11institutetext: Kapteyn Astronomical Institute, University of Groningen,
Landleven 12, 9747 AD Groningen, The Netherlands22institutetext: Université de
Strasbourg, CNRS, Observatoire astronomique de Strasbourg, UMR 7550, F-67000
Strasbourg, France33institutetext: Institute of Astronomy, University of
Cambridge, Madingley Road, Cambridge CB3 0HA, UK44institutetext: Max-Planck-
Institut für Astronomie, Königstuhl 17, D-69117 Heidelberg,
Germany55institutetext: Department of Astronomy, Stockholm University,
AlbaNova University Centre, SE-106 91 Stockholm, Sweden66institutetext: Dept.
of Physics and Astronomy, University of Victoria, P.O. Box 3055, STN CSC,
Victoria BC V8W 3P6, Canada77institutetext: Astronomisches Rechen-Institut,
Zentrum für Astronomie der Universität Heidelberg, Mönchhofstraße 12-14, 69120
Heidelberg, Germany88institutetext: Instituto de Astrofísica de Canarias,
Calle Vía Láctea s/n, 38206 La Laguna, Santa Cruz de Tenerife,
Spain99institutetext: Universidad de La Laguna, Avda. Astrofísico Francisco
Sánchez, 38205 La Laguna, Santa Cruz de Tenerife, Spain1010institutetext:
Isaac Newton Group of Telescopes, Apartado 321, 38700 Santa Cruz de la Palma,
Tenerife, Spain1111institutetext: Department of Physics and Astronomy,
University of Sheffield, Sheffield S3 7RH, UK1212institutetext: Department of
Physics & Astronomy, University of Tampa, 401 West Kennedy Boulevard, Tampa,
FL 33606, USA1313institutetext: Department of Physics and Astronomy, Dartmouth
College, Hanover, NH 03755, USA1414institutetext: Institute of Physics,
Laboratory of Astrophysics, Ecole Polytechnique Fédérale de Lausanne (EPFL),
Observatoire de Sauverny, 1290 Versoix, Switzerland1515institutetext: Núcleo
de Astronomía, Facultad de Ingeniería y Ciencias Universidad Diego Portales,
Ejército 441, Santiago, Chile1616institutetext: CAS Key Laboratory of Optical
Astronomy, National Astronomical Observatories, Chinese Academy of Sciences,
Beijing 100101, People’s Republic of China1717institutetext: Université Côte
d’Azur, Observatoire de la Côte d’Azur, CNRS, Laboratoire Lagrange, Nice,
France

# The Pristine survey - XXVI. The very metal-poor Galaxy††thanks: Based on
service-mode and visitor-mode observations made with the INT/IDS operated on
the island of La Palma by the Isaac Newton Group of Telescopes in the Spanish
Observatorio del Roque de los Muchachos of the Instituto de Astrofísica de
Canarias.

Chemodynamics through the follow-up of the Pristine-Gaia synthetic catalogue

Akshara Viswanathan Corresponding author 11<EMAIL_ADDRESS>Zhen Yuan
(袁珍) 22 Anke Ardern-Arentsen 33 Else Starkenburg 11 Nicolas F. Martin 2244
Kris Youakim 55 Rodrigo A. Ibata 22 Federico Sestito 66 Tadafumi Matsuno 77
Carlos Allende Prieto 8899 Freya Barwell 10101111 Manuel Bayer 11 Amandine
Doliva-Dolinsky 1212131322 Emma Fernández-Alvar 8899 Pablo M. Galán-de Anta
1010 Kiran Jhass 10101111 Nicolas Longeard 1414 José María Arroyo-Polonio 8899
Pol Massana 1010 Martin Montelius 11 Samuel Rusterucci 22 Judith Santos 1010
Guillaume F. Thomas 8899 Sara Vitali 10101515 Wenbo Wu 16168899 Paige Yarker
1010 Xianhao Ye 16168899 David S. Aguado 8899 Felipe Gran 1717 Julio Navarro
66

(Received ; accepted xx)

###### Abstract

Context. The Pristine-Gaia synthetic catalogue of reliable photometric
metallicities makes use of spectrophotometric information from Gaia DR3 XP
spectra to calculate metallicity-sensitive CaHK magnitude, which in turn
provides photometric metallicities for $\sim$30 million FGK stars using the
Pristine survey model and the survey’s training sample.

