td>-0.03</td><td>6.62</td><td>0.19</td><td>6.78</td><td>0.00</td><td>6.54</td><td>-0.04</td><td>6.61</td><td>0.07</td><td>6.52</td><td>-0.16</td></tr><tr><td>Si I</td><td>7.77</td><td>0.17</td><td>7.48</td><td>-0.01</td><td>7.90</td><td>0.06</td><td>7.71</td><td>0.07</td><td>7.79</td><td>0.19</td><td>7.73</td><td>-0.01</td></tr><tr><td>Ca I</td><td>6.42</td><td>-0.01</td><td>6.49</td><td>0.17</td><td>6.61</td><td>-0.06</td><td>6.49</td><td>0.02</td><td>6.41</td><td>-0.02</td><td>6.52</td><td>-0.05</td></tr><tr><td>Sc II</td><td>3.43</td><td>0.19</td><td>3.28</td><td>0.15</td><td>3.64</td><td>0.16</td><td>3.26</td><td>-0.02</td><td>3.36</td><td>0.12</td><td>3.49</td><td>0.11</td></tr><tr><td>Ti I</td><td>4.98</td><td>-0.06</td><td>5.01</td><td>0.08</td><td>5.43</td><td>0.15</td><td>5.11</td><td>0.03</td><td>5.20</td><td>0.16</td><td>5.14</td><td>-0.04</td></tr><tr><td>Ti II</td><td>5.24</td><td>0.20</td><td>5.25</td><td>0.32</td><td>5.49</td><td>0.21</td><td>5.17</td><td>0.09</td><td>5.29</td><td>0.25</td><td>5.41</td><td>0.23</td></tr><tr><td>Cr I</td><td>5.68</td><td>-0.05</td><td>5.72</td><td>0.10</td><td>5.98</td><td>0.01</td><td>5.77</td><td>0.00</td><td>5.75</td><td>0.02</td><td>5.85</td><td>-0.02</td></tr><tr><td>Cr II</td><td>6.10</td><td>0.37</td><td>6.11</td><td>0.49</td><td>—</td><td>—</td><td>5.90</td><td>0.13</td><td>6.18</td><td>0.45</td><td>6.02</td><td>0.15</td></tr><tr><td>Mn I</td><td>5.48</td><td>-0.04</td><td>5.45</td><td>0.04</td><td>5.92</td><td>0.16</td><td>5.57</td><td>0.01</td><td>5.66</td><td>0.14</td><td>5.64</td><td>-0.03</td></tr><tr><td>Ni I</td><td>6.35</td><td>0.04</td><td>6.20</td><td>0.00</td><td>6.77</td><td>0.22</td><td>6.35</td><td>0.00</td><td>6.47</td><td>0.16</td><td>6.36</td><td>-0.09</td></tr><tr><td>Sr I</td><td>3.18</td><td>0.22</td><td>—</td><td>—</td><td>—</td><td>—</td><td>—</td><td>—</td><td>—</td><td>—</td><td>3.13</td><td>0.03</td></tr><tr><td>Y I</td><td>2.77</td><td>0.47</td><td>2.44</td><td>0.25</td><td>2.70</td><td>0.16</td><td>—</td><td>—</td><td>2.33</td><td>0.03</td><td>2.64</td><td>0.20</td></tr><tr><td>Y II</td><td>3.18</td><td>0.88</td><td>—</td><td>—</td><td>3.07</td><td>0.53</td><td>2.42</td><td>0.08</td><td>2.61</td><td>0.31</td><td>2.60</td><td>0.16</td></tr><tr><td>Zr I</td><td>3.07</td><td>0.40</td><td>2.72</td><td>0.16</td><td>3.07</td><td>0.16</td><td>2.75</td><td>0.04</td><td>2.67</td><td>0.00</td><td>3.07</td><td>0.26</td></tr><tr><td>Zr II</td><td>3.41</td><td>0.74</td><td>—</td><td>—</td><td>—</td><td>—</td><td>2.76</td><td>0.05</td><td>3.25</td><td>0.58</td><td>3.32</td><td>0.51</td></tr><tr><td>Ba II</td><td>3.12</td><td>0.85</td><td>2.55</td><td>0.39</td><td>3.05</td><td>0.54</td><td>2.70</td><td>0.39</td><td>2.77</td><td>0.50</td><td>2.90</td><td>0.49</td></tr><tr><td>Ba II*</td><td>3.02</td><td>0.75</td><td>2.41</td><td>0.25</td><td>2.99</td><td>0.48</td><td>2.50</td><td>0.19</td><td>2.67</td><td>0.40</td><td>2.81</td><td>0.40</td></tr><tr><td>La II</td><td>1.95</td><td>0.76</td><td>1.22</td><td>0.14</td><td>1.80</td><td>0.37</td><td>1.44</td><td>0.21</td><td>1.38</td><td>0.19</td><td>1.54</td><td>0.21</td></tr><tr><td>Ce II</td><td>2.37</td><td>0.70</td><td>1.89</td><td>0.33</td><td>2.29</td><td>0.38</td><td>2.04</td><td>0.33</td><td>2.22</td><td>0.55</td><td>2.20</td><td>0.39</td></tr><tr><td>Nd II</td><td>2.54</td><td>1.03</td><td>1.99</td><td>0.59</td><td>2.23</td><td>0.48</td><td>1.93</td><td>0.38</td><td>2.07</td><td>0.56</td><td>2.02</td><td>0.37</td></tr><tr><td>Sm II</td><td>1.58</td><td>0.53</td><td>1.31</td><td>0.37</td><td>1.63</td><td>0.34</td><td>1.40</td><td>0.31</td><td>2.06</td><td>1.01</td><td>1.46</td><td>0.27</td></tr><tr><td>Eu II</td><td>0.83</td><td>0.22</td><td>0.55</td><td>0.05</td><td>1.00</td><td>0.15</td><td>0.73</td><td>0.08</td><td>0.70</td><td>0.09</td><td>0.83</td><td>0.08</td></tr></tbody></table>

*Note.* The NLTE-corrected Ba abundances are listed here.

at 5853.67, 6141.71, and 6496.91 Å in HD 119650 were measured as 154.6, 229.1, and 233.0 mÅ, respectively. For HD 125079, these values are 149.4, 255.6, and 235.5 mÅ. We accounted for the HFS of Sc (Lawler et al. 2019), Mn (Den Hartog et al. 2011), Ba (McWilliam 1998), La (Ivans et al. 2006), and Eu (Lawler et al. 2001) in this work. Additionally, considering that the non-local thermodynamic equilibrium (NLTE) effects are significant for Ba abundances (Mashonkina, Gehren & Bikmaev 1999; Reddy & Lambert 2017; Liu et al. 2021), the NLTE effects have been considered in determining the Ba abundances. Fig. 4 shows the fitting for the three Ba II lines of HD 92482.

In our study, we have determined detailed elemental abundances for four stars (i.e. HD 92482, HD 150430, HD 151101, and HD 177304) for the first time. As for the remaining 16 sample stars, four of them (HD 101079, HD 115927, HD 130255, and HD 176670) have had their chemical abundances analysed by Pereira et al. (2011), Forsberg et al. (2019), de Castro et al. (2016), and McWilliam (1990), respectively, but none provided the Ba abundances. Fernández-Villacañas, Rego & Cornide (1990) analysed Ba and Fe abundances for two stars (HD 77912 and HD 133208) using intermediate dispersion spectra. Hence, the Ba abundances for these six stars have been derived from high-resolution spectra for the first time. While the Ba abundances of the remaining 10 sample stars in our study have been explored before (e.g. Začs et al. 1997; Luck 2015; Kong et al. 2018b; Casamiquela, Tarricq & Soubiran 2020), most researchers used EW measurements to obtain the abundance results. In this work, we employed the spectral synthesis method to calculate the Ba abundance taking into account both the HFS and NLTE effects. Using HD 119650 and HD 125079 as examples, the standard deviations of Ba abundances

from line to line are approximately 0.05 and 0.06 dex, respectively. The NLTE corrections $\Delta_{\text{NLTE}}$ (= log $\epsilon$ (Ba)$_{\text{NLTE}}$ - log $\epsilon$ (Ba)$_{\text{LTE}}$) for different lines vary. The largest absolute values of the corrections for the two stars both correspond to the line at 6496.91 Å, reaching $\Delta_{\text{NLTE}}$ = -0.16 and -0.33 dex, respectively. For comparison, the LTE and NLTE Ba abundances of the sample stars are presented in Tables 6 and 7. Our results demonstrate that NLTE analyses are critical for enhancing the accuracy of stellar abundance measurements.

