# The oldest stars with low neutron-capture element abundances and origins in
ancient dwarf galaxies

Hillary Diane Andales,1 Ananda Santos Figueiredo,1 Casey Gordon Fienberg,1
Mohammad K. Mardini,2,3,4,1 Anna Frebel1,4
1Department of Physics and Kavli Institute for Astrophysics and Space
Research, Massachusetts Institute of Technology, Cambridge, MA 02139, USA
2Department of Physics, Zarqa University, Zarqa 13110, Jordan
3Jordanian Astronomical Virtual Observatory, Zarqa University, Zarqa 13110,
Jordan
4Joint Institute for Nuclear Astrophysics – Center for the Evolution of the
Elements (JINA-CEE), East Lansing, MI 48824, USA
E-mail<EMAIL_ADDRESS>
(Accepted XXX. Received YYY; in original form ZZZ)

###### Abstract

We present a detailed chemical abundance and kinematic analysis of six
extremely metal-poor ($-4.2\leq$ [Fe/H] $\leq-$2.9) halo stars with very low
neutron-capture abundances ([Sr/H] and [Ba/H]) based on high-resolution
Magellan/MIKE spectra. Three of our stars have [Sr/Ba] and [Sr/H] ratios that
resemble those of metal-poor stars in ultra-faint dwarf galaxies (UFDs). Since
early UFDs may be the building blocks of the Milky Way, extremely metal-poor
halo stars with low, UFD-like Sr and Ba abundances may thus be ancient stars
from the earliest small galactic systems that were accreted by the proto-Milky
Way. We label these objects as Small Accreted Stellar System (SASS) stars, and
we find an additional 61 similar ones in the literature. A kinematic analysis
of our sample and literature stars reveals them to be fast-moving halo
objects, all with retrograde motion, indicating an accretion origin. Because
SASS stars are much brighter than typical UFD stars, identifying them offers
promising ways towards detailed studies of early star formation environments.
From the chemical abundances of SASS stars, it appears that the earliest
accreted systems were likely enriched by a few supernovae whose light element
yields varied from system to system. Neutron-capture elements were sparsely
produced and/or diluted, with $r$-process nucleosynthesis playing a role.
These insights offer a glimpse into the early formation of the Galaxy. Using
neutron-capture elements as a distinguishing criterion for early formation, we
have access to a unique metal-poor population that consists of the oldest
stars in the universe.

###### keywords:

stars: abundances – stars: Population II – Galaxy: halo – nuclear reactions,
nucleosynthesis, abundances – galaxies: dwarfs

††pubyear: 2024††pagerange: The oldest stars with low neutron-capture element
abundances and origins in ancient dwarf galaxies–References

##  1 Introduction

The Milky Way’s ultra-faint dwarf satellite galaxies (UFDs) continue to offer
a vital window into the early universe. Besides providing insight on a broad
range of astrophysical questions such as the formation of galaxies and
structure (e.g., Frebel & Bromm, 2012; Wheeler et al., 2015; Kallivayalil et
al., 2018; Applebaum et al., 2021) and the small-scale properties of dark
matter (e.g., Simon & Geha, 2007; Calabrese & Spergel, 2016; Brandt, 2016;
Safarzadeh & Spergel, 2020; Orkney et al., 2022), UFDs are uniquely important
laboratories for studying the early universe. In particular, they provide a
glimpse into early star formation environments, early chemical enrichment
events, and the early assembly of the Milky Way (e.g, Kirby et al., 2008;
Simon et al., 2010; Roederer et al., 2018; Lee et al., 2019; Brauer et al.,
2022).

Formed 12 to 13 billion years ago at around $z\sim 10$ (Brown et al., 2012,
2014; Simon et al., 2023), UFDs are ancient and thus extremely metal-poor
(Frebel et al., 2010a, b; Ji et al., 2016c; Chiti et al., 2018, 2021, 2023).
In fact, the majority of UFD stars have metallicities between
$\mbox{[Fe/H]}=-2$ and $\mbox{[Fe/H]}=-3$, with some stars reaching levels as
low as $\mbox{[Fe/H]}=-4$ (Simon, 2019). Most of their stars formed before
reionization ($z\sim 6$) after which their star formation was quenched (Weisz
et al., 2014; Brown et al., 2014). Since the epoch of reionization, UFDs have
been gas-poor, quiescent, and largely chemically unaltered (Koch et al.,
2009). As a result, their long-lived metal-poor stars now carry clean
signatures of the earliest chemical enrichment events in the universe, as well
as hints about early star formation environments.

However, further progress in exploring the early universe using UFDs is
severely hampered by the simple fact that UFDs are too distant. Because these
tiny galaxies lie dozens or hundreds of kiloparsecs away, studying their
chemical makeup with high-resolution spectroscopy is notoriously challenging.
In principle, faint stars with visual magnitudes of $V\sim 19$ are still
observable but such observations would be near the technical limits of current
telescopes. Thus only a handful of UFD stars — the brightest cool red giants —
are available for high-resolution spectroscopy. Much fainter stars are much
more difficult to observe, even with medium-resolution spectroscopy. As a
result, the current chemical abundance datasets on UFDs are still too sparse
and inadequate; only a few dozen UFD stars have detailed chemical abundances
(Suda et al., 2017; Abohalima & Frebel, 2018). To get around this, one common
solution is to combine stars from various UFDs and treat that combination as
one "population" (e.g., Vargas et al., 2013). Yet, even with this workaround,
most of the chemical abundance data still have large uncertainties.

Unraveling the histories of UFDs requires more detailed probing, much more
than the current chemical abundance datasets can allow. We thus suggest an
alternative path forward: searching in the Galactic halo for erstwhile member
stars of small satellites that were accreted by the Milky Way billions of
years ago. After all, the Milky Way assembled hierarchically, and presumably
many small "building block" systems — including possible analogs of the
surviving UFDs — contributed to its early formation (e.g., Wetzel et al.,
2016; Griffen et al., 2016). Stars from these small accreted systems could
therefore be suitable analogs of the stars in surviving UFDs. Using these
analog stars is obviously advantageous for detailed chemical study since halo
stars are more numerous and more accessible for high-resolution spectroscopy.

To isolate these analogs from the broader collection of halo stars, we look
for stars with similar chemical properties as those in extant UFDs. First,
they must be metal-poor since UFD stars are typically Very Metal-Poor (VMP;
$\mbox{[Fe/H]}<-2$) or Extremely Metal-Poor (EMP; $\mbox{[Fe/H]}<-3$). Second,
unlike VMP and EMP stars in the halo, we would expect most UFD stars not to
show significant carbon enhancements (where carbon enhancement is defined as
[C/Fe] $>$ 0.7). This discrepancy is well-established and works in our favor
here (Frebel et al., 2010b). Third, they must have extraordinarily low
abundances of neutron-capture elements (e.g., Sr and Ba). Finally, they must
have retrograde motion consistent with an accretion origin. Previous works
have used some combination of these criteria to associate halo stars with
their possible origins in now-destroyed UFD-like satellites, effectively
finding some of our possible analog stars (Hansen et al., 2016a; Lee et al.,
2017; Yoon et al., 2018; Yuan et al., 2020). Intriguingly, Sestito et al.
(2024) even found a VMP star in the inner Galaxy with putative origins in an
accreted UFD.

Low neutron-capture abundance is the hallmark characteristic of UFD stars that
sets them apart from the general halo trend (Ji et al., 2019). Many chemical
abundance studies have consistently found low [Sr/Fe] and [Ba/Fe] across all
UFD stars, with their neutron capture abundances being as much as $\sim 1$ dex
lower than a halo star of the same metallicity (Frebel et al., 2010a, b, 2014;
Frebel, 2018; Ji et al., 2020). When normalizing the Sr and Ba abundances to H
instead of Fe to determine the overall neutron-capture enrichment level (i.e.,
by examining [Sr,Ba/H] instead of [Sr,Ba/Fe]), Frebel et al. (2014) found a
cleaner cut between halo stars and UFD stars, such that the latter are
predominantly found with abundance levels at $\mbox{[Sr/H]}\lesssim-4$ and
$\mbox{[Ba/H]}\lesssim-3.5$. Two exceptions to this rule are Reticulum II and
Tucana III which show an overabundance of neutron-capture elements due to a
rare $r$-process enrichment event (Ji et al., 2016b; Roederer et al., 2016;
Marshall et al., 2019). Even then, two of the stars in Reticulum II still
display the characteristic low neutron-capture abundances. Indeed, this
chemical criterion has been crucial in evaluating the membership of candidate
UFD stars (e.g., Chiti et al., 2023).

In this paper, we thus set out to identify metal-poor halo stars that possibly
originated from small satellite galaxies, and were presumably accreted by the
Milky Way long ago. We developed abundance criteria in line with chemical
signatures observed among UFD stars. We then apply these criteria to our
program stars and to other stars with chemical abundances from the literature.
Our goal was to assemble a set of stars that could be used as local probes of
the population of small satellites that assembled the Galactic halo.

We later label these probes "Small Accreted Stellar System stars" or "SASS
stars." These long-lived SASS stars offer an opportunity to study, in detail,
the nucleosynthesis histories of early small galaxies, their environments, and
early star formation. This complements the many ongoing studies of current
UFDs. Detailed comparisons of the chemical abundances of our newly identified
SASS stars with those in surviving UFDs would allow for further establishing
to what extent the surviving UFDs are related to the first galaxies, or at
least, the earliest assembled systems (Frebel et al., 2014).

In Section 2, we describe the initial target selection of our candidates and
the spectroscopic observations. Sections 3 and 4 detail the chemical abundance
analysis and our results. Section 5 presents the kinematic analysis of our
program stars. In Section 6, we discuss our refined selection criteria for
identifying SASS stars, taking into account a detailed kinematic analysis of
our sample stars and considering the context of the Galactic halo’s assembly.
In the same section, we present possible enrichment scenarios and clues about
the environments in early dwarf galaxies. We also offer future directions for
identifying more SASS stars. Finally in Section 7, we summarize our work.

##  2 Observations and Measurements

###  2.1 Initial target selection and existing observations

All of our program stars were initially selected for high-resolution follow-up
spectroscopy based on their brightness and low [Fe/H] as determined from the
medium-resolution spectra obtained during the Hamburg/ESO survey (Wisotzki et
al., 1996; Frebel et al., 2006). For this subset of stars, high-resolution
spectroscopy was then obtained over several years (described further below)
with the goal of carrying out a detailed chemical abundance analysis. From
this sample of high-resolution spectra, we visually selected stars that showed
very weak or no Ba line at 4554 Å. This preliminary selection was to establish
an exploratory sample for the current study on halo stars with possible
origins in long-accreted small satellite systems.

Table 1: Observing details and heliocentric velocities Star | RA | Dec | UT Dates | Slit | $t_{\text{exp}}$ | $g$ | $S/N$ | $S/N$ | $V_{\text{helio}}$ | $V_{\text{helio, Gaia}}$
---|---|---|---|---|---|---|---|---|---|---
| [J2000] | [J2000] |  |  | [s] | [mag] | [4000] | [4500] | [km s^-1] | [km s^-1]
HE 0104$-$5300 | 01:06:51.9 |  $-$52:44:10.5 | Jan 6, 2013 | $1\aas@@fstack{\prime\prime}0\times 5\aas@@fstack{\prime\prime}0$ | 1800 | 14.1 | 63 | 111 | 188.3 | 186.6
HE 1310$-$0536 | 13:13:31.1 |  $-$05:52:12.5 | Mar 11, 2014 | $0\aas@@fstack{\prime\prime}7\times 5\aas@@fstack{\prime\prime}0$ | 3600 | 14.1 | 40 | 81 | 110.9 | 108.3
HE 2155$-$2043 | 21:58:42.2 |  $-$20:29:15.7 | May 31, 2013 | $0\aas@@fstack{\prime\prime}7\times 5\aas@@fstack{\prime\prime}0$ | 4339 | 13.0 | 49 | 83 |  $-$94.8 |  $-$90.8
HE 2303$-$5756 | 23:06:55.1 |  $-$57:40:33.5 | Jun 22, 2014 | $0\aas@@fstack{\prime\prime}7\times 5\aas@@fstack{\prime\prime}0$ | 1200 | 13.1 | 80 | 122 |  $-$85.4 |  $-$92.8
HE 2319$-$5228 | 23:21:58.1 |  $-$52:11:43.2 | Jun 23, 2014 | $0\aas@@fstack{\prime\prime}7\times 5\aas@@fstack{\prime\prime}0$ | 900 | 13.0 | 37 | 45 | 292.8 | 284.2
HE 2340$-$6036 | 23:43:41.1 |  $-$60:19:22.4 | Sep 25, 2014 |  $0\aas@@fstack{\prime\prime}7$ $\times$ $5\aas@@fstack{\prime\prime}0$ | 1200 | 12.5 | 45 | 74 | 212.4 | 212.1

Since data collection occurred back in 2013 and 2014, the majority of our
program stars have already been observed by other groups for a variety of
purposes. Given our generally high(er) data quality, we nevertheless decided
to (re-)analyze these stars to obtain our independent measurements.