Aims. We perform the first low-to-medium-resolution spectroscopic follow-up of
bright (G¡15) and distant (upto 35 kpc) very and extremely metal-poor (V/EMP,
[Fe/H]¡-2.5) red giant branch stars from this catalogue , to evaluate the
quality of the photometric metallicities and study the chemodynamics of these
V/EMP stars.

Methods. We use Isaac Newton Telescope/Intermediate Dispersion Spectrograph
(INT/IDS) observations centred around the calcium triplet region ideal for
V/EMP stars for this spectroscopic follow-up.

Results. We find that 76% of our stars indeed have [Fe/H] ¡ -2.5 with these
inferred spectroscopic metallicities and only 3% are outliers with [Fe/H] ¿
-2.0. We report a success rate of 77% and 38% in finding stars with [Fe/H] ¡
-2.5 and -3.0 respectively. This is a huge improvement compared to the
literature in the selection of V/EMP stars based on photometric metallicities
and will allow for 10,000-20,000 homogeneously analysed EMP stars using the
WEAVE survey follow-up of Pristine EMP candidates. Using kinematics, we
associate 20%, 46%, and 34% of the stars to be confined to the disc plane, or
to have inner and outer halo orbits, respectively. Based on their integrals-
of-motion, we are able to associate these V/EMP stars to the metal-poor tail
of the metallicity distribution functions of known accretion events such as
Gaia-Enceladus-Sausage (GES), LMS-1/Wukong, Thamnos, Helmi streams,
Sagittarius, Sequoia, and other retrograde mergers. For the stars that orbit
close to the disc plane, we find that the prograde region with low vertical
action is overdense with a significance of 4$\sigma$ as compared to its
retrograde counterpart. We also find three new (brightest) members of the most
metal-poor stellar stream, C-19, one of which is 50∘ away from the main body
of the stream. This is the first member of C-19 found at positive height above
the disc plane. Our measured mean metallicity, velocity dispersion, and stream
width are consistent with the literature, but our results favour a slightly
farther distance ($\sim$21.5 kpc) for the stream.

Conclusions. With this work, we publish a catalogue (and 1D spectra) of 215
V/EMP stars from this first spectroscopic follow-up of the Pristine-Gaia
synthetic catalogue of photometric metallicities and showcase the power of
chemokinematic analysis of bright and distant red giant stars in the V/EMP
end.

###### Key Words.:

Galaxy: stellar content - Galaxy: halo - Galaxy: kinematics and dynamics -
stars: Population II - stars: Population III - techniques: spectroscopic

##  1 Introduction

The enduring presence of the most metal-poor stars within the Milky Way serves
as an invaluable window to the early Universe and the pristine conditions in
which these stars originated. These stellar relics are believed to have
coalesced from material enriched by the initial generations of stars, offering
crucial insights into the characteristics of their progenitors through their
chemical compositions (Beers & Christlieb, 2005; Frebel & Norris, 2015).
Furthermore, by studying the dynamical properties of these ancient stars, we
gain valuable insights into the early formation processes of our Galaxy.

The metal-poor halo of the Milky Way has become the heart of several accreted
structures, remnants of the Galaxy’s tumultuous merger history. These
structures include the Gaia-Enceladus-Sausage (GES), Sequoia, Thamnos, Helmi,
LMS-1/Wukong, Cetus-Palca, and Sagittarius streams (Belokurov et al., 2018;
Helmi et al., 2018; Barbá et al., 2019; Myeong et al., 2019; Koppelman et al.,
2019; Yuan et al., 2020; Naidu et al., 2020; Thomas & Battaglia, 2022; Ibata
et al., 1994), among others, each offering unique insights into past galactic
interactions. The recent discovery of numerous stellar streams further
underscores the ongoing accretion events from dwarf and ultra-faint galaxies,
as well as globular clusters, enriching the Galactic halo (Li et al., 2022a;
Ibata et al., 2021, 2023; Martin et al., 2022a, b; Viswanathan et al., 2023).
Additionally, recent studies suggest that a significant fraction of the halo
stars may have formed in-situ, representing a blend of components including an
$\alpha$-rich splashed hot thick disc and stars born in a primordial, hot, and
disordered state (e.g., Bonaca et al., 2017; Haywood et al., 2018; Di Matteo
et al., 2019; Belokurov et al., 2020; Belokurov & Kravtsov, 2022). This
comprehensive understanding of the Milky Way’s metal-poor population provides
a foundation for unraveling the complex tapestry of Galactic formation and
evolution.