The uncertainties on the abundance determination arise from two main sources: stochastic errors (random errors on the spectral-line measurements, uncertainties on the log $g f$) and systematic errors primarily due to uncertainties in the determination of the stellar atmospheric parameters. For elements such as Na, Mg Al, Si, Ca, Ti, Cr, Fe, Ni, Sr, Y, Zr, Ce, Nd, and Sm, the uncertainty on the EWs has been estimated to be about 3.5 mÅ (see Fig. 2). The magnitude of the random error on the abundance of an element X is estimated as $\sigma_{\text{ran}} = \sigma/\sqrt{N-1}$, where $\sigma$ is the root mean square (rms) around the mean X abundance and N is the number of lines of the X element (with $N \ge 2$). For elements Sc, Mn, Ba, La, and Eu, whose abundances were derived using the spectral synthesis method, the abundance scatter from multiple lines of one element could represent the random uncertainty (Tautvaišienė et al. 2021). If the abundance of an element in a sample star is determined using only one line, we do not provide the random error. Systematic uncertainties are dominated by the errors in the choice of stellar parameters. They were estimated by varying $T_{\text{eff}}$ by 100 K, log $g$ by 0.1 dex, [Fe/H] by 0.1 dex, and $\xi_i$ by 0.2 km s$^{-1}$. The total uncertainty on the abundance of the element X, log $\epsilon$ (X), is computed as the quadratic sum of the stochastic and systematic errors. The results are presented in Table 8 for two sample
<figure><img src="image_7.png" /><figcaption>**Figure 4.** Examples of spectral fittings of Ba lines at 5853.67, 6141.71, and 6496.91 Å for HD 92482. The observed and synthetic spectra are shown with filled circles and black lines, respectively. The blue and red lines are the synthetic spectra for $\Delta \log \epsilon(\text{Ba}) = \pm 0.2$ dex around the best fits.</figcaption></figure>

stars, and it can be found that the error is generally on the order of 0.15 dex or less.

## 4.2 Comparisons of abundance results with previous literature

To assess the reliability of the elemental abundances derived in this work, we compared our results with those from the literature for the common stars. To minimize uncertainties arising from different stellar atmospheric parameters adopted, we only compared the elemental abundances between previous works and our results for stars with discrepancies satisfying $\Delta\Gamma_{\text{eff}} < 200$ K, $\Delta \log g < 0.20$ dex, $\Delta[\text{Fe/H}] < 0.20$ dex, and $\Delta\xi_t < 0.40 \text{ km s}^{-1}$. The comparison results are listed in Table 9. Most of the elements exhibit mean deviations of $\lesssim 0.2$ dex between this work and previous studies. However, our Mg abundances for the three common stars are noticeably lower compared to previous research. The primary reason for this difference may lie in the usage of different spectral lines for determining Mg

abundances. We primarily utilize the spectral lines at 4571 and 5711 Å to determine Mg abundances in our sample stars. Notably, Luck (2014) did not provide the atomic lines used for abundance determination. Pereira et al. (2011) used eight lines to determine the Mg abundance of HD 101079. There is only one common line (5711 Å) between Pereira et al. (2011) and our study. For this line, we adopted the excitation potential $\chi = 4.34$ ev and oscillator strength $\log gf = -1.75$, and derived the equivalent width of EW = 135 mÅ, which closely match the values reported in Pereira et al. (2011) ($\chi = 4.34$ ev, $\log gf = -1.75$, and EW = 132 mÅ). Karinkuzhi & Goswami (2014) utilized two lines to ascertain the Mg abundance of HD 125079. Unfortunately, there is no common line between Karinkuzhi & Goswami (2014) and our analysis. There is also a significant deviation in the Sr abundance (~0.76 dex) for HD 125079 between this work and Karinkuzhi & Goswami (2014). The stellar atmosphere parameters for this star in our work are similar to those in Karinkuzhi & Goswami (2014). The Sr abundance is determined using only the Sr I line at $\lambda = 4607$ Å in both works. However, the corresponding $\log gf$ in our work ($\log gf = 0.28$, Liu et al. (2021)) is 0.62 larger than that in Karinkuzhi & Goswami (2014) ($\log gf = -0.57$). Therefore, the deviation in the Sr abundance for HD 125079 between the two works can be attributed to the different $\log gf$ values adopted.

## 5 DISCUSSION

In this study, we conducted a chemical abundance analysis of 20 Ba-stars spanning metallicities in the range of -0.86 to 0.38 dex. A low metallicity star, HD 130255 ([Fe/H] = -0.86 dex), is included in this work. The abundance analysis of these Ba-stars with [Fe/H] in a large range can help us better study the stellar formation and chemical evolution of Ba-stars.

### 5.1 The Ba classifications of HD 115927 and HD 160538

The [s/Fe] ratio serves as an important indicator for identifying whether a star is s-rich. This ratio is defined as the mean abundance of the s-process elements, including Sr, Y, Zr, Ba, La, Ce, and Nd, relative to Fe abundance (see for example fig. 16 in de Castro et al. 2016). A star is classified as s-rich if its [s/Fe] ratio is significantly higher than that in the other stars. Various studies have proposed different thresholds for this classification. For instance, Rojas et al. (2013) adopted [s/Fe] = 0.34 dex, while de Castro et al. (2016) suggested 0.25 dex. In Fig. 5, we show [s/Fe] as a function of [Fe/H] for the sample stars. The typical error is calculated based on the average [s/Fe] uncertainty of the sample stars HD 119650 and HD 125079. It can be seen from the figure that the s-process abundances of our sample stars fall within the range of $0.21 \le [\text{s/Fe}] \le 1.31$ dex. The [s/Fe] ratios of HD 107950 (0.23 dex) and HD 176524 (0.21 dex) are slightly lower than 0.25 dex. However, taking into account, the uncertainties on the abundances, these two sample stars could be considered as mild (or marginal) Ba-stars.