HE 0104$-$5300 was part of the large Barklem et al. (2005) and Jacobson et al.
(2015) samples that provided detailed chemical abundances. The spectrum used
in the present study is the same spectrum analyzed by Jacobson et al. (2015),
allowing for a detailed abundance comparison between their work and this work.

HE 1310$-$0536 was also observed by Hansen et al. (2014) and noted for its low
Ba and Sr abundances. Their data had a signal-to-noise ratio ($S/N$) of
$\sim$40 at 4300 Å (65 at 6700 Å) whereas our spectrum has a $S/N$ of 40 at
4000 Å and $\sim$80 at 4500 Å, around where the strongest Sr and Ba absorption
lines are located. Hansen et al. (2014) found [Sr/Fe] = $-$1.08 and [Ba/Fe] =
$-$0.50. As we will show in Section 4, we find somewhat lower values likely
due to our increased $S/N$.

Two of our carbon-rich stars, HE 2319$-$5228 and HE 2155$-$2043, were analyzed
by Hansen et al. (2016b) and were both found to have extremely low abundances
of heavy neutron-capture elements. Our spectra are of better quality than
those presented in Hansen et al. (2016b). Their spectrum of HE 2319$-$5228 had
a $S/N$ of 20 at 4000 Å compared to our $\sim$40\. Their spectrum of HE
2155$-$2043 was reported at 40, while we have $\sim$50.

For HE 2319$-$5228, they report upper limits of [Ba/Fe] < $-$3 and [Sr/Fe] <
$-$3\. For HE 2155$-$2043, they report [Sr/Fe] = 0.2, and no value for
[Ba/Fe].

HE 2155$-$2043 was also analyzed by Purandardas & Goswami (2021), but no $S/N$
was reported of the spectra taken from the Subaru Telescope archive. They
reported [Sr/Fe] = $-$0.13 and [Ba/Fe] = $-$1.51. Our spectrum has a
relatively high $S/N$ of $\sim$50 at 4000 Å and $\sim$80 at 4500 Å. In Section
4, we will compare abundances in more detail.

HE 2340$-$6036 was included in the $r$-process Alliance’s Holmbeck et al.
(2020) data release. A $S/N$ of 22 was reported at 4129 Å of their snapshot
spectrum, compared to our $S/N$ of 45 at 4000 Å and 74 at 4500 Å. They
reported weak Sr and Ba lines, finding [Sr/Fe] of $-$1.89 and [Ba/Fe] of
$-$1.34. We will report somewhat higher abundances in Section 4.

Finally, we present the first detailed chemical abundance analysis of HE
2303$-$5756.

Table 2: Equivalent widths measurements and line abundances of our program stars |  HE 0104$-$5756 |  |  HE 1310$-$0536 |  |  HE 2155$-$2043 |  |  HE 2303$-$5756 |  |  HE 2319$-$5228 |  |  HE 2340$-$6036 |
---|---|---|---|---|---|---|---|---|---|---|---|---
Species | $\lambda$ | $\chi$ | $\log gf$ | EW |  $\log\epsilon$(X) |  | EW |  $\log\epsilon$(X) |  | EW |  $\log\epsilon$(X) |  | EW |  $\log\epsilon$(X) |  | EW |  $\log\epsilon$(X) |  | EW |  $\log\epsilon$(X) |
|  [] |  [eV] | [dex] |  [m] |  [dex] |  |  [m] |  [dex] |  |  [m] |  [dex] |  |  [m] |  [dex] |  |  [m] |  [dex] |  |  [m] |  [dex] |
CH | 4314.00 | … | … | … | 4.80 |  | … | 6.29 |  | … | 6.00 |  | … | 6.30 |  | … | 6.27 |  | … | 4.26 |
CH | 4323.00 | … | … | … | 4.74 |  | … | 6.21 |  | … | 5.99 |  | … | 5.93 |  | … | 6.26 |  | … | 4.28 |
Na i | 5682.63 | 2.10 |  $-$0.71 | … | … |  | … | … |  | … | … |  | … | … |  | 28.1 | 4.78 |  | … | … |
Na i | 5688.20 | 2.10 |  $-$0.41 | … | … |  | … | … |  | … | … |  | … | … |  | 33.2 | 4.58 |  | … | … |
Na i | 5889.95 | 0.00 | +0.11 | 113.0 | 2.84 |  | 70.4 | 2.20 |  | 139.6 | 3.52 |  | 45.2 | 3.07 |  | 265.8 | 4.51 |  | 143.4 | 3.19 |
Na i | 5895.92 | 0.00 |  $-$0.19 | 107.6 | 3.04 |  | 48.4 | 2.14 |  | 106.4 | 3.22 |  | 22.7 | 2.93 |  | 242.4 | 4.69 |  | 127.1 | 3.23 |
Mg i | 4057.51 | 4.35 |  $-$0.90 | 11.2 | 4.56 |  | … | … |  | … | … |  | … | … |  | … | … |  | 14.0 | 4.61 |
Mg i | 4167.27 | 4.35 |  $-$0.74 | 18.7 | 4.65 |  | … | … |  | … | … |  | 6.9 | 4.88 |  | … | … |  | 22.2 | 4.68 |
Mg i | 4702.99 | 4.33 |  $-$0.44 | 24.4 | 4.44 |  | 10.4 | 4.00 |  | 35.1 | 4.76 |  | 10.5 | 4.74 |  | 58.9 | 5.03 |  | 36.5 | 4.60 |
Mg i | 5172.68 | 2.71 |  $-$0.36 | 141.1 | 4.38 |  | 88.1 | 3.60 |  | 147.2 | 4.70 |  | 89.5 | 4.81 |  | 176.0 | 4.83 |  | 160.4 | 4.45 |
Mg i | 5183.60 | 2.72 |  $-$0.17 | 160.1 | 4.46 |  | 105.4 | 3.76 |  | 163.0 | 4.75 |  | 104.0 | 4.90 |  | 187.4 | 4.75 |  | 177.2 | 4.45 |
Al i | 3944.00 | 0.00 |  $-$0.64 | … | 2.60 |  | … | 2.08 |  | … | … |  | … | 2.76 |  | … | … |  | … | 2.34 |
Al i | 3961.52 | 0.01 |  $-$0.34 | … | 2.13 |  | … | 1.31 |  | … | … |  | … | 2.71 |  | … | 2.82 |  | … | 2.30 |
Si i | 3905.52 | 1.91 |  $-$1.09 | … | 4.54 |  | … | … |  | … | … |  | … | 4.67 |  | … | 4.92 |  | … | 4.61 |
Si i | 4102.94 | 1.91 |  $-$3.14 | … | … |  | … | … |  | … | … |  | … | … |  | … | 4.23 |  | … | 4.98 |
This table is available in its entirety in the online journal. |  |  |  |  |  |
Table 3: Photometry and astrometry for our program stars Star |  Gaia Source ID | J | K | V | $A_{V}$ | $E(B-V)$ | parallax (mas)
---|---|---|---|---|---|---|---
HE 0104$-$5300 | 4927175937828177280 | $11.91\pm 0.03$ | $11.29\pm 0.02$ | $13.63\pm 0.03$ | 0.0326 | $0.0108\pm 0.0004$ | $0.0999\pm 0.0112$
HE 1310$-$0536 | 3635533208672382592 | $12.53\pm 0.02$ | $12.05\pm 0.02$ | $14.38\pm 0.02$ | 0.1149 | $0.0368\pm 0.0016$ | $0.0118\pm 0.0232$
HE 2155$-$2043 | 6823219517280229120 | $11.57\pm 0.02$ | $11.02\pm 0.02$ | $13.23\pm 0.01$ | 0.0799 | $0.0257\pm 0.0006$ | $0.1240\pm 0.0301$
HE 2303$-$5756 | 6493230587155263360 | $12.23\pm 0.02$ | $11.91\pm 0.02$ | $13.24\pm 0.02$ | 0.0455 | $0.0147\pm 0.0005$ | $0.9229\pm 0.0166$
HE 2319$-$5228 | 6501398446721935744 | $11.55\pm 0.02$ | $10.96\pm 0.02$ | $13.25\pm 0.02$ | 0.0277 | $0.0087\pm 0.0005$ | $0.1616\pm 0.0158$
HE 2340$-$6036 | 6488378098744387968 | $11.05\pm 0.02$ | $10.40\pm 0.02$ | $12.83\pm 0.01$ | 0.0363 | $0.0119\pm 0.0002$ | $0.1666\pm 0.0126$

###  2.2 High-resolution spectroscopy

All of our stars were observed at Las Campanas Observatory in January and May
of 2013, as well as March, June, and September of 2014 using the Magellan
Inamori Kyocera Echelle (MIKE) instrument (Bernstein et al., 2003) on the
Magellan-Clay telescope. Data for all program stars except HE 0104$-$5300 was
taken with a slit size of $0\aas@@fstack{\prime\prime}7\times
5\aas@@fstack{\prime\prime}0$, resulting in a resolving power of $\sim$28,000
in the red wavelength and $\sim$35,000 in the blue. Data for HE 0104$-$5300
was obtained with a $1\aas@@fstack{\prime\prime}0\times
5\aas@@fstack{\prime\prime}0$ slit, yielding a resolving power of $\sim$22,000
and $\sim$28,000 for the red and blue spectra, respectively. Exposure times
ranged from 900 to roughly 4500 seconds.

We used the CarPy MIKE pipeline (the version available at the time) to reduce
all spectra (Kelson, 2003). Using the SMH custom analysis software (Casey,
2014), we normalized and stitched all blue and red echelle orders together
before Doppler correcting the spectra by cross-correlation. We employed
HD122563 as a template using the Ca triplet lines around 8500 Å. Heliocentric
velocities were determined using standard Python packages. Observational
details and radial velocity measurements are listed in Table 1. Figure 1 shows
three selected spectral regions of the final spectra of all our stars: around
the Sr line at 4077 Å, the CH G-band around 4300 Å, and the Ba line at 4554 Å.

We also list Gaia radial velocities in Table 1 for comparison with our derived
values. The measurements generally agree within a few km s-1. The largest
differences are found for HE 2319$-$5228 and HE 2303$-$5756 for which the
radial velocities differ by nearly 10 km s-1, suggesting that these stars are
part of binary systems. Searching for additional radial velocity measurements
in the SIMBAD database reveals further support for their binarity. We note
that neither of the two stars appears to be a double-lined spectroscopic
binary so our spectra are not affected by any light from a presumably faint
companion.

HE 2319$-$5228 is very carbon-enhanced which could be due to a mass transfer
event from its previously more massive companion. Since it is not $s$-process-
enriched, as is common for carbon-rich stars in binary systems, it may have
had a relatively massive erstwhile AGB companion. An alternative enrichment
scenario that could explain the observed abundances is a massive, rapidly
rotating "spinstar" that produced large amounts of carbon and then enriched
the natal gas cloud of HE 2319$-$5228 (Chiappini et al., 2006; Cescutti et
al., 2013; Meynet et al., 2006; Hirschi, 2006). If such spinstars were present
among later generations of stars (not just among Pop III stars), then HE
2319$-$5228 may have had a spinstar companion that transferred carbon-rich
mass into it.

Figure 1: Spectra of our program stars around the Sr line at 4077 Å, the
carbon G-band around 4300 Å, and the Ba line at 4554 Å. The top three spectra
(teal) are halo stars from our sample (HE 2303$-$5756, HE 2155$-$2043, HE
0104$-$5300), whereas the bottom three spectra (red) depict stars with
putative origins in long-accreted small satellite systems (HE 2319$-$5228, HE
2340$-$6036, HE 1310$-$0536). Spectra were shifted vertically to aid visual
representation. Note that HE 2303$-$5756 (teal) and HE 2319$-$5228 (red)
appear to be in binary systems.