The prevailing narrative across various cosmological simulations suggests that
Very Metal-Poor (VMP, [Fe/H] ¡ -2.0) stars residing within the central regions
of the Milky Way, including the bulge, represent some of the oldest stars in
our Galaxy. These stars serve as invaluable tracers of the early Galactic
assembly, offering critical insights into its formation history (Diemand et
al., 2005; Tumlinson, 2010; Starkenburg et al., 2017a; El-Badry et al., 2018).
On the observational front, numerous VMP stars have been scrutinised for their
chemistry and kinematics, with particular emphasis on regions such as the
bulge and the disc (Howes et al., 2014; Lucey et al., 2019; Arentsen et al.,
2020), providing essential data for understanding their origins. The chemical
characteristics of these stellar populations unveil a diverse assortment of
objects that contributed to the formation of the inner Galaxy. Some stars
exhibit chemical signatures indicative of formation in systems resembling
ultra-faint dwarf galaxies, while others display traits consistent with birth
within globular cluster-like environments (Schiavon et al., 2017).
Additionally, as pointed out earlier, there is evidence pointing towards a
significant presence of in-situ metal-poor stars within the inner Galaxy
(Belokurov & Kravtsov, 2022, 2023; Rix et al., 2022; Ardern-Arentsen et al.,
2024). This amalgamation of observational and theoretical findings paints a
picture of the early stages of Galactic evolution, underscoring the complex
interplay between various stellar populations and their environments.

Finding the most metal-deficient ([Fe/H] ¡ -3.0) stars has long been
recognised as a formidable challenge due to their exceptionally rare
occurrence rate (Bond, 1981). These stars are predominantly located in the
Milky Way’s halo component, and the ratio of halo-to-disc stars in the solar
neighborhood is approximately 10-3 (Bahcall & Soneira, 1980). Moreover, the
number of stars diminishes exponentially by a factor of roughly $\sim$10 or
more for each dex decrease in metallicity (Hartwick, 1976). In the solar
neighborhood, this translates to expectations of encountering one star with
[Fe/H] = -3.0 among every 65,000 stars, and one star with [Fe/H] = -3.5 among
200,000 stars (Youakim et al., 2020). To effectively address these challenges
and advance our understanding of the old halo component with sufficient
statistical power, efficient selection techniques are imperative.

Over the past four decades, numerous dedicated searches have been conducted to
assemble large samples of the most metal-poor stars. Various techniques have
been employed to identify these stars, including the pursuit of high-proper
motion stars exhibiting ultraviolet excesses (Ryan & Norris, 1991), the
detection of objects with diminished Ca II H & K lines in extensive objective-
prism surveys (Beers et al., 1985; Christlieb et al., 2008), and the
utilization of metallicity-sensitive (narrow-band) photometry (Schlaufman &
Casey, 2014; Starkenburg et al., 2017b; Wolf et al., 2018; Cenarro et al.,
2019; Almeida-Fernandes et al., 2022). Moreover, very and extremely metal-poor
(V/EMP) stars have been identified in greater abundance through large-scale
spectroscopic surveys such as the Sloan Digital Sky Survey (SDSS, York et al.,
2000), the Large sky Area Multi-Object fiber Spectroscopic Telescope (LAMOST,
Zhao et al., 2006; Li et al., 2018), the RAdial Velocity Experiment (RAVE,
Steinmetz et al., 2006), the GALactic Archaeology with HERMES spectroscopic
survey (GALAH, Buder et al., 2021), and the Gaia Radial Velocity Spectrometer
(RVS, Recio-Blanco et al., 2023; Matsuno et al., 2022; Viswanathan et al.,
2024). These photometric and spectroscopic surveys are often complemented by
dedicated follow-up efforts in the VMP end (e.g. Aguado et al., 2019; Yong et
al., 2021; Li et al., 2022b; Witten et al., 2022).

While astrometric data from the Gaia mission (Gaia Collaboration et al., 2022,
2023) has revealed significant stellar kinematic signatures attributed to
earlier accretion events, there persists a crucial requirement for precise
metallicities, particularly in the extremely metal-poor regime (EMP, [Fe/H] ¡
-3.0), for many stars. Such detailed metallicity measurements are essential
for constructing a comprehensive chemodynamic framework, enabling a thorough
understanding of the early structures that persist within the present Milky
Way halo.