In our sample, there are two stars, i.e. HD 115927 and HD 160538, having no Ba intensity (Warner 1965; Lu 1991). Forsberg et al. (2019) derived high abundance ratios for HD 115927 with [Zr/Fe] = 0.68 dex, [La/Fe] = 1.00 dex, and [Ce/Fe] = 1.10 dex. Similarly, based on the abundances of Sr, Y, Zr, Ba, La, and Ce reported by Merle et al. (2016), the [s/Fe] of HD 160538 is 0.30 dex. However, these authors did not evidence the Ba-rich nature of HD 115927 and HD 160538. Accordingly, these two stars can be categorized as strong and mild Ba-stars, respectively. We further determined the
<table><caption>**Table 8.** The uncertainties of the abundance analysis for the sample stars HD 119650 and HD 125079.</caption><thead><tr><th rowspan="3">Species</th><th rowspan="3">σ/√N − 1</th><th colspan="5">HD 119650</th><th colspan="5">HD 125079</th></tr><tr><th>ΔT<sub>eff</sub></th><th>Δlog g</th><th>Δ[Fe/H]</th><th>Δξ<sub>i</sub></th><th>σ<sub>total</sub></th><th>σ/√N − 1</th><th>ΔT<sub>eff</sub></th><th>Δlog g</th><th>Δ[Fe/H]</th><th>Δξ<sub>i</sub></th><th>σ<sub>total</sub></th></tr><tr><th>(+100 K)</th><th>(+0.1 dex)</th><th>(+0.1 dex)</th><th>(+0.2 km s<sup>−1</sup>)</th><th></th><th>(+100 K)</th><th>(+0.1 dex)</th><th>(+0.1 dex)</th><th>(+0.2 km s<sup>−1</sup>)</th><th></th><th></th></tr></thead><tbody><tr><td>Fe I</td><td>0.01</td><td>0.04</td><td>0.01</td><td>0.01</td><td>0.09</td><td>0.10</td><td>0.01</td><td>0.08</td><td>0.01</td><td>0.08</td><td>0.00</td><td>0.11</td></tr><tr><td>Fe II</td><td>0.02</td><td>0.11</td><td>0.05</td><td>0.03</td><td>0.06</td><td>0.14</td><td>0.01</td><td>0.05</td><td>0.02</td><td>0.09</td><td>0.01</td><td>0.11</td></tr><tr><td>Na I</td><td>0.03</td><td>0.08</td><td>0.01</td><td>0.02</td><td>0.07</td><td>0.11</td><td>0.03</td><td>0.07</td><td>0.01</td><td>0.03</td><td>0.00</td><td>0.08</td></tr><tr><td>Mg I</td><td>–</td><td>0.08</td><td>0.02</td><td>0.02</td><td>0.05</td><td>0.10</td><td>0.01</td><td>0.09</td><td>0.03</td><td>0.05</td><td>0.02</td><td>0.11</td></tr><tr><td>Al I</td><td>0.04</td><td>0.08</td><td>0.00</td><td>0.00</td><td>0.02</td><td>0.09</td><td>–</td><td>0.06</td><td>0.00</td><td>0.00</td><td>0.00</td><td>0.06</td></tr><tr><td>Si I</td><td>0.01</td><td>0.04</td><td>0.02</td><td>0.02</td><td>0.04</td><td>0.06</td><td>0.01</td><td>0.03</td><td>0.01</td><td>0.02</td><td>0.01</td><td>0.04</td></tr><tr><td>Ca I</td><td>0.01</td><td>0.11</td><td>0.01</td><td>0.01</td><td>0.10</td><td>0.15</td><td>0.01</td><td>0.09</td><td>0.03</td><td>0.06</td><td>0.01</td><td>0.11</td></tr><tr><td>Sc II</td><td>0.04</td><td>0.05</td><td>0.05</td><td>0.04</td><td>0.06</td><td>0.11</td><td>0.04</td><td>0.04</td><td>0.06</td><td>0.09</td><td>0.09</td><td>0.15</td></tr><tr><td>Ti I</td><td>0.01</td><td>0.15</td><td>0.00</td><td>0.02</td><td>0.12</td><td>0.19</td><td>0.02</td><td>0.12</td><td>0.01</td><td>0.08</td><td>0.00</td><td>0.15</td></tr><tr><td>Ti II</td><td>0.02</td><td>0.01</td><td>0.04</td><td>0.04</td><td>0.09</td><td>0.11</td><td>0.01</td><td>0.01</td><td>0.03</td><td>0.07</td><td>0.03</td><td>0.08</td></tr><tr><td>Cr I</td><td>0.02</td><td>0.14</td><td>0.01</td><td>0.02</td><td>0.13</td><td>0.19</td><td>0.02</td><td>0.11</td><td>0.02</td><td>0.10</td><td>0.01</td><td>0.15</td></tr><tr><td>Cr II</td><td>0.03</td><td>0.07</td><td>0.05</td><td>0.03</td><td>0.07</td><td>0.12</td><td>0.02</td><td>0.04</td><td>0.04</td><td>0.07</td><td>0.02</td><td>0.10</td></tr><tr><td>Mn I</td><td>0.01</td><td>0.04</td><td>0.05</td><td>0.04</td><td>0.06</td><td>0.10</td><td>0.03</td><td>0.06</td><td>0.01</td><td>0.08</td><td>0.06</td><td>0.12</td></tr><tr><td>Ni I</td><td>0.01</td><td>0.03</td><td>0.02</td><td>0.02</td><td>0.10</td><td>0.11</td><td>0.01</td><td>0.08</td><td>0.00</td><td>0.06</td><td>0.01</td><td>0.10</td></tr><tr><td>Sr I</td><td>–</td><td>–</td><td>–</td><td>–</td><td>–</td><td>–</td><td>–</td><td>0.13</td><td>0.01</td><td>0.17</td><td>0.02</td><td>0.22</td></tr><tr><td>Y I</td><td>–</td><td>0.08</td><td>0.02</td><td>0.06</td><td>0.04</td><td>0.11</td><td>–</td><td>0.14</td><td>0.00</td><td>0.01</td><td>0.01</td><td>0.14</td></tr><tr><td>Y II</td><td>0.02</td><td>0.00</td><td>0.04</td><td>0.03</td><td>0.10</td><td>0.11</td><td>0.02</td><td>0.01</td><td>0.03</td><td>0.10</td><td>0.04</td><td>0.11</td></tr><tr><td>Zr I</td><td>0.03</td><td>0.13</td><td>0.01</td><td>0.04</td><td>0.04</td><td>0.15</td><td>–</td><td>0.15</td><td>0.00</td><td>0.01</td><td>0.01</td><td>0.15</td></tr><tr><td>Zr II</td><td>0.03</td><td>0.01</td><td>0.04</td><td>0.03</td><td>0.06</td><td>0.09</td><td>0.06</td><td>0.01</td><td>0.05</td><td>0.09</td><td>0.03</td><td>0.13</td></tr><tr><td>Ba II</td><td>0.05</td><td>0.02</td><td>0.01</td><td>0.04</td><td>0.11</td><td>0.13</td><td>0.06</td><td>0.05</td><td>0.05</td><td>0.06</td><td>0.07</td><td>0.13</td></tr><tr><td>La II</td><td>0.00</td><td>0.02</td><td>0.04</td><td>0.03</td><td>0.06</td><td>0.08</td><td>0.07</td><td>0.02</td><td>0.05</td><td>0.06</td><td>0.04</td><td>0.11</td></tr><tr><td>Ce II</td><td>0.03</td><td>0.02</td><td>0.04</td><td>0.04</td><td>0.09</td><td>0.11</td><td>0.04</td><td>0.02</td><td>0.03</td><td>0.12</td><td>0.03</td><td>0.13</td></tr><tr><td>Nd II</td><td>0.02</td><td>0.03</td><td>0.04</td><td>0.04</td><td>0.10</td><td>0.12</td><td>0.03</td><td>0.03</td><td>0.04</td><td>0.09</td><td>0.04</td><td>0.11</td></tr><tr><td>Sm II</td><td>0.03</td><td>0.03</td><td>0.04</td><td>0.04</td><td>0.06</td><td>0.09</td><td>0.04</td><td>0.02</td><td>0.04</td><td>0.05</td><td>0.03</td><td>0.08</td></tr><tr><td>Eu II</td><td>0.03</td><td>0.01</td><td>0.04</td><td>0.04</td><td>0.02</td><td>0.07</td><td>0.02</td><td>0.00</td><td>0.05</td><td>0.04</td><td>0.02</td><td>0.07</td></tr></tbody></table>