##  3 Chemical Abundance Analysis

###  3.1 Line measurements

To determine the chemical abundances of our stars, we utilize a line list from
linemake111Available at https://github.com/vmplacco/linemake (Placco et al.,
2021) that contains absorption lines of various elements. We measure the
strength of clean spectral lines by fitting Gaussian profiles to each
absorption feature, after which we obtain their equivalent widths (EWs). Only
lines between $\qty{3500}{}$ and $\qty{7000}{}$ were used. Below
$\qty{4000}{}$ the $S/N$ becomes lower, making good line measurements
challenging. To ensure that EWs are on the linear part of the curve of growth,
only lines with reduced EWs less than $-4.5$ were used. Some exceptions were
made for elements with few lines, such as Na, but line abundances were always
individually checked for consistency.

EW measurements are not useful for spectral features that are blended, lines
with hyperfine structure, and molecular bands (e.g., CH). In those cases, we
used the spectrum synthesis approach to obtain the line abundance where we fit
a synthetic spectrum with a known abundance to the observed spectrum. Spectrum
synthesis was performed for features of C, Si, Al, Sc, Mn, Co, Sr, Ba, and Eu.
We also determined $3\sigma$ upper limits on the abundances of any pertinent,
undetected lines. Table 2 lists all our measured EWs (where available) and
their corresponding abundances.

###  3.2 Stellar Parameters

We determine photometric stellar parameters using parallaxes and proper motion
data from Gaia (where available) and our line measurements of Fe i and Fe ii
lines. First, we determine the effective temperature ($T_{\text{eff}}$) using
the $V-K$ color-temperature-metallicity formula presented in Casagrande et al.
(2010). Using the United States Naval Observatory UCAC4 Catalog222Accessible
via NASA/IPAC Infrared Science Archive (Zacharias et al., 2013), we utilize J
and K magnitudes from 2MASS (Two Micron All-Sky Survey) and V magnitudes from
APASS (AAVSO Photometric All-Sky Survey). To correct for dust reddening, we
use the extinction ($A_{V}$) and reddening $E(B-V)$ values from the dust maps
by Schlafly & Finkbeiner (2011). Values are listed in Table 3. As necessary,
we iterate on the temperature determination with our simultaneous metallicity
determination.

We calculate the surface gravity ($\log g$) by minimizing the trend between Fe
i and Fe ii lines. The values are listed in Table 4. We then test the
robustness of our results by re-deriving $\log g$ using a different method,
specifically the one used in Mardini et al. (2019). This method uses
dereddened $g$ magnitudes from APASS (compiled on the UCAC4 catalog),
parallaxes $\pi$ from Gaia DR3 (Gaia Collaboration et al., 2022), and
fundamental equations with respect to the Sun. The surface gravity values
derived with both methods agree within $1\sigma$.

Generally, our stellar parameters agree well with those obtained by other
studies (e.g., Jacobson et al. 2015). For one star, HE 1310$-$0536, our
$T_{\text{eff}}$ of $\sim$4700 K is lower than that derived by Hansen et al.
(2014) ($T_{\text{eff}}$ = 5000 K). However, a visual comparison of this
star’s H $\beta$ absorption line profile with those of stars with known
temperatures — namely CS22892-52 ($T_{\text{eff}}$ = 4620 K) and HE
1523$-$0901 ($T_{\text{eff}}$ = 4370 K) — well supports $T_{\text{eff}}$ <
5000 K.

For our abundance calculations, we adopt 1D plane-parallel model atmospheres
with $\alpha$-enhancement (generated using the values listed in Table 4) from
Castelli & Kurucz (2004). We used the latest version of the MOOG radiative
transfer code (Sneden, 1973; Sobeck et al., 2011) and assumed local
thermodynamic equilibrium (LTE). This MOOG version includes scattering
effects. We then obtain the Fe abundance for each star using their respective
effective temperatures and surface gravities, as determined above and
including any iterations.

In the process of determining the Fe abundances, we also derived the
microturbulence ($v_{\text{mic}}$) for each star by forcing no trend between
Fe i line abundance and reduced equivalent widths. All our abundance ratios
([X/Fe] and [X/H]) are calculated relative to the solar abundances determined
by Asplund et al. (2009). Our final and adopted stellar parameters are given
in Table 4.

Table 4: Stellar parameters Star | $T_{\text{eff}}$ | $\log g$ | [Fe/H] | $v_{\text{mic}}$
---|---|---|---|---
| [K] | [cgs] | [dex] | [km s^-1]
HE 0104$-$5300 | 4791 | 1.82 | $-3.48\pm 0.15$ | 1.95
HE 1310$-$0536 | 4668 | 0.57 | $-4.22\pm 0.14$ | 1.84
HE 2155$-$2043 | 4994 | 1.82 | $-3.40\pm 0.17$ | 2.02
HE 2303$-$5756 | 6349 | 3.60 | $-2.93\pm 0.10$ | 1.53
HE 2319$-$5228 | 4836 | 2.09 | $-3.38\pm 0.12$ | 1.70
HE 2340$-$6036 | 4689 | 1.93 | $-3.59\pm 0.17$ | 1.91

##  4 Chemical Abundance Results

Table 5: Chemical abundances of our program stars Species | $N$ | $\log\epsilon(X)$ | stderr | [X/H] | [X/Fe] |  Species | $N$ | $\log\epsilon(X)$ | stderr | [X/H] | [X/Fe]
---|---|---|---|---|---|---|---|---|---|---|---
HE 0104$-$5300 | HE 1310$-$0536
C (CH) | 2 | 4.77 | 0.10 |  $-$3.66 |  $-$0.17 |  C (CH) | 2 | 6.25 | 0.10 |  $-$2.18 | 2.04
C (CH)corr | … | … | … | … |  $-$0.11 |  C (CH)corr | … | … | … | … | 2.67
Na i | 2 | 2.94 | 0.13 |  $-$3.30 | 0.18 |  Na i | 2 | 2.17 | 0.05 |  $-$4.07 | 0.15
Mg i | 5 | 4.50 | 0.05 |  $-$3.10 | 0.38 |  Mg i | 6 | 3.84 | 0.06 |  $-$3.76 | 0.46
Al i | 2 | 2.37 | 0.42 |  $-$4.08 |  $-$0.60 |  Al i | 2 | 1.70 | 0.51 |  $-$4.75 |  $-$0.53
Si i | 2 | 4.67 | 0.16 |  $-$2.84 | 0.64 |  Si i | 1 | 3.68 | 0.15 |  $-$3.83 | 0.39
Ca i | 14 | 3.22 | 0.05 |  $-$3.12 | 0.36 |  Ca i | 6 | 2.78 | 0.05 |  $-$3.56 | 0.66
Sc ii | 6 |  $-$0.29 | 0.05 |  $-$3.44 | 0.05 |  Sc ii | 1 |  $-$1.19 | 0.15 |  $-$4.34 |  $-$0.12
Ti i | 12 | 1.74 | 0.05 |  $-$3.21 | 0.27 |  Ti i | 3 | 1.19 | 0.24 |  $-$3.76 | 0.47
Ti ii | 21 | 1.91 | 0.05 |  $-$3.04 | 0.45 |  Ti ii | 14 | 1.28 | 0.05 |  $-$3.67 | 0.55
Cr i | 5 | 1.89 | 0.13 |  $-$3.75 |  $-$0.26 |  Cr i | 1 | 1.07 | 0.15 |  $-$4.57 |  $-$0.35
Mn i | 1 | 1.07 | 0.15 |  $-$4.36 |  $-$0.88 |  Mn i | 1 | 0.36 | 0.15 |  $-$5.07 |  $-$0.85
Fe i | 116 | 4.02 | 0.05 |  $-$3.48 | 0.00 |  Fe i | 43 | 3.28 | 0.10 |  $-$4.22 | 0.00
Fe ii | 13 | 4.27 | 0.05 |  $-$3.23 | 0.25 |  Fe ii | 2 | 3.25 | 0.33 |  $-$4.25 |  $-$0.03
Co i | 3 | 1.82 | 0.05 |  $-$3.17 | 0.31 |  Co i | 2 | 0.86 | 0.05 |  $-$4.13 | 0.09
Ni i | 6 | 2.92 | 0.12 |  $-$3.30 | 0.18 |  Ni i | 6 | 1.95 | 0.07 |  $-$4.27 |  $-$0.05
Zn i | 1 | 1.64 | 0.15 |  $-$2.92 | 0.56 |  Zn i | 1 |  $<$ 0.57 | … |  $<-$3.98 |  $<$ 0.24
Sr ii | 2 |  $-$1.02 | 0.05 |  $-$3.89 |  $-$0.40 |  Sr ii | 1 |  $-$2.99 | 0.15 |  $-$5.86 |  $-$1.64
Ba ii | 2 |  $-$2.71 | 0.05 |  $-$4.89 |  $-$1.41 |  Ba ii | 2 |  $-$2.92 | 0.15 |  $-$5.10 |  $-$0.88
Eu ii | 1 |  $<-$2.62 | … |  $<-$3.14 |  $<$ 0.34 |  Eu ii | 1 |  $<-$3.20 | … |  $<-$3.72 |  $<$ 0.50
HE 2155$-$2043 | HE 2303$-$5756
C (CH) | 2 | 6.00 | 0.10 |  $-$2.43 | 0.97 |  C (CH) | 2 | 6.11 | 0.35 |  $-$2.32 | 0.61
C (CH)corr | … | … | … | … | 1.07 |  C (CH)corr | … | … | … | … | 0.61
Na i | 2 | 3.37 | 0.21 |  $-$2.87 | 0.54 |  Na i | 2 | 3.00 | 0.09 |  $-$3.24 |  $-$0.31
Mg i | 8 | 4.72 | 0.05 |  $-$2.88 | 0.53 |  Mg i | 7 | 4.82 | 0.05 |  $-$2.78 | 0.15
Al i | 1 | 2.29 | 0.10 |  $-$4.16 |  $-$0.76 |  Al i | 2 | 2.74 | 0.05 |  $-$3.71 |  $-$0.79
Si i | 1 | 4.06 | 0.20 |  $-$3.45 |  $-$0.04 |  Si i | 1 | 4.67 | 0.15 |  $-$2.84 | 0.09
Ca i | 11 | 3.32 | 0.05 |  $-$3.02 | 0.38 |  Ca i | 9 | 3.76 | 0.05 |  $-$2.58 | 0.34
Sc ii | 5 |  $-$0.22 | 0.05 |  $-$3.37 | 0.03 |  Sc ii | 4 | 0.29 | 0.05 |  $-$2.86 | 0.06
Ti i | 8 | 1.93 | 0.08 |  $-$3.02 | 0.38 |  Ti i | 3 | 2.62 | 0.05 |  $-$2.33 | 0.60
Ti ii | 17 | 1.70 | 0.05 |  $-$3.25 | 0.15 |  Ti ii | 14 | 2.23 | 0.05 |  $-$2.72 | 0.21
Cr i | 3 | 1.93 | 0.05 |  $-$3.71 |  $-$0.30 |  Cr i | 4 | 2.62 | 0.05 |  $-$3.02 |  $-$0.09
Mn i | 1 | 1.29 | 0.20 |  $-$4.14 |  $-$0.73 |  Mn i | 1 | 1.94 | 0.15 |  $-$3.49 |  $-$0.57
Fe i | 110 | 4.10 | 0.05 |  $-$3.40 | 0.00 |  Fe i | 67 | 4.57 | 0.05 |  $-$2.93 | 0.00
Fe ii | 10 | 4.18 | 0.05 |  $-$3.32 | 0.08 |  Fe ii | 5 | 4.53 | 0.05 |  $-$2.97 |  $-$0.04
Co i | 7 | 2.05 | 0.05 |  $-$2.94 | 0.47 |  Co i | 6 | 2.47 | 0.05 |  $-$2.52 | 0.41
Ni i | 11 | 2.92 | 0.05 |  $-$3.30 | 0.10 |  Ni i | 13 | 3.52 | 0.05 |  $-$2.70 | 0.23
Zn i | 2 | 1.77 | 0.21 |  $-$2.79 | 0.61 |  Zn i | 1 |  $<$ 2.16 | … |  $<-$2.40 |  $<$ 0.53
Sr ii | 2 |  $-$0.76 | 0.06 |  $-$3.63 |  $-$0.23 |  Sr ii | 2 |  $-$0.47 | 0.05 |  $-$3.34 |  $-$0.41
Ba ii | 1 |  $-$2.38 | 0.08 |  $-$4.56 |  $-$1.16 |  Ba ii | 1 |  $<-$1.75 | … |  $<-$3.93 |  $<-$1.00
Eu ii | 1 |  $<-$2.50 | … |  $<-$3.02 |  $<$ 0.38 |  Eu ii | 1 |  $<-$1.42 | … |  $<-$1.94 |  $<$ 0.99
HE 2319$-$5228 | HE 2340$-$6036
C (CH) | 2 | 6.27 | 0.10 |  $-$2.16 | 1.21 |  C (CH) | 2 | 4.27 | 0.15 |  $-$4.16 |  $-$0.56
C (CH)corr | … | … | … | … | 1.22 |  C (CH)corr | … | … | … | … |  $-$0.55
Na i | 4 | 4.64 | 0.07 |  $-$1.60 | 1.78 |  Na i | 2 | 3.21 | 0.05 |  $-$3.03 | 0.56
Mg i | 3 | 4.87 | 0.10 |  $-$2.73 | 0.65 |  Mg i | 5 | 4.56 | 0.05 |  $-$3.04 | 0.55
Al i | 1 | 2.82 | 0.20 |  $-$3.63 |  $-$0.25 |  Al i | 2 | 2.32 | 0.05 |  $-$4.13 |  $-$0.53
Si i | 2 | 5.07 | 0.05 |  $-$2.44 | 0.94 |  Si i | 2 | 4.79 | 0.23 |  $-$2.72 | 0.88
Ca i | 17 | 3.96 | 0.05 |  $-$2.38 | 0.99 |  Ca i | 13 | 3.40 | 0.05 |  $-$2.94 | 0.65
Sc ii | 6 | 0.06 | 0.05 |  $-$3.09 | 0.29 |  Sc ii | 6 |  $-$0.21 | 0.05 |  $-$3.36 | 0.24
Ti i | 7 | 1.99 | 0.05 |  $-$2.96 | 0.42 |  Ti i | 9 | 1.58 | 0.05 |  $-$3.37 | 0.22
Ti ii | 17 | 2.34 | 0.05 |  $-$2.61 | 0.77 |  Ti ii | 21 | 1.95 | 0.05 |  $-$3.00 | 0.59
Cr i | 3 | 2.11 | 0.17 |  $-$3.53 |  $-$0.15 |  Cr i | 4 | 1.54 | 0.08 |  $-$4.10 |  $-$0.50
Mn i | 2 | 1.40 | 0.15 |  $-$4.03 |  $-$0.66 |  Mn i | 1 | 0.70 | 0.15 |  $-$4.73 |  $-$1.14
Fe i | 59 | 4.12 | 0.05 |  $-$3.38 | 0.00 |  Fe i | 86 | 3.91 | 0.05 |  $-$3.59 | 0.00
Fe ii | 6 | 4.29 | 0.05 |  $-$3.21 | 0.16 |  Fe ii | 8 | 4.23 | 0.07 |  $-$3.27 | 0.32
Co i | 3 | 1.66 | 0.05 |  $-$3.33 | 0.05 |  Co i | 3 | 1.51 | 0.13 |  $-$3.48 | 0.11
Ni i | 1 | 2.71 | 0.15 |  $-$3.51 |  $-$0.13 |  Ni i | 2 | 2.41 | 0.06 |  $-$3.81 |  $-$0.22
Zn i | 1 | 1.73 | 0.25 |  $-$2.83 | 0.55 |  Zn i | 1 |  $<$ 1.57 | … |  $<-$2.99 |  $<$ 0.60
Sr ii | 1 |  $-$2.31 | 0.15 |  $-$5.18 |  $-$1.81 |  Sr ii | 2 |  $-$2.13 | 0.15 |  $-$5.00 |  $-$1.41
Ba ii | 1 |  $-$3.20 | 0.15 |  $-$5.38 |  $-$2.01 |  Ba ii | 3 |  $-$2.26 | 0.15 |  $-$4.44 |  $-$0.85
Eu ii | 1 | … | … | … | … |  Eu ii | 1 |  $<-$2.77 | … |  $<-$3.29 |  $<$ 0.31
For Zn, upper limits were determined from the line at 4810 Å. For Eu, they
were derived from the line at 4129 Å.