In this work, we aim to increase the amount of chemodynamically analysed
bright (and distant) metal-poor red giant branch (RGB) stars. We use the
Pristine-Gaia synthetic catalogue of photometric metallicities to select V/EMP
stars. These metallicities are estimated using the Gaia XP low-resolution
spectrophotometry by determining the flux around the metallicity-sensitive
narrowband Ca II H$\&$K lines (CaHK hereafter), and running them with the
Pristine survey model. We present the first results of our low-to-medium-
resolution follow-up spectroscopy of 215 stars in the Pristine-Gaia synthetic
catalogue to assess the performance of the survey’s photometric pre-selection.
We subsequently use Gaia’s astrometry, distances, radial velocities (where
available) and the follow-up spectra’s radial velocities to calculate the
dynamics, and discuss their chemo-kinematic properties. Section 2 summarises
the target selection strategy for this spectroscopic follow-up program.
Section 3 presents the spectral analysis methods. Section 4 presents the
success rates of the photometric metallicities and discuss the chemo-
kinematics of the V/EMP stars. Section 5 discusses in detail the implications
of disc-like V/EMP stars in the sample and new tentative members of the most
metal-poor C-19 stellar stream. In section 6, we summarise the conlusions and
outlook.

##  2 Target Selection

In this section, we will describe how the V/EMP stars were selected using the
Pristine-Gaia synthetic photometric metallicity catalogue. For a comprehensive
overview of the Pristine survey, we refer the readers to Starkenburg et al.
(2017b). Additionally, for detailed information on the Pristine-Gaia synthetic
catalogue and the Pristine data release 1, we refer the readers to the work by
Martin, Starkenburg et al. (2023, hereafter MS23).

###  2.1 The Pristine survey and success rates

The Pristine survey focuses on observing the sky of the northern hemisphere
using the MegaCam wide-field imager installed on the Canada France Hawaii
Telescope situated on Mauna Kea. This survey employs a narrowband filter
centred on the metallicity-sensitive CaHK in the near UV $(3968.5$ and
$3933.7\AA)$. When combined with Sloan Digital Sky Survey (SDSS) g, r, i
filters, or the Gaia broadband BP-RP (blue and red photometric) and Gaia G
bands, this narrowband filter has been demonstrated to provide very reliable
estimates of stellar metallicity and an exceptional tool for identifying V/EMP
stars.

The Pristine photometry, in conjunction with SDSS broadband photometry, has
paved the way for medium and high-resolution spectroscopic follow-up studies.
These investigations specifically target stars with the lowest metallicity
estimates derived from Pristine CaHK observations. Through dedicated medium-
and high-resolution spectroscopic follow-up studies of Pristine-selected
candidates for EMP stars, we have found strong evidence that the majority of
these stars are indeed (very) metal-poor. Spectroscopic analyses of these
candidates have confirmed that approximately 20% of the EMP star candidates
possess [Fe/H] ¡ -3.0 and 70% of the VMP star candidates are below [Fe/H] ¡
-2.0 when carefully considering quality flags during target selection. This
conclusion is supported by the work of Youakim et al. (2017); Aguado et al.
(2019) through medium-resolution follow-up and Caffau et al. (2017);
Starkenburg et al. (2018); Bonifacio et al. (2019); Caffau et al. (2020); Venn
et al. (2020); Kielty et al. (2021); Lardo et al. (2021); Lucchesi et al.
(2022); Caffau et al. (2023) through high-resolution follow-up studies.
However, it is essential to account for potential variability in these stars
and undertake thorough photometric quality cuts to ensure accurate
characterization of their metallicities (Lombardo et al., 2023).

###  2.2 Pristine-Gaia synthetic catalogue of photometric metallicities

Figure 1: Colour absolute magnitude diagram (CaMD) of the INT observable
targets that pass the quality cuts described in section 2.2 in gray
overplotted with the RGB stars that were observed and presented in this paper
in purple and green. Purple star symbols correspond to good quality parallax
used to select the RGB stars and green triangle symbols correspond to bad
quality parallax used to select the RGB stars for which the distances were
derived using photometry and isochrone fitting. Blue dashed lines refer to age
13 Gyr and metallicity [M/H] -2.2 PARSEC isochrone shifted by $\pm 0.2$ mag in
$BP_{0}-RP_{0}$ and $\pm 0.75$ mag in $M_{G}$. Blue polygon is the absolute
magnitude limited ($M_{G}<2.0$) target selection area.