<table><caption>**Table 9.** Comparisons of element abundances with literature values.</caption><thead><tr><th>Star name</th><th>[Na I/Fe]</th><th>[Mg I/Fe]</th><th>[Al I/Fe]</th><th>[Si I/Fe]</th><th>[Ca I/Fe]</th><th>[Sc II/Fe]</th><th>[Ti I/Fe]</th><th>[Ti II/Fe]</th><th>[Cr I/Fe]</th><th>[Cr II/Fe]</th><th>[Mn I/Fe]</th><th>[Ni I/Fe]</th><th>References</th></tr></thead><tbody><tr><td>HD 77912</td><td>0.49</td><td>-0.06</td><td>0.07</td><td>0.08</td><td>0.02</td><td>0.11</td><td>-0.04</td><td>0.35</td><td>0.15</td><td>0.47</td><td>-0.06</td><td>-0.01</td><td>1</td></tr><tr><td></td><td>0.57</td><td>0.20</td><td>0.24</td><td>0.28</td><td>0.17</td><td>0.08</td><td>0.15</td><td>–</td><td>0.24</td><td>–</td><td>0.27</td><td>0.11</td><td>2</td></tr><tr><td>HD 101079</td><td>0.26</td><td>-0.02</td><td>-0.04</td><td>0.07</td><td>-0.04</td><td>0.16</td><td>-0.07</td><td>0.16</td><td>-0.03</td><td>0.20</td><td>0.05</td><td>0.03</td><td>1</td></tr><tr><td></td><td>0.20</td><td>0.14</td><td>0.23</td><td>0.19</td><td>0.10</td><td>–</td><td>-0.06</td><td>–</td><td>0.02</td><td>–</td><td>–</td><td>0.03</td><td>3</td></tr><tr><td>HD 125079</td><td>0.12</td><td>0.00</td><td>-0.22</td><td>0.07</td><td>0.05</td><td>0.19</td><td>-0.03</td><td>0.12</td><td>0.03</td><td>0.21</td><td>-0.03</td><td>-0.01</td><td>1</td></tr><tr><td></td><td>0.34</td><td>0.05</td><td>–</td><td>–</td><td>0.03</td><td>0.02</td><td>0.00</td><td>0.45</td><td>-0.06</td><td>–</td><td>-0.22</td><td>0.07</td><td>4</td></tr><tr><th>Star name</th><th>[Sr I/Fe]</th><th>[Y I/Fe]</th><th>[Y II/Fe]</th><th>[Zr I/Fe]</th><th>[Zr II/Fe]</th><th>[Ba II/Fe]</th><th>[La II/Fe]</th><th>[Ce II/Fe]</th><th>[Nd II/Fe]</th><th>[Sm II/Fe]</th><th>[Eu II/Fe]</th><th>References</th></tr><tr><td>HD 77912</td><td>–</td><td>0.27</td><td>0.46</td><td>0.19</td><td>0.52</td><td>0.56</td><td>0.25</td><td>0.29</td><td>0.49</td><td>0.51</td><td>0.15</td><td>1</td></tr><tr><td></td><td>0.47</td><td>0.36</td><td>–</td><td>0.14</td><td>–</td><td>–</td><td>0.43</td><td>0.81</td><td>0.47</td><td>0.31</td><td>0.24</td><td>2</td></tr><tr><td>HD 101079</td><td>0.14</td><td>0.36</td><td>0.71</td><td>0.32</td><td>0.59</td><td>0.48</td><td>0.44</td><td>0.53</td><td>0.70</td><td>0.36</td><td>0.03</td><td>1</td></tr><tr><td></td><td>–</td><td>–</td><td>0.63</td><td>–</td><td>0.44</td><td>–</td><td>0.43</td><td>0.28</td><td>0.51</td><td>–</td><td>–</td><td>3</td></tr><tr><td>HD 125079</td><td>0.83</td><td>0.89</td><td>0.97</td><td>0.90</td><td>1.09</td><td>1.15</td><td>1.03</td><td>1.06</td><td>1.02</td><td>0.61</td><td>0.34</td><td>1</td></tr><tr><td></td><td>1.59</td><td>–</td><td>1.05</td><td>–</td><td>–</td><td>1.06</td><td>–</td><td>0.93</td><td>1.16</td><td>0.56</td><td>–</td><td>4</td></tr><tr><td>HD 160538</td><td>–</td><td>0.25</td><td>–</td><td>0.16</td><td>–</td><td>0.25</td><td>0.14</td><td>0.33</td><td>0.59</td><td>0.37</td><td>0.05</td><td>1</td></tr><tr><td></td><td>0.26</td><td>0.60</td><td>–</td><td>–</td><td>0.33</td><td>–</td><td>0.43</td><td>0.18</td><td>0.00</td><td>–</td><td>–</td><td>5</td></tr></tbody></table>

Notes. References: 1. This work; 2. Average abundances provided by Luck (2014); 3. Pereira et al. (2011); 4. Karinkuzhi & Goswami (2014); 5. Merle et al. (2016).

abundances of additional elements and discussed the Ba classifications for these two stars by taking into account their [s/Fe]. As shown in Fig. 5, the [s/Fe] ratio of the Ba-unknown star HD 115927 ([s/Fe] = 0.99 dex) is similar to that ([s/Fe] = 1.02 dex) of the strong Ba-star HD 125079 (Lu 1991). Thus, we classify HD 115927 as a strong Ba-star. For the Ba-likely star HD 160538, its [s/Fe] ratio (=0.33 dex) is in the [s/Fe] range of the mild Ba-stars (Lu 1991). Consequently, we classify HD 160538 as a mild Ba-star. In summary, among our sample stars, HD 92482, HD 115927, and HD 125079 belong to strong Ba-stars, while the other 17 are mild Ba-stars. This enlarges the sample of Ba-stars, including both strong and mild classifications.

## 5.2 Abundances of Na to Ni

The chemical abundances of Ba-stars provide valuable insights into nucleosynthesis processes across stars of different masses and the chemical evolution of binary or triple systems. Fig. 6 shows the abundance ratios [X/Fe] as a function of metallicity for elements ranging from Na to Ni. Additionally, we included the abundances of 253 Ba-giants analysed by Junqueira & Pereira (2001), Liang et al. (2003), Allen & Barbut (2006a, 2006b), Smiljanic et al. (2007), Liu et al. (2009), Pereira & Drake (2009), Pereira et al. (2011), de Castro et al. (2016), Mahanta et al. (2016), Yang et al. (2016) for comparison. As shown in the figure, the abundance distribution
<figure><img src="image_8.png" /><figcaption>**Figure 5.** The ratios of [s/Fe] are plotted against [Fe/H] of the sample stars. The red square and star refer to the two sample stars HD 115927 and HD 160538, respectively. The filled circles represent the other 18 sample stars. The dashed line at [s/Fe] = 0.25 dex represents the threshold for determining whether a star is s-rich. The s-abundance of a star refers to the mean value of Sr, Y, Zr, Ba, La, Ce, and Nd abundances. The typical error bar is plotted in the bottom left-hand corner of the figure.</figcaption></figure>

<figure><img src="image_9.png" /><figcaption>**Figure 6.** Abundance ratios [X/Fe] versus [Fe/H] for the elements from Na to Ni of the sample stars in this work (filled circles) and those in previous studies (crosses): Junqueira & Pereira (2001), Liang et al. (2003), Allen & Barbuy (2006a), Smiljanic et al. (2007), Liu et al. (2009), Pereira & Drake (2009), Pereira et al. (2011), de Castro et al. (2016), Mahanta et al. (2016), Yang et al. (2016).</figcaption></figure>

of most elements among our 20 Ba-stars (represented by filled circles) closely aligns with those from previous studies (indicated by crosses).