In Table 5, we present the results of our detailed chemical abundance analysis
of 17 elements in our sample stars. In Figure 2, we show our sample stars in
comparison with a general set of metal-poor halo stars (Cayrel et al., 2004;
Barklem et al., 2005; Yong et al., 2013) and with metal-poor stars in various
UFDs333Compilation by AP Ji available on GitHub (Koch et al., 2008; Simon et
al., 2010; François et al., 2016; Frebel, 2010; Norris et al., 2010; Ji et
al., 2016b; Koch et al., 2013; Gilmore et al., 2013; Ishigaki et al., 2014;
Frebel et al., 2014; Ji et al., 2016a; Roederer et al., 2016; Frebel et al.,
2016; Kirby et al., 2017; Spite et al., 2018; Chiti et al., 2018; Nagasawa et
al., 2018; Ji et al., 2019; Marshall et al., 2019; Ji et al., 2020; Hansen et
al., 2020; Waller et al., 2022; Chiti et al., 2022). In the following, we
discuss notable results and abundance trends. Further discussion of the
origins of our stars and their possible relations to early, small dwarf
galaxies will come in Section 6.

As mentioned in Section 3.2, we derived $\log g$ using two methods, i.e.,
spectroscopically versus using photometry and parallax. To err on the side of
caution, we tested whether changes in stellar parameters affected the chemical
abundances. We found that the abundances were not significantly altered. Since
our abundance measurements are robust to changes in $\log g$, from this point
onward we will only report abundances obtained using the spectroscopically-
derived $\log g$.

Figure 2:  Light element abundances for our program stars (squares and
triangles with black outlines), UFD stars (olive squares; see Section 4 for
references), and halo stars (gray points; Cayrel et al. 2004; Barklem et al.
2005; Yong et al. 2013). The abundances of the stars denoted by triangles (HE
0104$-$5300, HE 2155$-$2043, HE 2303$-$5756) generally follow the abundance
trends found in halo stars. Meanwhile, the stars denoted by squares (HE
1310$-$0536, HE 2340$-$6036, HE 2319$-$5228) show more variation and
deviations from the general halo trend.

###  4.1 Carbon abundances

The most metal-poor halo stars with $\mbox{[Fe/H]}<-2.5$ have long been known
to have predominantly very large carbon abundances (e.g., Frebel et al. 2005;
Placco et al. 2014b; Hansen et al. 2016a; Lee et al. 2017; Yoon et al. 2018).
In general, carbon enhancement is defined as $\mbox{[C/Fe]}>0.7$ (Aoki et al.,
2007). As stars ascend the giant branch, carbon abundances get depleted. We
correct for this fact using models provided in Placco et al. (2014b). Our
sample shows a large variety of [C/Fe], as measured from the carbon G-band
head around 4313 Å and the feature at 4323 Å (see Figure 1). We used spectrum
synthesis on each spectral feature to infer the star’s carbon abundance, after
which we averaged the measurements on both features to obtain the star’s final
carbon abundance.

Three of our program stars, HE 1310$-$0536, HE 2155$-$2043, and HE 2319$-$5228
can thus be classified as carbon-enhanced metal-poor (CEMP) stars. HE
1310$-$0536, with the lowest [Fe/H] of $-$4.22, strongly follows the halo
trend having [C/Fe] = 2.7 (corrected). The other two stars have corrected
abundances around [C/Fe] $\sim$ 1.2 and also support the halo trend. In
contrast, two of our sample stars have very low, subsolar carbon abundances
(HE 2340$-$6036 and HE 0104$-$5300) at [Fe/H] $\sim-$3.5.

###  4.2 Light and iron-peak element abundance trends

In general, our stars’ chemical abundances agree well with the overall
abundance trends of halo and dwarf galaxy stars. However, it is interesting to
note that we do see a significant spread among our six stars. We recall in
this context that these stars were selected only visually for their having a
weak or no Ba line in their spectrum. This selection alone might imply a
certain level of heterogeneity among light and iron-peak elements as measured
in our sample but it may not warrant it either. In Figure 2, it can be seen
that for several elements (Al, Ca, Ti, Co, Ni), the stars denoted by squares
(HE 1310$-$0536, HE 2340$-$6036, HE 2319$-$5228) differ somewhat from the
stars denoted by triangles (HE 0104$-$5300, HE 2155$-$2043, HE 2303$-$5756).
We chose to group these stars by shape according to their behaviour in the
[Sr/Ba] versus [Ba/Fe] plot (see Figure 3 and Section 4.3). Stars denoted by
squares behave like UFD stars in Figure 3 while those denoted by triangles
agree with typical halo star [Sr/Ba] behavior.

Generally, the stars denoted by triangles reasonably follow the overall halo
trends in their light element abundances. Meanwhile, the stars denoted by
squares tend to deviate from them. For Zn, we only have two measurements, so
no strong conclusions can be derived. However, the upper limit of HE
1310$-$0536, our most metal-poor star, is significantly lower than the halo
trend. We thus add Zn to the list of elements for which the two star groups
show different behavior.

###  4.3 Heavy element abundance trends

Sr was detected in all six of our sample stars. All have [Sr/H] $<-$3, and
three have a remarkably low abundance of [Sr/H] $\lesssim-$5\. Ba abundances
were determined for all stars except HE 2303$-$5756, for which only an upper
limit of [Ba/H] $<-$3.93 was found. The other five stars all had [Ba/H]
$<-$4.0, with three having [Ba/H] $<-$5.0. For [Sr/H] and [Ba/H], all stars
follow the general halo trend.

To classify our stars and learn about the origin(s) of their neutron-capture
elements, we plot [Sr/Ba] versus [Ba/Fe] in Figure 3. This offers a way to
empirically learn about the nucleosynthesis processes that may have produced
these elements, as it remains unclear what underlying nucleosynthesis
processes cause the observed differences between typical halo stars and stars
in UFDs. As has been shown in Frebel et al. (2014) and Ji et al. (2019), in
this type of plot, the UFD stars are removed from the halo stars that form the
main branch. Instead, they form their own branch at lower [Sr/Ba] values. This
clustering is immediately evident in Figure 3.

We find that three of our stars, HE 2303$-$5756, HE 2155$-$2043, and HE
0104$-$5300, are located along the main halo branch. Meanwhile, the three
other stars, HE 2319$-$5228, HE 1310$-$0536, and HE 2340$-$6036, are squarely
within the region occupied by the UFD stars. This telling distinction between
the two halves of our sample stars, as shown in Figure 3, suggests that three
of our stars are similar to stars in UFDs (hence the square symbols) while the
other three are similar to typical halo stars (hence the triangle symbols).

In line with our broader goal of building a set of metal-poor halo stars that
possibly originated from small satellite galaxies, we searched the literature
for objects similar to our UFD star-like targets (square symbols in Figure 3).
Using JINAbase (Abohalima & Frebel, 2018), we specifically looked for halo
stars known to have low Sr abundances, specifically $\mbox{[Sr/H]}<-4.5$. We
found 61 of them (detailed in Table 9), which we added as orange diamonds to
the bottom panel of Figure 3 (Li et al., 2015; Barklem et al., 2005; Roederer
et al., 2014; Cayrel et al., 2004; Hansen et al., 2015; Lai et al., 2008;
Placco et al., 2014a; Hollek et al., 2011; Jacobson et al., 2015; Cohen et
al., 2013; Norris et al., 2007; Bonifacio et al., 2009; Ryan et al., 1996;
Collet et al., 2006; Yong et al., 2013; Keller et al., 2014; Frebel et al.,
2015; Caffau et al., 2011; Cohen et al., 2008; Frebel et al., 2007; Mardini et
al., 2024). These panels are identical except for the addition of literature
stars on the bottom panel. We opted to show these two datasets separately for
clarity. Further discussion on the literature stars is presented in Section 6.

We then considered the overall abundance trends of Sr and Ba in various stars.
We chose to focus on [Sr/H] and [Ba/H] values in our plots as these abundance
ratios provide nucleosynthetically decoupled axes (as opposed to [Sr/Fe] and
[Ba/Fe]). For completeness, we nevertheless show [Sr/Fe] and [Ba/Fe] values.
In Figure 4, we show our stars’ Sr and Ba abundances, in comparison with a
general set of metal-poor halo stars (gray points; same set as in Figure 2),
with metal-poor stars in various UFDs (olive squares; same set as in Figure
2), and with the halo stars known to have low Sr abundances (orange diamonds;
same set as in Figure 3).

In Figure 4, we keep the same shape grouping as derived in Figure 3 and
explained in Section 4.3. For [Sr/H], the square-denoted stars are clustered
away from the triangle-denoted stars, lying at the bottom end of the overall
trend set by halo stars. They also have slightly lower [Fe/H] values than the
triangle-denoted stars. However, for [Ba/H], there is no separation between
the two groups. They all cluster towards the lowest Ba abundances as set by
the halo stars. As noted briefly in Section 4.2, the shape grouping also seems
to correlate with light element abundances. We discuss this behavior further
in Section 6.