As shown in Fig. 6, the abundance characteristics of these elements in Ba-stars are similar to those in normal field stars. It is evident that the [Na/Fe] and [Al/Fe] ratios exhibit no discernible trend with increasing [Fe/H]. The α-elements Mg, Si, Ca, and Ti are dominantly generated in massive stars with $M > 10 M_⊙$ (Woosley & Weaver 1995). In Fig. 6, the [X/Fe] ratios for all the α-elements demonstrate a slight decrease with increasing [Fe/H]. In addition, the (normal) Mg abundances play in favour of low-mass nature of the polluter AGB stars. For the iron-peak elements Sc, Cr, Mn, and Ni, although

some scatter is observed, their abundances remain consistent with those stars with corresponding metallicities. It is important to note that the elemental abundances of neutral lines are adopted wherever available.

### 5.3 Abundances of s-process elements

The chemical abundances of s-process elements serve as important tools in studying nucleosynthesis processes and the formation history of Ba-stars. Fig. 7 displays the [X/Fe] ratios for the n-capture elements Sr, Y, Zr, Ba, La, Ce, Nd, and Sm for the 20 sample stars (represented by filled circles) in comparison with Ba-stars analysed in previous studies (represented by crosses). The abundance distribution of these s-process elements shows a good agreement with the previous literature. From the figure, we can see that most of the Ba-stars exhibit overabundances of s-process elements to varying degrees. It is worth noting that s-process nucleosynthesis contributes about 30 per cent of the solar Sm (Sneden et al. 2008). Therefore, although its enhancements are appreciable in most Ba-stars, as shown in Fig. 7, Sm is not considered as a truly s-process element. Furthermore, the Galactic Chemical Evolution (GCE) models are not able to reproduce the abundance levels of s-elements generally observed in Ba-stars (Kobayashi, Karakas & Lugaro 2020).

The lighter s-process elements Sr, Y, and Zr constitute the first peak of s-process abundances. Among our 20 sample stars, Sr abundances can be determined using Sr I line at 4607 Å for only eight stars, as no reliable abundance determinations are possible for any Sr II lines in the spectra. Among these eight stars, seven exhibit slight enhancement of Sr with $0.03 \le [\text{Sr I/Fe}] \le 0.83$ dex, while the star HD 107950 displays a mild underabundance of Sr with $[\text{Sr I/Fe}] = -0.16$ dex. In LTE, the Sr abundances derived from Sr I line at 4607 Å are underestimated by about 0.1–0.2 dex (Bergemann et al. 2012). The discrepancy between LTE abundances derived from neutral and ionized lines can also be seen for Y. It has been found that for stars close to solar metallicity, the NLTE effects for all Y II lines are positive and do not exceed 0.12 dex, while the Y I lines can reach up to ~0.5 dex (Alexeeval, Wang & Zhao 2023) therefore we adopt the Y abundance from Y II lines. It can be seen from Fig. 7 that Y and Zr abundances are significantly overabundant relative to the solar values.

The heavier s-process elements Ba, La, Ce, and Nd constitute the second peak of s-process abundances. The abundances of heavier s-process elements in our Ba-stars exhibit a large scatter, and the patterns in this work are consistent with those in previous literature. For our 20 sample stars, the abundances of the typical s-process element Ba are measured from lines at 5853.67, 6141.71, and 6496.91 Å. The three strong Ba-stars HD 92482, HD 115927, and HD 125079 exhibit strong enhancement of Ba with $[\text{Ba/Fe}] = 1.44, 1.07$, and $1.15$ dex, respectively. The other 17 Ba-stars display mild enhancements of Ba abundances with $0.19 \le [\text{Ba/Fe}] \le 0.75$ dex. As expected, all the sample stars exhibit enhancement of not only Ba but also the other heavier s-process elements.

The [hs/ls] ratios of Ba-stars serve as observational constraints on neutron exposures. Consequently, these ratios have been widely used to indicate s-process efficiency and are frequently incorporated by AGB nucleosynthesis models. Fig. 8 shows the [hs/ls] ratios as a function of [Fe/H] for our sample stars, along with additional Ba-stars from the literature, similar to Fig. 6. Here, [hs] and [ls] represent the mean abundances of heavier s-process elements (Ba, La, Ce, and Nd) and lighter s-process elements (Sr, Y, and Zr), respectively. From the figure, we can see that the [hs/ls] distribution
<figure><img src="image_10.png" /><figcaption>**Figure 7.** Abundance ratios [X/Fe] versus [Fe/H] for the n-capture elements of Ba-stars. The meanings represented by the filled circles and crosses are the same as in Fig. 6, but with the addition of data from Roriz et al. (2021b).</figcaption></figure>

<figure><img src="image_11.png" /><figcaption>**Figure 8.** Abundance ratios [hs/ls] versus [Fe/H] for Ba-stars. The red square, star, upper and lower triangle represent the stars HD 92482, HD 150430, HD 151101, and HD 177304, respectively. The filled circles and crosses are the same as in Fig. 6. The solid line is a linear fit for the whole sample.</figcaption></figure>

of our sample stars aligns with the data in the literature. The four stars HD 92482, HD 150430, HD 151101, and HD 177304, whose chemical abundances are analysed for the first time in the present work, have [hs/ls] ratios within the range of the literature, indicating that these four objects are typical Ba-stars. The [hs/ls] ratios exhibit an anticorrelation with metallicity, as discussed previously by Cseh et al. (2018) and Roriz et al. (2021b), which indicates that at lower metallicities ([Fe/H] < -0.6), the hs elements are the dominant products of n-capture processes in AGB models. In contrast, at higher metallicities, the ls elements become predominant (Busso et al. 2001). Negative values for [hs/ls] suggest that ls elements are more abundantly produced at the corresponding metallicities.

## 5.4 R-process characteristics of Ba-stars

The typical r-process element Eu in our 20 sample stars exhibits slight overabundances with $0.02 \le [Eu/Fe] \le 0.57$ dex. For all the Ba-stars, the abundance scatter of Eu is obviously smaller than that of

<figure><img src="image_12.png" /><figcaption>**Figure 9.** Abundance ratios [Ba/Fe] versus [Eu/Fe] for the Ba-stars. The red square, star, upper and lower triangle represent the stars HD 92482, HD 150430, HD 151101, and HD 177304, respectively. The filled circles refer to the other 16 sample stars. The dashed line is a one-to-one correlation.</figcaption></figure>

both the lighter and the heavier s-process elements, which suggests that Eu has not been significantly affected by the AGB ejecta, as depicted in Fig. 7.

Generally, Ba-stars display enhancements of s-process elements, with their abundance anomalies often explained by binary or triple system scenarios (Allen & Barbuy 2006a; Stancliffe 2021). However, by comparing the observed abundances of 12 Ba-stars with parametric model, Cui et al. (2014) found six Ba-stars having significant r-process characteristics. They divided the Ba-stars into normal and r-rich Ba-stars. Lugaro et al. (2012) proposed a stable multiple-star-system comprising a former AGB star and a former SN II, to account for the stars with enrichment in both s- and r-process elements. The i-process, initially suggested by Cowan & Rose (1977), is also a significant hypothesis for explaining the formation of r/s stars
<table><caption>**Table 10.** Signed distances of the sample stars.</caption><thead><tr><th>Star name</th><th>d<sub>s</sub></th><th>Star name</th><th>d<sub>s</sub></th></tr></thead><tbody><tr><td>HD 77912</td><td>0.76</td><td>HD 133208</td><td>0.79</td></tr><tr><td>HD 85503</td><td>0.83</td><td>HD 136138</td><td>0.95</td></tr><tr><td>HD 92482</td><td>1.26</td><td>HD 150430</td><td>0.62</td></tr><tr><td>HD 101079</td><td>0.99</td><td>HD 151101</td><td>0.84</td></tr><tr><td>HD 107950</td><td>0.74</td><td>HD 160507</td><td>1.01</td></tr><tr><td>HD 115927</td><td>1.24</td><td>HD 160538</td><td>0.76</td></tr><tr><td>HD 119650</td><td>0.74</td><td>HD 168532</td><td>0.78</td></tr><tr><td>HD 125079</td><td>1.16</td><td>HD 176524</td><td>0.65</td></tr><tr><td>HD 130255</td><td>0.82</td><td>HD 176670</td><td>0.93</td></tr><tr><td>HD 131873</td><td>0.83</td><td>HD 177304</td><td>0.74</td></tr></tbody></table>

(Roederer et al. 2016; Goswami P. & Goswami 2020; Karinkuzhi et al. 2021). However, Choplin, Siess & Goriely (2021) found that the i-process occurring in AGB stars cannot fully explain the chemical abundances of giant stars in binaries involving pollution by a former AGB companion.

Ba and Eu are considered as typical s-process and r-process elements, respectively therefore many researchers have analysed the s- and r-process characteristics of stars based on the [Ba/Fe] and [Eu/Fe] ratios (e.g. Beers & Christlieb 2005; Masseron et al. 2010). Fig. 9 shows the abundance ratios [Ba/Fe] versus [Eu/Fe] for our 20 Ba-stars. From the figure, we can see that all of the sample stars have [Ba/Eu] > 0, indicating that they all belong to normal Ba-stars. Among the four stars whose chemical abundances were determined for the first time, HD 92482 exhibits the highest [Ba/Fe] and [Ba/Eu] values, indicating its most pronounced s-process characteristics. HD 150430, HD 151101, and HD 177304 also lack r-process characteristics.

In addition to the approach that relies on two typical elements, Karinkuzhi et al. (2021) proposed using the abundances of eight heavy elements to define a signed distance, $d_s$, for re-classifying carbon-enhanced metal-poor (CEMP) stars as either CEMP-r, CEMP-s, or CEMP-rs. Based on the new classification approach, they confirmed the CEMP-s stars that had been identified using the abundance ratios of the typical s-process element La and the r-process element Eu. Additionally, they reclassified two stars, previously classified as CEMP-rs, as CEMP-s objects.

Following the methodology outlined by Karinkuzhi et al. (2021), we discuss the s- and r-characteristics of the sample stars. This criterion is based on the signed distances ($d_s$) of a sample star from the solar r-process abundance. The $d_s$ values are computed as follows:

$$d_s = \frac{1}{N} \sum_{i=1}^{N} (A_{i,*} - A_{i,r,\odot}), \qquad (3)$$

where, $A_{i,*} = \log_{10}(n_{i,*}/n_{H,*})+12$ with $n_{i,*}$ the number density of a n-capture element i. $A_{i,r,\odot}$ is the solar r-process abundance which is adopted from Goriely (1999), and has been scaled to the Eu abundance of a target star. The element set i considered here is $i = \{\text{Sr, Y, Zr, Ba, La, Ce, Nd, and Sm}\}$. In this new classification scheme, $d_s$ increases for stars with higher s-process element enrichment. This indicates that the observed abundance pattern deviates more significantly from the standard solar r-process. Following the definition described by Karinkuzhi et al. (2021), $d_s = 0.6$ is set as separation between normal and r-rich Ba-stars. Table 10 lists the derived $d_s$ values of the 20 sample stars. As indicated in the table, all our sample stars are categorized as normal Ba-stars, which means that the Ba enhancement observed in our

<table><caption>**Table 11.** Mass estimations of companions of the Ba-stars.</caption><thead><tr><th>Star name</th><th>Mass (M<sub>☉</sub>)</th><th>Star name</th><th>Mass (M<sub>☉</sub>)</th></tr></thead><tbody><tr><td>HD 77912</td><td>2.0</td><td>HD 133208</td><td>4.0</td></tr><tr><td>HD 85503</td><td>2.5</td><td>HD 136138</td><td>3.0</td></tr><tr><td>HD 92482</td><td>2.0</td><td>HD 150430</td><td>4.0</td></tr><tr><td>HD 101079</td><td>4.0</td><td>HD 151101</td><td>5.0</td></tr><tr><td>HD 107950</td><td>4.0</td><td>HD 160507</td><td>4.0</td></tr><tr><td>HD 115927</td><td>3.0</td><td>HD 160538</td><td>1.5</td></tr><tr><td>HD 119650</td><td>2.5</td><td>HD 168532</td><td>4.0</td></tr><tr><td>HD 125079</td><td>2.5</td><td>HD 176524</td><td>5.0</td></tr><tr><td>HD 130255</td><td>4.0</td><td>HD 176670</td><td>5.0</td></tr><tr><td>HD 131873</td><td>5.0</td><td>HD 177304</td><td>5.0</td></tr></tbody></table>

<figure><img src="image_13.png" /><figcaption>**Figure 10.** Comparison of the observed abundance patterns of the six sample stars (represented by filled circles with error bars) with the FRUITY AGB models (shown as curves), which exhibit the closest fits to the observational data following the application of dilution. ‘Ext’ refers to the models with an extend $^{13}$C pocket (labelled as ‘TAIL’ in the FRUITY models).</figcaption></figure>

sample stars should be attributed to material transfer from the companions.

Karinkuzhi et al. (2021) applied this approach in the context of CEMP stars. In their Fig. 9, the CEMP-s stars have $d_s$ values in the range of about 0.6–1.4. As shown in Table 10, the results for our Ba-stars are also in a similar range. This indicates that the approach and threshold used for classifying CEMP stars may also be applicable in the identification of Ba-stars. HD 92482 has the highest $d_s$ value among all the stars in our sample, indicating that it is the most s-rich star.

## 5.5 Comparison of the observations to the models of AGB stars

By using low-mass star models, Cristallo et al. (2011) investigated nucleosynthesis processes in AGB stars and provided the FRUITY data base dedicated to the nucleosynthesis occurring in AGB stars. The models span a mass range from 1.3 to 6.0 M<sub>☉</sub> and metallicities Z between $2 \times 10^{-5}$ and $3 \times 10^{-2}$. The key assumptions of these models are: the primary neutron source is the radiative burning of $^{13}$C within the pockets, and the efficiency of the third dredge-up (TDU) is inversely correlated with metallicity. At a fixed metallicity, surface chemical enrichment of AGB stars mainly depends on the mass of dredged-up material and the extent of the convective envelope.

In this study, we compared the n-capture abundance patterns of our sample stars with the prediction of (non-rotation) FRUITY AGB
models. It should be noted that these models cannot be directly compared with the observations of Ba-stars because the effect of dilution after mass transfer from the former AGB companions cannot be ignored. For the Ba-giants, the dilution mainly results from the mass transfer and mixing. The mixing is mostly due to the first dredge-up, which mixes the AGB material transferred from the companions ($M_{AGB,trans}$) and the original envelope mass ($M_{Ba,env}$). In this case, the dilution factor is given by $dil = (M_{AGB,trans} + M_{Ba,env})/M_{AGB,trans}$ (Becker & Iben 1979; Bisterzo et al. 2012). The theoretical abundances of AGB models for individual Ba-stars were derived following the methodology outlined in Cseh et al. (2022). The diluted model abundances ($[X/Fe]_{dil}$) were calculated as follows:

$$[X/Fe]_{dil} = \log(10^{[X/Fe]_{ini}} \times (1 - \delta) + 10^{[X/Fe]_{AGB}} \times \delta), \quad (4)$$

where $[X/Fe]_{ini}$ represents the initial abundances in the models, $[X/Fe]_{AGB}$ indicates the final abundances presented in the AGB models. The value of $\delta$ (= 1/dil) was derived by scaling the [Ba/Fe] of each AGB model to that of the program star. Subsequently, the $[X/Fe]_{dil}$ of each element of each model can be calculated by using the same $\delta$ value. To identify the best-match model for a Ba-star, we preferentially selected models with similar metallicity to the target star and then employed the $\chi^2$ method. We select models with $\delta < 1$, as $\delta = 1$ is unrealistic since it implies that the AGB material is deposited on top of the Ba-star without any mixing.