We note for completeness that we detected no Eu lines in any of our stars but
we provide upper limits in Table 5.

Figure 3: Top: [Sr/Ba] vs [Ba/Fe] for our sample stars in comparison with halo
stars (gray points) and UFD stars (olive squares). The diagonal line guides
the eye and separates halo stars from UFD stars. The UFD stars towards the
bottom right corner that appear to extend the main halo trend are the highly
$r$-process enhanced stars in Reticulum II (Ji et al., 2016a) as this is the
region occupied by strongly enhanced $r$-process stars. Bottom: Same as above
but with additional halo stars (orange diamonds) that have been selected with
a Sr cut of [Sr/H] $<-4.5$. They appear to overlap with the region covered by
UFD stars. Figure 4: Sr and Ba abundances for our sample stars in comparison
to halo stars (gray points), UFD stars (olive squares), and low Sr stars in
the halo (orange diamonds). The program stars denoted by square symbols are
those with [Sr/Ba] and [Ba/Fe] ratios similar to those of UFD stars.
Meanwhile, the program stars denoted by triangles have halo-like [Sr/Ba] and
[Ba/Fe] ratios. The pink-shaded regions on the bottom panels guide the eye to
the lowest neutron-capture abundances (see Section 6.1 for discussion).

###  4.4 Abundance uncertainties

From our measurements of individual spectral lines for each element and the
corresponding standard deviations $\sigma$, we calculated standard errors for
each of our abundance measurements. However, for elements with a small number
of lines ($2<N<30$), the resulting standard error will be underestimated. We
thus applied small $N$ corrections to arrive at our final values. We list our
standard errors in Table 5. In several cases, however, the standard error was
unrealistically small ($<0.05$ dex). For these cases, we thus adopt a minimum
nominal abundance uncertainty of 0.05 dex, a figure informed by the $S/N$ and
other typical uncertainties. For elements with only a single line detection,
we estimated an uncertainty between $0.10-0.15$ dex depending on the quality
of the local spectrum and the $S/N$. Generally, our uncertainties range
between $0.05-0.30$ dex, reflecting the relatively good data quality. For
elements without any line detections, we estimated their abundances with a
3$\sigma$ upper limit, without including any uncertainty.

To determine the systematic uncertainties in our measurements, we varied each
stellar parameter by its uncertainty to assess how the abundance changes in
response to changes in stellar parameter values. We varied the effective
temperature by $100$ K, $\log g$ by $0.3$ dex, and the microturbulence by
$0.3$ dex. We then recorded the change in abundance for each element. Our
uncertainties are listed in Table 6 for HE 2155$-$2043 and HE 2303$-$5756, two
example stars that illustrate this effect. The total uncertainties were found
by adding all of the previously measured uncertainties in quadrature.
Generally, the total uncertainties range from 0.10 to 0.30 dex.

###  4.5 Abundance comparison with literature values

In Table 7, we compare our abundance results with those from previous works
that already analyzed our program stars. For HE 2155$-$2043, we compare our
measurements with those from Purandardas & Goswami (2021) as well as the
handful of measurements provided in Hansen et al. (2016b). For HE 0104$-$5300,
we compare our results with those from Jacobson et al. (2015).

The manner in which each paper derived surface gravity varied, thus leading to
differing numbers between this study and the others. This explains, for
instance, the large discrepancy in [Mg/Fe] between our work and Purandardas &
Goswami (2021). There is also a significant difference in [C/Fe], likely
because we normalized the carbon G-band region of our respective spectra
differently. We tested several different normalization schemes but we were not
able to reproduce their reported value of [C/Fe] $\sim 2$. Meanwhile, Hansen
et al. (2016b) reported [C/Fe] for HE 2155$-$2043 and found a value closer to
ours.

As for the comparison with Jacobson et al. (2015), our measurements for HE
0104$-$5300 agree reasonably well. Overall, for both stars, there is a
moderate to good agreement between our results and the literature values at
the $0.2-0.3$ dex level.

Table 6: Example abundance uncertainties for HE 2155$-$2043 and HE 2303$-$5756 Species | Statistical | $\Delta T_{\text{eff}}$ | $\Delta\log(g)$ | $\Delta v_{\text{mic}}$ | Total
---|---|---|---|---|---
| Unc. |  $[+100$ K] |  $[+0.3$ dex] |  $[+0.3$ dex] | Unc.
HE 2155$-$2043
C (CH) | 0.10 |  $-$0.11 |  $-$0.12 | 0.01 | 0.19
Na i | 0.21 | 0.09 |  $-$0.06 |  $-$0.15 | 0.27
Mg i | 0.05 | 0.09 |  $-$0.07 |  $-$0.08 | 0.15
Al i | 0.10 | 0.25 |  $-$0.07 |  $-$0.11 | 0.30
Si i | 0.20 | 0.06 | 0.06 |  $-$0.13 | 0.25
Ca i | 0.05 | 0.07 |  $-$0.03 |  $-$0.01 | 0.09
Sc ii | 0.05 | 0.06 | 0.08 |  $-$0.01 | 0.11
Ti i | 0.08 | 0.12 |  $-$0.02 |  $-$0.01 | 0.15
Ti ii | 0.05 | 0.07 | 0.10 |  $-$0.06 | 0.14
Cr i | 0.05 | 0.12 |  $-$0.02 |  $-$0.03 | 0.13
Mn i | 0.15 | 0.16 | 0.02 |  $-$0.03 | 0.22
Fe i | 0.05 | 0.12 |  $-$0.03 |  $-$0.08 | 0.16
Fe ii | 0.05 | 0.02 | 0.10 |  $-$0.01 | 0.11
Co i | 0.05 | 0.09 |  $-$0.02 |  $-$0.07 | 0.13
Ni i | 0.05 | 0.13 |  $-$0.05 |  $-$0.15 | 0.21
Zn i | 0.21 | 0.07 | 0.05 | 0.01 | 0.23
Sr ii | 0.06 | 0.10 | 0.14 |  $-$0.27 | 0.32
Ba ii | 0.08 | 0.09 | 0.08 | 0.00 | 0.14
HE 2303$-$5756
C (CH) | 0.35 | 0.15 | 0.15 | 0.13 | 0.43
Na i | 0.09 | 0.07 |  $-$0.01 | 0.01 | 0.11
Mg i | 0.05 | 0.05 |  $-$0.03 | 0.00 | 0.08
Al i | 0.05 | 0.02 | 0.00 | 0.01 | 0.05
Si i | 0.15 | 0.10 | 0.00 | 0.00 | 0.18
Ca i | 0.05 | 0.06 |  $-$0.01 | 0.01 | 0.08
Sc ii | 0.05 | 0.05 | 0.09 | 0.02 | 0.12
Ti i | 0.05 | 0.08 |  $-$0.01 | 0.02 | 0.10
Ti ii | 0.05 | 0.05 | 0.08 | 0.00 | 0.11
Cr i | 0.05 | 0.10 |  $-$0.01 | 0.02 | 0.11
Mn i | 0.15 | 0.03 |  $-$0.07 |  $-$0.04 | 0.17
Fe i | 0.05 | 0.08 |  $-$0.01 | 0.01 | 0.10
Fe ii | 0.05 | 0.02 | 0.11 | 0.01 | 0.12
Co i | 0.05 | 0.09 |  $-$0.04 | 0.00 | 0.11
Ni i | 0.05 | 0.13 |  $-$0.05 | 0.02 | 0.15
Sr ii | 0.06 |  $-$0.03 | 0.04 |  $-$0.05 | 0.09
Table 7: Comparisons with results from other studies

HE 2155$-$2043
---
| Purandardas | Hansen | This Study
| et al. 2021 | et al. 2016 |
Species |  $\log\epsilon$(X) | [X/Fe] | [X/Fe] |  $\log\epsilon$(X) | [X/Fe]
C (CH) | 7.05 | 2.05 | 0.7 | 6.00 | 0.97
Na i | 3.79 | 0.98 | … | 3.37 | 0.54
Mg i | 5.74 | 1.57 | … | 4.72 | 0.53
Al i | … | … | … | 2.29 |  $-$0.76
Si i | … | … | … | 4.06 |  $-$0.04
Ca i | 3.35 | 0.44 | … | 3.32 | 0.38
Sc ii |  $-$0.36 |  $-$0.02 | … |  $-$0.22 |  $-$0.03
Ti i | … | … | … | 1.93 | 0.38
Ti ii | 1.67 | 0.21 | … | 1.70 | 0.15
Cr i | 2.11 |  $-$0.10 | … | 1.93 |  $-$0.30
Mn i | 2.04 | 0.04 | … | 1.29 |  $-$0.73
Fe i | 4.07 | 0.00 | 0.00 | 4.10 | 0.00
Fe ii | 4.01 | … | … | 4.18 | 0.08
Co i | 2.46 | 0.90 | … | 2.05 | 0.47
Ni i | 3.09 | 0.30 | … | 2.92 | 0.10
Zn i | … | … | … | 1.77 | 0.61
Sr ii |  $-$0.66 |  $-$0.04 | 0.2 |  $-$0.76 |  $-$0.23
Ba ii |  $-$2.95 |  $-$1.64 | … |  $-$2.38 |  $-$1.16
Eu ii | … | … | … |  $<-$2.50 |  $<$0.38

HE 0104$-$5300
---
| Jacobson | This Study
| et al. 2015 |
Species |  $\log\epsilon$(X) | [X/Fe] |  $\log\epsilon$(X) | [X/Fe]
C (CH) | 4.77 | 0.13 | 4.77 |  $-$0.17
Na i | 2.74 | 0.29 | 2.94 | 0.18
Mg i | 4.37 | 0.57 | 4.50 | 0.38
Al i | 2.06 |  $-$0.60 | 2.37 |  $-$0.60
Si i | 4.65 | 0.93 | 4.67 | 0.64
Ca i | 3.01 | 0.46 | 3.22 | 0.36
Sc ii |  $-$0.83 |  $-$0.19 |  $-$0.29 | 0.05
Ti i | 1.37 | 0.21 | 1.74 | 0.27
Ti ii | 1.41 | 0.25 | 1.91 | 0.45
Cr i | 1.39 |  $-$0.46 | 1.89 |  $-$0.26
Mn i | 1.12 |  $-$0.52 | 1.07 |  $-$0.88
Fe i | 3.71 | 0.00 | 4.02 | 0.00
Fe ii | 3.72 | 0.01 | 4.27 | 0.25
Co i | 1.34 | 0.14 | 1.82 | 0.31
Ni i | 2.30 |  $-$0.12 | 2.92 | 0.18
Zn i | … | … | 1.64 | 0.56
Sr ii |  $-$1.71 |  $-$0.79 |  $-$1.02 |  $-$0.40
Ba ii |  $-$3.25 |  $-$1.64 |  $-$2.71 |  $-$1.41
Eu ii | … | … |  $<-$2.62 |  $<$0.34

##  5 Kinematic Analysis

In addition to chemical abundances, we also use stellar kinematics to
reconstruct each star’s orbital history and learn more about its origin
scenario. We employ the same procedure detailed in Mardini et al. (2022a). The
authors used Gaia astrometric data in combination with a time-varying Galactic
potential modeled in The orient444Available at https://github.com/Mohammad-
Mardini/The-ORIENT. The potential was constructed from snapshots of a large-
scale cosmological simulation of a Milky Way-like halo (see Mardini et al.
2020 for more details).

For this work, we determined distances to our sample stars using data from
Gaia DR3 (Gaia Collaboration et al., 2022). Despite Gaia’s unprecedented
precision, a considerable number of its parallax measurements have large
uncertainties such that distances derived by simply inverting the parallax
($d=\frac{1}{p}$) can be unreliable. To correct this, we adopt the method
described in Mardini et al. (2022b) where we first employ a zero-point
correction using a pseudocolor where available. We then probabilistically
derive a final distance estimate using a space density prior, also in
combination with the zero-point correction.

However, as Luri et al. (2018) recommended, distances should ideally be
derived using a full Bayesian approach. Following this suggestion, we thus
infer a final distance measurement using the more rigorous probabilistic
method described by Bailer-Jones et al. (2018). As such, we adopt an
exponentially decreasing space density prior (SDP; as opposed to an isotropic
prior) to calculate our final distance estimate. We then run a Monte Carlo
simulation generating 10,000 realizations for each star and assuming a normal
distribution centered around the corrected parallax and its standard
deviation. The mean of the resulting posterior distribution is the final
distance for each star. Our final distances are presented in Table 8.