Taking the six stars, HD 119650, HD 125079, HD 92482, HD 150430, HD 151101, and HD 177304, as examples, Fig. 10 shows the comparisons of abundance patterns between the target stars and the best-match models. From the figure, we can see that the diluted model patterns of AGB stars with $M = 2.0-5.0 M_{\odot}$ closely resemble the observed abundance patterns of these six sample stars, suggesting that the n-capture elements of these Ba-stars are accreted from the WD companions with an initial mass range from 2.0 to $5.0 M_{\odot}$.

Regarding the dilution factors we find, the calculated $\delta$ values are mostly below 0.3. However, some systems require larger $\delta$ values, which is similar to the findings of Cseh et al. (2022). We present the $\delta$ values for selected sample stars in Fig. 10. From the figure, it is clear that the $\delta$ values of models with stellar masses $M \le 4.0 M_{\odot}$ are around or below 0.30, whereas for models with higher stellar masses, e.g. $M = 5.0 M_{\odot}$, the $\delta$ values are much larger.

The estimated masses of the companions of all our sample stars are listed in Table 11. As shown in the table, the initial masses of companions for the 17 mild Ba-stars range from 1.5 to $5.0 M_{\odot}$, with an average of approximately $3.8 M_{\odot}$. Meanwhile, the initial masses of companions for the three strong Ba-stars range from 2.0 to $3.0 M_{\odot}$, with an average value of $2.5 M_{\odot}$, which is slightly lower than that of the companions of mild Ba-stars. Consequently, our results are consistent with numerous prior investigations demonstrating the low mass characteristic of TP-AGB stars, which have contaminated the envelopes of Ba-stars (Allen & Barbuy 2006a; Cseh et al. 2018; Den Hartogh et al. 2023; Goswami & Goswami 2023; Roriz et al. 2024).

In order to check the reliability of mass estimations by comparisons of the observed abundance patterns of Ba-stars with the FRUITY models, we furthermore compared the [hs/ls] ratios of Ba-stars with those of FRUITY models with different stellar masses, and presented the comparisons in Fig. 11. From the figure, we can see that models with stellar masses between 1.5-5.0 $M_{\odot}$ can match the [hs/ls] ratios for [Fe/H] $\lesssim -0.6$. However, for higher [Fe/H], these

<figure><img src="image_14.png" /><figcaption>Figure 11. [hs/ls] ratios as a function of [Fe/H] for the Ba-stars and the FRUITY AGB models with the indicated stellar masses (shown as different coloured curves). The symbols representing the Ba-stars are the same as in Fig. 8.</figcaption></figure>

<figure><img src="image_15.png" /><figcaption>Figure 12. The Toomre diagram of $(U_{LSR}^2 + W_{LSR}^2)^{1/2}$ versus $V_{LSR}$ for our sample stars. The semicircular line delineates the boundary for thin disc stars, with $V_{tot} = (U_{LSR}^2 + V_{LSR}^2 + W_{LSR}^2)^{1/2} < 85$ km s$^{-1}$ (Chen et al. 2004).</figcaption></figure>

models do not cover the entire data range. In the figure, the data for HD 130255, HD 131873, and HD 176670 fall near the line of the $5.0 M_{\odot}$ model, aligning with the models selected in Table 11. It is noteworthy that the abundance pattern of the $5.0 M_{\odot}$ ([Fe/H] = -0.85) model best matches that of HD 130255; however, it needs unexpected $\delta = 1.67$ (>1). Therefore, the mass of the companion star is estimated to be $4.0 M_{\odot}$ with a reasonable $\delta$ of 0.50 (<1), although the corresponding model-observation match is slightly worse than that of the $5.0 M_{\odot}$ model. HD 125079 lies on the line of the $2.5 M_{\odot}$ model. HD 77912 and HD 136138 also fall within the region corresponding to models with stellar masses lower than $4.0 M_{\odot}$. The distribution of [hs/ls] ratios of these six sample stars discussed above all align with the models we selected in Table 11. However, the remaining sample stars either span a broad range within the FRUITY models or fall outside their range, making the [hs/ls] ratios ineffective for selecting suitable AGB models for these stars.
<table><caption>**Table 12.** Comparisons of radial velocities with literature values.</caption><thead><tr><th>Star name</th><th>V<sub>r</sub> (km s<sup>-1</sup>)<br>(our estimates)</th><th>V<sub>r</sub> (km s<sup>-1</sup>)<br>Gaia DR3</th><th>V<sub>r</sub> (km s<sup>-1</sup>)<br>(from the literature)</th><th>References</th></tr></thead><tbody><tr><td>HD 77912</td><td>15.9</td><td>16.1</td><td>15.4</td><td>1</td></tr><tr><td>HD 85503</td><td>13.5</td><td>13.7</td><td>14.1</td><td>2</td></tr><tr><td>HD 92482</td><td>-7.2</td><td>-7.8</td><td>-</td><td>-</td></tr><tr><td>HD 101079</td><td>-0.8</td><td>-3.9</td><td>-2.77</td><td>3</td></tr><tr><td>HD 107950</td><td>-14.7</td><td>-14.2</td><td>-14.2</td><td>1</td></tr><tr><td>HD 115927</td><td>1.4</td><td>4.2</td><td>2.7</td><td>4</td></tr><tr><td>HD 119650</td><td>-4.7</td><td>-8.3</td><td>-5.7</td><td>5</td></tr><tr><td>HD 125079</td><td>-4.0</td><td>-7.0</td><td>-4.5</td><td>6</td></tr><tr><td>HD 130255</td><td>41.5</td><td>41.0</td><td>40.6</td><td>7</td></tr><tr><td>HD 131873</td><td>16.8</td><td>17.0</td><td>16.6</td><td>4</td></tr><tr><td>HD 133208</td><td>-18.3</td><td>-18.4</td><td>-18.2</td><td>1</td></tr><tr><td>HD 136138</td><td>-5.9</td><td>-10.8</td><td>-7.28</td><td>8</td></tr><tr><td>HD 150430</td><td>6.8</td><td>7.7</td><td>4.9</td><td>9</td></tr><tr><td>HD 151101</td><td>-1.4</td><td>-1.4</td><td>-1.0</td><td>10</td></tr><tr><td>HD 160507</td><td>-15.8</td><td>-10.4</td><td>-14.5</td><td>11</td></tr><tr><td>HD 160538</td><td>-10.1</td><td>-12.0</td><td>-9.0</td><td>10</td></tr><tr><td>HD 168532</td><td>-1.5</td><td>-17.3</td><td>-15.9</td><td>9</td></tr><tr><td>HD 176524</td><td>-16.5</td><td>-10.0</td><td>-8.75</td><td>10</td></tr><tr><td>HD 176670</td><td>-17.3</td><td>-17.7</td><td>-17.7</td><td>9</td></tr><tr><td>HD 177304</td><td>-10.6</td><td>-11.3</td><td>-</td><td>-</td></tr></tbody></table>

Note. References: 1. Takeda et al. (2008); 2. Hourihane et al. (2023); 3. Pereira et al. (2011); 4. Jönsson et al. (2017); 5. Gontcharov (2006); 6. Karinkuzhi & Goswami (2014) 7. de Castro et al. (2016); 8. Massarotti et al. (2008); 9. Famaey et al. (2005); 10. Tautvaišienė et al. (2020); 11. Wang et al. (2011).