We then calculate the positions and velocities of our sample stars using Gaia
coordinates, proper motion, and radial velocity. We adopt a position of the
Sun of $R_{\odot}=8.178\pm 0.013$ kpc away from the Milky Way’s center
(GRAVITY Collaboration et al., 2019) at a vertical distance $z_{\odot}=20.8\pm
0.3$ pc above the Galactic plane and solar velocities as given by Bennett &
Bovy (2019) and Schönrich et al. (2010). We adopt $V_{\text{LSR}}=220$ km.s^-1
as measured by Kerr & Lynden-Bell (1986). Table 8 summarizes the kinematic
parameters of our sample stars.

The strong retrograde $V$ motions of all of our stars (i.e., $V<0$),
especially those of HE 2303$-$5756 and HE 2155$-$2043 with $V\sim-450$
km.s^-1, suggest that our sample is indeed composed of halo stars, a finding
which aligns with their extremely metal-poor nature. These results suggest a
possible early accretion origin of our stars. We also tested using
$V_{\text{LSR}}=232$ km.s^-1 (McMillan, 2017) and $V_{\text{LSR}}=238$ km.s^-1
(Schönrich, 2012) to assess the impact of different $V_{\text{LSR}}$ values on
our conclusions regarding the origins of our stars. Using these new values
changes $V$ by only 10 km.s^-1. Thus, changing $V_{\text{LSR}}$ does not
affect our conclusion regarding the retrograde motion of our stars.

Given all these inputs, we derive the Galactocentric coordinates $(X,Y,Z)$
with a right-handed frame and rectangular Galactic velocities $(U,V,W)$. With
these kinematic parameters, we then back-integrate each star’s orbit with a
time-varying Galactic potential, namely orient potential #483868 (Mardini et
al., 2022a). To investigate their orbits statistically, we generate 1,000
realizations of each star’s orbital evolution. In Figure 5, we show one such
realization for our UFD star-like targets, namely HE 1310$-$0536, HE
2340$-$6036, and HE 2319$-$5228 (square symbols in relevant figures).

Of particular note is each star’s $Z_{\text{max}}$, its maximum distance above
or below the Galactic plane. On the right column of Figure 5, we show the
distribution of $Z_{\text{max}}$ values over 1,000 realizations. In Table 8,
we list our program stars’ median $Z_{\text{max}}$ values. All of our program
stars have significant $Z_{\text{max}}$, generally ranging between $7-48$ kpc,
once again confirming that these are halo stars. Many of them also have rather
elliptical orbits about the Galactic Center.

For HE 1310$-$0536, we were unable to derive a reliable distance due to large
uncertainties in the Gaia data ($d=11.9\pm 2.3$ kpc). To err on the side of
caution, we thus carried out our kinematic analysis (described above) of HE
1310$-$0536 three times, with three possible distances taking into account its
uncertainty (9, 12, and 14 kpc). This was to test the effects of the distance
uncertainty on our orbital history results. We find that in all three cases,
the star appears to be part of the halo, rather than the disk.

Table 8: Kinematic parameters of our sample stars Star | $X$ | $Y$ | $Z$ | $U$ | $V$ | $W$ | Distance | $Z_{\text{max}}$
---|---|---|---|---|---|---|---|---
| [kpc] | [kpc] | [kpc] | [km.s^-1] | [km.s^-1] | [km.s^-1] | [kpc] | [kpc]
HE 0104$-$5300 | $-6.73$ | $-2.81$ | $-6.54$ | $-177.98$ | $-303.15$ | $-117.25$ | $7.4\pm 0.6$ | 31
HE 1310$-$0536 | $-3.88$ | $-4.60$ | $9.51$ | $-137.98$ | $-263.95$ | $64.35$ | $11.9\pm 2.3$ | 14
HE 2155$-$2043 | $-4.11$ | $2.64$ | $-5.81$ | $-216.27$ | $-467.79$ | $-245.61$ | $6.4\pm 1.2$ | 48
HE 2303$-$5756 | $-7.65$ | $-0.34$ | $-0.86$ | $-213.62$ | $-430.29$ | $155.35$ | $1.1\pm 0.0$ | 12
HE 2319$-$5228 | $-5.97$ | $-1.25$ | $-4.37$ | $57.12$ | $-190.84$ | $-245.15$ | $5.3\pm 0.4$ | 11
HE 2340$-$6036 | $-6.10$ | $-1.92$ | $-4.02$ | $124.93$ | $-281.53$ | $-60.58$ | $5.6\pm 0.4$ | 7
Figure 5: Orbital evolution over the last 8 billion years of our three UFD
star-like targets. For every star, we show one realization of its orbital
history, back-integrated using the time-dependent orient potential #483868.
Left: View of the Galactic plane from above in X-Y coordinates (in kpc).
Center: Side view showing the Z height of the orbit as a function of distance
from the Galactic Center. Right: Histogram showing the distribution of
$Z_{\text{max}}$ (maximum distance above or below the Galactic plane) over
1000 realizations of the star’s orbital evolution. Figure 6: Abundances of HE
2319$-$5228 as a function of atomic number. Overlaid is a Starfit supernova
nucleosynthesis model (teal line) that best matches the data (red points). Sc
is not considered in the fit given uncertainties in the Sc production, as
recommended by Starfit. Overall, the fit is moderate given that Cr and Mn are
significantly overproduced by the model. It corresponds to a progenitor with
$\sim$27 solar masses and an explosion energy of 0.9 Bethe. Applying exemplary
NLTE corrections to multiple elements did not significantly change the fitting
outcome.

##  6 On the origins of halo stars with low neutron-capture element abundances

###  6.1 Small Accreted Stellar System (SASS) stars

In Figure 2, we showed that our six sample stars — all metal-poor and
deficient in neutron-capture elements — generally follow the light element
abundance trends set by typical halo stars. In Section 5, we also showed that
all our stars have halo kinematics with a strong propensity for significantly
retrograde orbits. Yet, when considering their [Sr/Ba] vs. [Ba/Fe] behavior
(as shown in Figure 3), the sample splits into two distinct groups. Three
stars fall in the region typically occupied by UFD stars, while the other
three stars overlap with the main halo star branch. This division persists
noticeably in the [Sr/H] vs. [Fe/H] panel in Figure 4. To some degree, this
division is also seen in some of the light elements, particularly in Ca, Co,
and Ni.

Generally, stars in UFDs have a strong tendency to cover regions of the lowest
[Sr/H] and [Ba/H] values, i.e., [Sr/H] $<-4.0$ and [Ba/H] $<-4.0$ as shown in
Frebel et al. (2014). To arrive at more stringent results, we restrict the
following discussion to an even lower neutron-capture abundance cutoff at
[Sr/H] $<-4.5$ and [Ba/H] $<-4.0$. These low Sr and Ba regions are colored
pink in Figure 4 to guide the eye.

Indeed, half of our sample stars (denoted with squares for their UFD-like
[Sr/Ba] vs. [Ba/Fe] behavior) fall in the pink low-Sr region in the [Sr/H]
versus [Fe/H] panel of Figure 4. Meanwhile, the other three stars (denoted
with triangles) do not. For the [Ba/H] versus [Fe/H] panel, however, all
sample stars more or less cluster in the pink low-Ba region typically covered
by UFD stars. All these stars are extremely metal-poor.

Altogether, the three characteristics of (1) low [Fe/H] values paired with a
corresponding [Sr/Ba] value (as in Figure 3), (2) low [Sr/H] values, and (3)
retrograde halo orbits suggest that three of our stars have origins extremely
similar to those of typical UFD stars. This similarity implies that these
stars (HE 2319$-$5228, HE 2340$-$6036, HE 1310$-$0536) may have originated
from small dwarf galaxies long accreted by the Milky Way. On the other hand,
the other three stars (HE 0104$-$5300, HE 2155$-$2043, HE 2303$-$5756) may
have different origin scenarios. Perhaps they came from larger dwarf galaxies
(analogous to the classical dwarf spheroidal galaxies) which underwent some
amount of chemical evolution before eventually being accreted by the Galaxy.
Their Sr and Ba behavior resembles that of several extant satellite systems
(e.g., Venn et al. 2004).

Informed by the above arguments and by observations of stars in UFDs, we can
now assume that stars with the lowest Fe and Sr values originated in some of
the smallest, earliest systems that formed in the universe. We will refer to
these stars (the low-Sr stars represented as squares in the various figures)
as stars from Small Accreted Stellar Systems (SASS). Now we will go one step
further and assume that if a star in the Galactic halo has an extraordinarily
low Sr abundance, then it must also be a SASS star. As will be discussed
later, we indeed find many such SASS stars in the literature. Together, these
stars open up new ways of studying early star formation environments, the
early assembly history of the Milky Way, and the population of the earliest
galaxies.

In the following subsections (6.2, 6.3, 6.4, 6.5), we discuss the environments
of early dwarf galaxies and establish that SASS stars indeed originated from
such environments. In later subsections (6.6, 6.7), we discuss SASS stars in
the literature and the broader potential of SASS stars as probes of the early
universe.

###  6.2 Light element abundances

As noted in Subsection 4.2, SASS stars tend to show larger deviations from the
well-established halo trends in their light element abundances — a behavior
found repeatedly among the most iron-poor stars. These iron-poor stars are
thought to have formed from gas that was enriched by individual and/or unusual
supernova progenitors and/or gas that underwent inhomogeneous mixing
(Christlieb et al., 2004; Frebel et al., 2005; Placco et al., 2015; Frebel &
Norris, 2015). This behavior offers principal support for our claim that SASS
stars formed in the earliest systems; they experienced enrichment events that
were individual and more stochastic.

###  6.3 Carbon abundances

Another feature of the most iron-poor stars is strong carbon enhancement. Two
of our three SASS stars are strongly carbon-enhanced, with HE 2319$-$5228
having [C/Fe] $=1.2$ and HE 1310$-$0536 having [C/Fe] $=2.7$. Once again, this
lends further support to our claim that SASS stars came from the earliest
systems. However, our third SASS star, HE 2340$-$6036, has a low carbon
abundance with [C/Fe] $=-0.55$. The origin of such a low value generally
remains unclear, but it could be due to extreme inhomogeneous mixing.

Interestingly, stars in UFDs tend to have only mild carbon enhancement or even
subsolar levels of carbon (Frebel et al., 2016; Norris et al., 2010). Only one
of our SASS stars, HE 2340$-$6036, has UFD-like carbon abundance levels.
Meanwhile, our other two SASS stars are grossly different in that they are
much more carbon-enhanced. We note, however, that several individual UFD stars
are known to have such extremely high carbon abundances (Lai et al., 2011).
More data might help shed light on the overall C behavior in UFDs. In the
meantime, we continue to assume that SASS stars are indeed not unlike the
surviving UFD stars.

###  6.4 Supernova nucleosynthesis yields

We further test the idea of early pre-enrichment by considering supernova
nucleosynthesis yields and how they match with the now-observed abundances in
our SASS stars. We used Starfit555Available at https://starfit.org (Heger &
Woosley, 2010), a fitting tool that matches the observed light element
($Z<30$) abundance signatures with models of Pop III supernova nucleosynthesis
yields. This provides clues to the properties of the (possibly Pop III)
progenitors that operated in the birth environments of our SASS stars.

Studies of the ultra-faint dwarf galaxies Hercules and Leo IV suggest that
they may have been enriched by the supernovae of putative Population III stars
(Koch et al., 2008; Simon et al., 2010). We nominally find Pop III progenitor
masses of 10 to 23 M⊙ and accompanying explosion energies of $\sim$1050 to
1051 ergs. An example fit is shown in Figure 6. As can be seen, the fit for HE
2319$-$5228 is overall of moderate quality. The same is true for the fits of
other stars. These rather low explosion energies are in line with our claim
that SASS stars formed in early systems. Otherwise, significantly larger
energies might have blown the systems apart and prevented further star
formation.