In this section, we identified five FRUITY models with a stellar mass of 5.0 M☉ that best fit the observations. Typically, the masses of Ba-star companions are considered to be below 3.0 or 4.0 M☉ (e.g. Husti et al. 2009; Cseh et al. 2018; Shejeelammal et al. 2020; Roriz et al. 2021a). However, Cseh et al. (2022) raised the question of where the Ba-stars polluted by higher mass AGB companions are. They predicted that the stars with [s/Fe] < 0.25 are expected from theoretical models of mass over 4.0 M☉. We also notice that there are two sample stars (HD 77912 and HD 168532) whose masses are higher than the companions we estimated. The masses of these two stars are estimated to be 4.5 and 6.0 M☉ in this study. Generally, the typical masses of Ba-stars are found to be lower than 3.0 M☉. However, based on the orbital parameters and abundances of Ba-stars, Jorissen et al. (2019) predicted that the masses of mild Ba-stars could extend up to 5.0 M☉. Additionally, Roriz et al. (2021a) pointed out that, according to a simple Salpeter initial mass function (IMF), there should be 14 per cent of Ba-stars with masses up to 4.0–6.0 M☉. These two unusual cases may reflect the complexity of mass transfer during the evolution of binary star systems. To better understand these phenomena, further theoretical and observational efforts are required.

## 6 KINEMATIC ANALYSIS

The analysis of star motions in the vicinity of the Sun relies on measurements of their velocities along three key axes: *U*, *V*, and *W*. These axes correspond to directions pointing towards the Galactic Centre, the path of Galactic rotation, and the North Galactic Pole, respectively. By leveraging the precise stellar data of parallax, proper motion, and radial velocity, we can accurately calculate their velocities in space. In our study, we acquired proper motion and parallax data for the 19 sample stars from Gaia DR3 (Gaia Collaboration et al. 2023). For HD 131873 lacking proper motion

and parallax measurements in Gaia DR3, we acquired the data from SIMBAD data base.

The determination of the space velocity relative to the Sun follows the approach outlined by Johnson & Soderblom (1987). Adjustments to the space velocity to align with the Local Standard of Rest are made using a solar motion of ($U_{LSR}$, $V_{LSR}$, $W_{LSR}$)$_☉$ = (11.1, 12.2, 7.3) km s$^{-1}$ (Schönrich, Binney & Dehnen 2010). The Toomre diagram of ($U_{LSR}^2$+ $W_{LSR}^2$)$^{1/2}$ as function of $V_{LSR}$ is presented in Fig. 12. The semicircular line delineates the boundary for thin disc stars, with $V_{tot}$ = ($U_{LSR}^2$ + $V_{LSR}^2$ + $W_{LSR}^2$)$^{1/2}$ < 85 km s$^{-1}$, as proposed by Chen., Nissen & Zhao (2004). From the figure, we can see that the majority of the stars in the sample are classified within the thin disc, with the sole exception being the star HD 130255, which can be classified as belonging to the thick disc (see also de Castro et al. 2016). HD 130255 has the lowest metallicity ([Fe/H] = -0.86) among our sample stars, which is consistent with the kinematic results.

Radial velocities of the 20 sample stars were determined through the Doppler shift of the elemental absorption lines and were subsequently corrected for heliocentric motion. The derived results are presented in Table 12. For comparison, values from Gaia DR3 and other literature sources are also provided. All the sample stars are low-velocity objects ($V_t$ < 50 km s$^{-1}$). No literature reported radial velocity variations for the sample stars. Our radial velocities for HD 101079, HD 115927, HD 119650, HD 125079, HD 136138, HD 150430, HD 160507, HD 160538, HD 168532, and HD 176524 show significant differences from the Gaia DR3 and literature values, which suggest that these stars are in binary systems. For the remaining stars, we note only small differences from the literature values, and those stars might actually be related to binaries with pole-on orbits or very eccentric orbits, where significant radial velocity variations occur only in a small phase range (Pourbaix et al. 2004; de Castro et al. 2016).
## 7 CONCLUSIONS

Based on the high resolution and high S/N spectra obtained from OHP, we derive the chemical abundances of 20 elements (Na, Mg, Al, Si, Ca, Sc, Ti, Cr, Mn, Fe, Ni, Sr, Y, Zr, Ba, La, Ce, Nd, Sm, and Eu) for 20 Ba-stars. The stellar atmospheric parameters of the sample stars cover a range of $4150 \le T_{\text{eff}} \le 5334$ K, $1.35 \le \log g \le 3.20$ dex, $-0.86 \le [\text{Fe/H}] \le 0.38$ dex, and $1.20 \le \xi_i \le 2.05$ km s$^{-1}$. For the four sample stars, namely HD 92482, HD 150430, HD 151101, and HD 177304, it is the first time that a detailed elemental abundance has been determined. Furthermore, among the remaining 16 sample stars, Ba abundance has been derived for the first time in six of them. The conclusions are as follows:

(i) It is found that the [s/Fe] ratio of HD 115927 is similar to that of the strong Ba-star HD 125079, thus this star can be classified as a strong Ba-star. The [s/Fe] ratio of HD 160538 is in the range of mild Ba-stars.

(ii) The abundances of light odd-Z metal elements Na and Al, as well as Fe-peak elements Sc, Cr, Mn, and Ni of Ba-stars remain similar to those of stars with corresponding metallcities. The [α/Fe] ratios of Mg, Si, Ca, and Ti show slight decrease with increasing [Fe/H].

(iii) Most of the s-process elements in the Ba sample stars exhibit varying degrees of overabundance. The first peak of s-process elements (Sr, Y, and Zr) in most sample stars shows slight enrichment, while the second peak of s-process elements (Ba, La, Ce, and Nd) in all Ba-stars shows significant overabundance. The [hs/ls] ratios increase with decreasing metallicity, which aligns with theoretical nucleosynthesis models of AGB stars using $^{13}$C as the neutron source. The abundance scatter of the r-process element Eu in Ba-stars is smaller than that of s-process elements, suggesting that very little Eu can be produced by the s-process in AGB stars.

(iv) The signed distances for the sample stars indicate that all of them are normal Ba-stars, which means that the enhancements of s-process elements, such as Ba, La, and Nd, should be attributed to material transfer from the companions.

(v) We derived the masses of the TP-AGB stars that previously polluted the Ba-stars. The observations can be effectively explained by the FRUITY nucleosynthesis models with low masses.

(vi) Considering the kinematic perspective, we assessed the likelihood of these stars being part of the thin disc, thick disc, or halo. Our analysis revealed that the majority of these stars are associated with the thin disc, with HD 130255 ([Fe/H] = -0.86) being the only one potentially affiliated with the thick disc.

We anticipate that the derived chemical abundances of the 20 sample stars could contribute to a deeper comprehension of the chemical evolution of Ba-stars. Further observational and theoretical studies involving a larger sample of Ba-stars are imperative to explore the abundance characteristics of this stellar group more comprehensively.

## ACKNOWLEDGEMENTS

We thank the anonymous referee for thoughtful suggestions that have improved the manuscript. We appreciate Associate Professor Kefeng Tan from NAOC for the helpful discussions and constructive suggestions. This work is supported by the National Natural Science Foundation of China (NSFC) under grant Nos. 11988101, 12273055, 11927804, 11673007, 11733006, 12173013, 12090040, and 12090044, the National Key R&D Program of China under grant No. Fundamental Research Funds of China West Normal University

under Grant No. 22kE041, and the S&T Program of Hebei under grant Nos. A2018106014 and A2021205006. We acknowledge the support from the 2-m Chinese Space Station Telescope project.

## DATA AVAILABILITY

This research is based on the spectral data retrieved from the ELODIE archive at Observatoire de Haute-Provence (OHP, http://atlas obs hp.fr/elodie/).

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