###  6.5 Neutron-capture element abundances

A remaining question pertains to the origins of low neutron-capture element
abundances — a property shared by both UFD and SASS stars. Generally, low
neutron-capture element abundances imply either (1) a limited production of
said elements, or (2) a limited number of enrichment events in combination
with poor mixing and distribution of metals. Possible options for events and
processes include the limited $r$-process operating in supernovae (Wanajo &
Ishimaru, 2006; Arcones & Montes, 2011; Frebel, 2018) and other heavy element
production processes associated with supernovae such as spinstars (Maeder et
al., 2015; Hansen et al., 2016a), all in concert with the commensurate carbon
enhancement events. Crucially, we do not consider neutron star mergers as they
would strongly overproduce neutron-capture elements, likely generating
abundances higher than those observed in our SASS stars. Not only that, we
found upper limits of [Eu/H] $<-$3 for five out of our six program stars. With
such low Eu levels, we can reasonably exclude a rare, prolific $r$-process
event such as a neutron star merger. Our low Eu levels are more in line with
the Eu abundances found in typical UFDs.

Nevertheless, it is interesting to consider whether any SASS stars have a
[Sr/Ba] ratio similar to that produced by either the $s$\- or the $r$-process.
Typical $r$-process stars have [Sr/Ba] $\sim-0.4$ to $-0.7$ (Frebel & Ji,
2023; Ji et al., 2016c) while $s$-process stars have [Sr/Ba] $<-1$. None of
our stars has [Sr/Ba] $<-1$. However, one SASS star has [Sr/Ba] =$-0.87$ (HE
1310$-$0536) and another one has [Sr/Ba] =$-0.56$ (HE 2340$-$6036). These
stars have extremely low levels of [Sr/H] ($-5.86$ for HE 1310$-$0536; $-5.00$
for HE 2340$-$6036) and [Ba/H] ($-5.10$ for HE 1310$-$0536; $-4.44$ for HE
2340$-$6036). This suggests the potential existence of early $r$-process
nucleosynthesis, or at least, a dilution process that led to such low overall
levels. Perhaps this nucleosynthetic event was associated with an early
supernova such as a magnetorotational hypernova (Yong et al., 2021). More
observations of similar stars could help shed light on this question, which we
address further below. For completeness, we note that the third SASS star, HE
2319$-$5228, has [Sr/Ba] = $0.2$. Its neutron-capture elements likely have a
different origin.

###  6.6 SASS stars in the literature

As we established in the previous subsections, our three SASS stars likely
originated in the earliest systems in the universe. A large set of such stars
would therefore allow us to study in detail the early environments of dwarf
galaxies. Motivated by this, we sought to expand our collection of 3 SASS
stars by searching the literature for similar stars. To this end, we queried
JINAbase for stars with low [Fe/H] and [Sr/H] $<-4.5$. We found 61 such stars
which we show as orange diamonds in Figures 3, 4, and 7. These literature SASS
stars are listed in Table 9 together with selected chemical abundances,
JINAbase identification numbers, and radial velocities.

In Figure 4, our literature SASS stars lie within the low-Sr region of the
[Sr/H] versus [Fe/H] plot. This is expected since we selected these stars
precisely for their having [Sr/H] $<-4.5$. Though these stars were only
selected for their low [Sr/H], it is interesting to note that these same stars
strictly have low barium abundances; they all have [Ba/H] $<-4.0$. This makes
SASS stars generally straightforward to identify using neutron-capture element
abundances. However, while a low [Ba/H] value is a necessary characteristic
for a SASS star, this criterion by itself is not sufficient for the selection
of SASS stars. In addition to being poor in Sr (and by extension, Ba), SASS
stars must also have low metallicities, i.e., [Fe/H] $\sim$ $-$5.0 to $-$3.0.
This low-metallicity criterion selects for stars that formed in the early
universe.

Now we consider the reverse logical path. Looking at their [Sr/Ba] ratio in
Figure 3, all these literature SASS stars fall into the UFD region (just as
our three program SASS stars). In Figure 7, we show the light element
abundances of these stars (with abundances as collected from JINAbase). Like
our three program SASS stars, the literature SASS stars have a tendency to be
more scattered in their light element abundances, compared to the typical halo
trend set by stars with higher Sr values. This behavior was foreshadowed by
our program SASS stars and is more apparent now with the literature SASS
stars. We note that as more stars become available, a more detailed analysis
of the overall behavior would deliver more accurate results.

Finally, we carried out the same kinematic analysis detailed in Section 5 on
the literature SASS stars to assess their orbital histories. Like our program
SASS stars, we find a strong tendency for these stars to be on highly
retrograde orbits, with essentially no prograde motion present.

###  6.7 Studying the early universe with SASS stars

We thus conclude that halo stars with [Sr/H] $<-4.5$ can be considered SASS
stars and that they likely originated from early small systems that were
accreted by the Milky Way at the earliest times. After all, the early Milky
Way assembled from galactic building block-type systems much like UFDs. Thus
it should not be surprising that at least the low-mass stars are still present
in our Galaxy and that we can observe them. Accordingly, we suggest using
[Sr/H] $<-$4.5 as an important selection tool in identifying the truly oldest
stars, namely the SASS stars, in our Galaxy. Furthermore, these stars will be
characterized by a suitable UFD-like [Sr/Ba] ratio (together with an
appropriate [Ba/Fe] value), as well as a highly retrograde orbit.

From the chemical abundances of SASS stars, it appears that the earliest
systems accreted by the Milky Way were enriched by a very small number of
supernovae whose (perhaps unique) yields were somewhat inhomogeneously mixed.
This inhomogeneity produced variations in light element abundances, including
carbon abundances, from system to system. Neutron-capture elements were
sparsely produced and/or diluted likely in multiple ways, with $r$-process
nucleosynthesis (at least up to Ba) playing a role. This suggests that most,
if not all, elements were produced in small amounts from the earliest times
onward. High-resolution cosmological simulations of early (baryonic) structure
formation may be able to more fully model the population of the earliest
galaxies and building blocks based on these observational constraints.

Tracing the oldest surviving stars within large galaxies, such as the Milky
Way, would be another avenue towards confirming our observational findings
that there appears to be a significant population of the very oldest stars
present in the halo. Adding the Sr abundance as a criterion in addition to the
commonly used Fe abundance may finally offer a path towards identifying not
just the most metal-poor but indeed the oldest stars in the universe.

We note that SASS stars may not be the only unique survivors of this very
early time. For instance, stars with slightly higher Sr values may well be
part of this extremely ancient population. However, we apply a more stringent
cut here in this work to ensure the selection of a relatively clean and easily
identifiable representative set of this unique population of the very oldest
stars in the universe.

Going forward, SASS stars offer a new path to identifying ancient stars that
are relatively close by. These nearby stars, which are easier to observe than
the faint and distant dwarf galaxy stars, will allow for more detailed probing
of early galactic environments and the galaxy formation era. With many more
halo stars available to probe early star formation, we may eventually be able
to establish the full history of the systems that built the Milky Way. In
addition, we may better understand the conditions of the earliest galaxies
which hosted the oldest and most metal-poor stars.

Table 9:  SASS stars from the literature Name | [Fe/H] | [C/H] | [Mg/H] | [Sr/H] | [Ba/H] | JINAbase ID | RV | Reference
---|---|---|---|---|---|---|---|---
HE 2248$-$3345 |  $-$2.74 |  $-$2.55 |  $-$2.82 |  $-$4.62 |  $-$4.15 | 901 | 145.4 | BAR05
HE 2249$-$1704 |  $-$2.95 |  $-$3.18 |  $-$2.75 |  $-$5.00 |  $-$4.92 | 2604 | $-134.1$ | COH13
CS29502-042 |  $-$3.00 |  $-$2.94 |  $-$2.98 |  $-$5.07 |  $-$4.93 | 842 | $-138.05$ | CAY04
HE 2304$-$4153 |  $-$3.02 |  $-$3.71 |  $-$2.94 |  $-$4.55 |  $-$4.68 | 913 | $-221.7$ | BAR05
HE 0411$-$5725 |  $-$3.08 |  $-$2.61 |  $-$2.75 |  $-$4.58 |  $-$4.38 | 2557 | 286.1 | COH13
HE 0316+0214 |  $-$3.13 |  $-$3.89 |  $-$2.49 |  $-$4.52 |  $-$4.57 | 268 | $-161.0$ | BAR05
BS16084-160 |  $-$3.15 |  $-$3.23 |  $-$3.01 |  $-$5.26 |  $-$5.30 | 638 | $-132.32$ | LAI08
HE 1416$-$1032 |  $-$3.20 |  $<-$3.51 |  $-$3.04 |  $-$4.89 |  $-$4.49 | 2623 | $5.2$ | COH13
CS22943$-$137 |  $-$3.22 | N/A |  $-$2.59 |  $<-$4.62 |  $<-$3.70 | 709 | $-20.0$ | RYA96
CD $-$24°17504 |  $-$3.23 |  $-$1.83 |  $-$2.84 |  $-$4.75 |  $-$4.15 | 3936 | 136.4 | MAR24
LAMOST J0006+1057 |  $-$3.26 |  $-$3.58 |  $-$2.59 |  $-$5.33 |  $-$5.01 | 40 | $-281.6$ | LI15a
SMSS J031556.09$-$473442.1 |  $-$3.26 |  $-$3.07 |  $-$2.78 |  $-$5.24 |  $-$4.14 | 2372 | 192.3 | JAC15
HE 0008$-$3842 |  $-$3.35 |  $-$4.26 |  $-$3.02 |  $-$4.73 |  $-$5.13 | 47 | 128.9 | BAR05
HE 0344$-$0243 |  $-$3.35 |  $-$3.32 |  $-$2.95 |  $-$5.08 |  $-$4.63 | 2647 | $-110.5$ | COH13
SMSS J003327.36$-$491037.9 |  $-$3.36 |  $<-$3.76 |  $-$2.81 |  $-$4.93 |  $-$4.42 | 2351 | 84.7 | JAC15
CS22949-048 |  $-$3.37 |  $-$3.38 |  $-$3.15 |  $-$4.75 |  $-$5.00 | 925 | $-160.8$ | ROE14b
CS30492-016 |  $-$3.40 |  $-$2.65 |  $-$2.62 |  $<-$5.10 |  $<-$3.33 | 1453 | 46.0 | ROE14b
CS30325-094 |  $-$3.40 |  $-$3.21 |  $-$2.94 |  $-$5.44 |  $-$5.23 | 596 | $-138.0$ | CAY04
SMSS J023147.96$-$575341.7 |  $-$3.42 |  $-$3.07 |  $-$3.13 |  $-$5.49 |  $-$5.06 | 1923 | 219.2 | JAC15
HE 0139$-$2826 |  $-$3.46 |  $-$2.98 |  $-$2.90 |  $-$4.97 |  $-$4.68 | 2022 | $-253$ | PLA14a
HE 1347$-$1025 |  $-$3.48 |  $-$3.21 |  $-$3.08 |  $-$4.51 |  $-$4.15 | 2622 | 48.6 | COH13
HE 1356$-$0622 |  $-$3.49 |  $-$3.41 |  $-$2.93 |  $-$5.26 |  $-$4.72 | 2631 | 93.5 | COH13
CS22942-002 |  $-$3.53 |  $-$3.26 |  $-$3.05 |  $-$5.34 |  $-$4.85 | 99 | $-155.0$ | ROE14b
SMSS J085924.06$-$120104.9 |  $-$3.63 |  $-$3.83 |  $-$3.02 |  $-$4.64 |  $-$4.86 | 2387 | $-187.6$ | JAC15
SDSS J1322+0123 |  $-$3.64 |  $-$3.15 |  $-$3.39 |  $-$4.88 |  $-$4.94 | 2263 | 105.0 | PLA15a
LAMOST J0126+0135 |  $-$3.57 |  $-$4.08 |  $-$3.15 |  $-$5.18 |  $-$4.71 | 173 | $-256.0$ | LI15a
SMSS J184825.29$-$305929.7 |  $-$3.65 |  $-$3.40 |  $-$3.15 |  $-$5.17 |  $-$5.23 | 2471 | 110.2 | JAC15
HE 2318$-$1621 |  $-$3.67 |  $-$3.13 |  $-$3.47 |  $-$4.67 |  $-$5.28 | 2080 | $-39.95$ | PLA14a
CS22952$-$015 |  $-$3.68 |  $-$4.52 |  $-$3.57 |  $-$4.52 |  $-$5.37 | 968 | $-7.0$ | ROE14b
HE 0302$-$3417a |  $-$3.70 |  $-$3.22 |  $-$3.15 |  $-$5.05 |  $-$5.80 | 2026 | 121.7 | HOL11
HE 2331$-$7155 |  $-$3.70 |  $-$2.34 |  $-$2.48 |  $-$4.53 |  $-$4.58 | 2082 | 210.6 | HAN15
HE 0056$-$3022 |  $-$3.72 |  $-$3.52 |  $-$3.39 |  $-$4.68 |  $-$5.23 | 1490 | 29.7 | ROE14b
HE 0926$-$0546 |  $-$3.73 |  $<-$3.11 |  $-$3.41 |  $-$4.95 |  $-$4.56 | 2637 | 152.9 | COH13
HE 1116$-$0634 |  $-$3.73 |  $-$3.65 |  $-$2.91 |  $-$5.99 |  $-$5.54 | 2507 | 116.09 | HOL11
HE 0048$-$6408 |  $-$3.75 |  $-$4.03 |  $-$3.35 |  $-$4.57 |  $-$5.23 | 2018 | $-33.99$ | PLA14a
HE 1012$-$1540 |  $-$3.76 |  $-$1.77 |  $-$2.36 |  $-$4.56 |  $-$4.45 | 392 | 225.6 | ROE14b
CS22960$-$048 |  $-$3.78 |  $-$3.44 |  $-$3.19 |  $-$5.91 |  $-$5.50 | 815 | $-85.0$ | ROE14b
SMSS J010651.91$-$524410.5 |  $-$3.79 |  $-$3.66 |  $-$3.23 |  $-$4.58 |  $-$5.43 | 2355 | 189.5 | JAC15
CS30339$-$073 |  $-$3.80 |  $-$3.73 |  $-$3.48 |  $-$5.04 |  $-$5.48 | 1451 | 171.0 | ROE14b
BS16076$-$006 |  $-$3.81 |  $-$3.35 |  $-$3.25 |  $<-$5.30 |  $<-$4.86 | 474 | 203.12 | BON09
SMSS J004037.56$-$515025.2 |  $-$3.83 |  $-$3.92 |  $-$3.24 |  $-$4.82 |  $-$4.84 | 2352 | 72.9 | JAC15
CS22189$-$009 |  $-$3.85 |  $-$3.62 |  $-$3.48 |  $-$4.81 |  $-$5.44 | 222 | $-42.0$ | ROE14b
CS22963$-$004 |  $-$3.85 |  $-$3.58 |  $-$3.51 |  $-$5.04 |  $-$4.70 | 236 | 294.4 | ROE14b
CS22891$-$200 |  $-$3.88 |  $-$3.53 |  $-$3.24 |  $-$5.24 |  $-$4.81 | 671 | 137.7 | ROE14b
HE 1300+0157 |  $-$3.88 |  $-$2.54 |  $-$3.50 |  $<-$5.51 |  $<-$4.74 | 2665 | 73.4 | FRE07a
SMSS J184226.25$-$272602.7 |  $-$3.89 |  $<-$4.18 |  $-$3.27 |  $-$5.66 |  $-$5.16 | 2469 | $-247.0$ | JAC15
HE 0945$-$1435 |  $-$3.90 |  $<-$1.87 |  $-$3.76 |  $<-$4.77 |  $<-$3.87 | 2035 | 121.8 | HAN15
HE 1201$-$1512 |  $-$3.92 |  $-$2.72 |  $-$3.66 |  $<-$4.73 |  $<-$3.81 | 2040 | 237.2 | YON13
SMSS J005953.98$-$594329.9 |  $-$3.94 |  $-$2.73 |  $-$3.32 |  $-$5.04 |  $-$4.58 | 2353 | 375.3 | JAC15
HE 1424$-$0241 |  $-$3.96 |  $<-$3.17 |  $-$3.57 |  $<-$5.62 |  $-$4.92 | 2646 | 59.8 | COH08
CS22172$-$002 |  $-$4.00 |  $-$3.77 |  $-$3.68 |  $-$5.07 |  $-$5.08 | 257 | 255.0 | CAY04
HE 0057$-$5959 |  $-$4.08 |  $-$3.22 |  $-$3.57 |  $-$5.14 |  $-$4.54 | 2019 | 375.3 | YON13
HE 2239$-$5019 |  $-$4.20 |  $<-$2.45 |  $-$3.70 |  $<-$4.75 |  $<-$4.15 | 893 | 368.7 | HAN15
CS22885$-$096 |  $-$4.21 |  $-$3.81 |  $-$3.57 |  $-$6.16 |  $-$6.05 | 701 | $-249.1$ | ROE14b
BD +44∘ 493 |  $-$4.26 |  $-$3.08 |  $-$3.39 |  $-$4.83 |  $-$5.16 | 1192 | $-150.6$ | ROE14b
SDSS J102915.14+172927.9 |  $-$4.73 |  $<-$3.68 |  $-$4.70 |  $<-$5.00 | … | 2554 | -34.5 | CAF11a
CD $-$38∘ 245 |  $-$4.50 |  $<-$4.78 |  $-$3.93 |  $-$5.22 |  $-$5.59 | 1509 | 46.4 | ROE14b
HE 0557$-$4840 |  $-$4.80 |  $-$3.14 |  $-$4.57 |  $<-$5.77 |  $<-$4.73 | 368 | 211.94 | NOR07
SD1313$-$0019 |  $-$5.00 |  $-$2.04 |  $-$4.56 |  $<-$5.28 |  $<-$4.78 | 2262 | 267.0 | FRE15
HE 0107$-$5240 |  $-$5.28 |  $-$1.58 |  $-$5.13 |  $<-$5.80 |  $<-$4.46 | 2021 | 44.78 | CHR04
SMSS J031300.36-670839.3 |  $<-$7.30 |  $-$2.40 |  $-$4.30 |  $<-$6.70 |  $<-$6.10 | 2136 | $-67.93$ | KEL14
HE 2319$-$5228 |  $-$3.38 |  $-$2.16 |  $-$2.73 |  $-$5.18 |  $-$5.38 |  | 292.8 | this study
HE 2340$-$6036 |  $-$3.59 |  $-$4.16 |  $-$3.04 |  $-$5.00 |  $-$4.44 |  | 212.4 | this study
HE 1310$-$0536 |  $-$4.22 |  $-$2.18 |  $-$3.76 |  $-$5.86 |  $-$5.10 |  | 188.3 | this study
References are as follows: LI15a: Li et al. (2015), BAR05: Barklem et al. (2005), ROE14b: Roederer et al. (2014), |
CAY04: Cayrel et al. (2004), HAN15: Hansen et al. (2015), LAI08: Lai et al. (2008), PLA14a: Placco et al. (2014a), |
HOL11: Hollek et al. (2011), JAC15: Jacobson et al. (2015), COH13: Cohen et al. (2013), NOR07: Norris et al. (2007), |
BON09: Bonifacio et al. (2009), RYAN96: Ryan et al. (1996), COL06: Collet et al. (2006), YON13: Yong et al. (2013), |
KEL14: Keller et al. (2014), FRE15: Frebel et al. (2015), CAF11a: Caffau et al. (2011), COH08: Cohen et al. (2008), |
FRE07a: Frebel et al. (2007), MAR24: Mardini et al. (2024) |
Figure 7:  Same as Figure 2 but with halo stars having low [Sr/H] $<-4.5$
added as orange diamonds. While the metallicity is lower for those stars, the
scatter is larger than that for the other halo stars. See Section 6.6 for
discussion.

##  7 Summary

In this paper, we studied a sample of metal-poor stars with low neutron-
capture element abundances to gain insight into the origins of each of our
program stars, and to investigate the potential connections of these halo
stars to those in present-day UFDs and other dwarf galaxies. Our findings can
be broadly summarized as follows:

  1. 1.

High-resolution optical spectra of our target stars were obtained with
Magellan/MIKE. Stellar parameters ($T_{\text{eff}}$, $\log g$, and
$v_{\text{mic}}$) were determined as part of performing a detailed chemical
abundance analysis. Depending on the element, we obtained abundances using
either equivalent widths or spectrum synthesis. We obtained abundances of a
total of 17 light and heavy elements. All our stars were found to have [Fe/H]
between $-$2.9 and $-$4.3.

  2. 2.

The abundances of our six stars generally agree well with halo trends for most
elements. However, for some elements (specifically Al, Ca, Ti, Co, Ni), our
three low-Sr stars ([Sr/H] $<-4.5$; namely HE 1310$-$0536, HE 2340$-$6036, HE
2319$-$5228) show larger scatter and deviations from the main halo trend. In
addition, two of our stars (HE 2340$-$6036 and HE 0104$-$5300) have C
abundances that are far below average.

  3. 3.

The Sr and Ba abundances of our stars are extremely low, with our low-Sr
program stars (square symbols in relevant figures) having values of [Sr/H]
$<-$ 4.5 and [Ba/H] $<-$4.0. In terms of neutron-capture abundances, these
three stars are highly similar to UFD stars (see Figures 3 and 4). The other
three program stars (triangle symbols in relevant figures) are separated in a
distinct cluster in chemical abundance space, following the general halo trend
set by higher-metallicity stars.

Interestingly, this separation does not exist for Ba abundances, with all six
program stars having [Ba/H] $\lesssim$ 4.0, and following the main trend.
Therefore, Sr abundance is a telling sign of whether a star is similar to a
UFD star, and by extension, old.

  4. 4.

The low-Sr program stars (square symbols in relevant figures) possess [Sr/Ba]
ratios that do not reflect the main halo trend, but instead fall into the
regime of UFD stars (see Figure 3). We identify these three stars as having
similar origins to the metal-poor stars observed in UFDs. We label these
stars, and others with similarly low [Sr/H] values, as "SASS" stars. SASS
stars are named as such because they likely originated in the early universe
from Small Accreted Stellar Systems. Since these SASS stars are nearer to us
compared to the very faint stars in surviving UFDs, identifying them in the
halo allows for more detailed studies of accreted systems.

  5. 5.

From a kinematic analysis, we found that all program stars had extremely large
and retrograde $V$ velocities. This retrograde motion is indicative of objects
that were accreted by the Milky Way rather than formed within it. This
supports our conclusion that the stars with the lowest Sr and Ba abundances
are likely some of the earliest accreted objects.

  6. 6.

We searched the literature for additional SASS stars using the [Sr/H] $<-$4.5
cut and found 61 more such stars in the halo. A significant fraction of stars
with [Fe/H] $<-$3 show these very low neutron-capture abundances. This fact
supports the idea that the most metal-poor stars are also the oldest stars in
our Galaxy (or at least the earliest stars that were accreted by the proto-
Galaxy).

  7. 7.

The population of SASS stars is easily identifiable with low Fe and Sr
abundances. This easy identification of ancient stars will ultimately allow
insights into the earliest star-forming environments, the origins and
histories of the surviving UFDs, as well as the earliest assembly phase of the
Milky Way.

  8. 8.

The earliest systems were likely enriched by a very small number of supernovae
whose yields were somewhat inhomogeneously mixed. This produced variations in
light element abundances, including carbon abundances, from system to system.
Neutron-capture elements were sparsely produced and/or diluted likely in
different ways, with $r$-process nucleosynthesis playing a role.

Indeed, UFDs offer insights into the conditions of the early universe. In
particular, they provide glimpses into early star formation environments,
early chemical enrichment events, and even the earliest phases of the Milky
Way’s formation. However, their limited number of observationally accessible
stars naturally limits their further exploration. Now, SASS stars may fill
that gap and provide a new important line of access to the early cosmos. These
stars, which now lurk in the Galactic halo, may shed further light on the
nature and origins of the oldest stars in the universe.

## Acknowledgements

HDA, ASF, and CGF gratefully acknowledge support from the MIT Undergraduate
Research Opportunities Program Office. AF acknowledges support from NSF grants
AST-1716251 and AST-2307436. MKM and AF acknowledge partial support from NSF
grant OISE 1927130 (International Research Network for Nuclear
Astrophysics/IReNA).

This work has made use of data from the European Space Agency (ESA) mission
Gaia (https://www.cosmos.esa.int/gaia), processed by the Gaia Data Processing
and Analysis Consortium (DPAC,
https://www.cosmos.esa.int/web/gaia/dpac/consortium). Funding for the DPAC has
been provided by national institutions, in particular the institutions
participating in the Gaia Multilateral Agreement. This work made use of NASA’s
Astrophysics Data System Bibliographic Services. This research has made use of
the SIMBAD database, operated at CDS, Strasbourg, France. This work made use
of Astropy:666http://www.astropy.org a community-developed core Python package
and an ecosystem of tools and resources for astronomy (Astropy Collaboration
et al., 2013, 2018, 2022).

This work made use of general-purpose Python libraries, namely numpy (Harris
et al., 2020), matplotlib (Hunter, 2007), pandas (Wes McKinney, 2010), and
scipy (Virtanen et al., 2020).

## Data Availability

The individual line measurements are provided as supplementary material. The
reduced spectra can be obtained by reasonable request to the corresponding
author.

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