The creation of the Pristine-Gaia synthetic catalogue is described in detail
in MS23. We summarise the method here briefly. With the newest data release of
Gaia (Gaia Collaboration et al., 2022, DR3), the spectrophotometric BP/RP
information (De Angeli et al., 2023) was used to construct a comprehensive
catalogue of synthetic CaHK magnitudes, which emulate the narrow-band
photometry employed in the Pristine survey. Several other recent works
released a catalogue of metallicity estimates based on these Gaia XP spectra
(Lucey et al., 2023; Andrae et al., 2023; Zhang et al., 2023). At the same
time, the excellent Gaia accuracy allows us to reprocess and recalibrate the
entire Pristine CaHK dataset comprising of approximately 11,000 images
obtained since 2015, resulting in an updated survey that now covers more than
6,500 square degrees.

The improved photometric catalogue exhibits enhanced accuracy, achieving a
precision level of 13 millimagnitudes (mmag), a notable improvement from the
estimated 40 mmag during the first year of data. As part of the updated model,
the Pristine method for deriving the photometric metallicity of a star from
CaHK magnitudes and broadband magnitudes now relies solely on Gaia broadband
information (G, GBP, GRP), as opposed to SDSS previously. Additionally, a more
reliable iterative approach for addressing extinction correction has been
implemented. This involves incorporating extinction correction on the Gaia
broadband magnitudes and CaHK synthetic (or Pristine CaHK) narrowband
magnitudes by considering the star’s photometric temperature and metallicity.
To mitigate the impact of photometrically variable sources, which can
introduce spurious metallicities, a variability model based on the photometric
uncertainties of the 1.8 billion Gaia sources was also implemented in the
catalogue. Employing both the Pristine CaHK magnitudes and the BP/RP based
synthetic CaHK magnitudes within the Pristine model, two catalogues of
photometric metallicities for reliable stars were made public: the Pristine-
Gaia synthetic catalogue and the Pristine data release 1 (DR1) catalogue of
photometric metallicities, encompassing stars common to both Pristine and the
BP/RP catalogue of Gaia DR3. The latter served as the first data release of
the Pristine survey and provides deeper and better signal-to-noise (S/N) data
for stars in common.

Both these catalogues facilitate the construction of reliable samples of
metal-poor stars, with a particular emphasis on tracking V/EMP stars. The
Pristine-Gaia synthetic catalogue provides photometric metallicities across a
vast portion of the sky, while the Pristine BP/RP catalogue, limited to the
Pristine survey’s footprint, offers notably high-quality metallicities,
extending to significantly fainter stars.

In this work, we use the Pristine-Gaia synthetic catalogue of photometric
metallicities to select V/EMP targets that have reliable distances and are on
the red giant branch (hereafter RGB) of stellar evolution for a dedicated
spectroscopic follow-up. We choose to follow-up red giants as they are
intrinsically bright and probe large distances out into the Galactic halo,
allowing us to study the chemodynamics of the Milky Way out to large distances
and down to very low metallicities. For this target selection, we use the
following quality cuts on the parent sample, most of which are recommended by
MS23. Because some of the cuts recommended by MS23 were introduced after this
follow-up program commenced, these cuts are only implemented later in the
manuscript as stricter quality cuts in section 4.1. The quality cuts initially
used for target selection are defined as follows:

Figure 2: Mollweide projection of the Galactic coordinates for the V/EMP stars
with [Fe/H]¡-2.5 in the Pristine-Gaia synthetic catalogue, displayed in gray.
The large scale patterns visible in the map are mainly attributed to the
scanning law of the Gaia satellite. The purple stars represent the subset of
stars observed and analysed in this spectroscopic follow-up.

  * •

The follow-up is done from the northern hemisphere. Declination greater than
-30 degrees - which is the observability of the INT/IDS telescope facility
used for the follow-up (dec¿-30).

  * •

Photometric metallicity less than -2.5 dex (FeH_CaHKsyn¡-2.5). Note that this
cut was based on the first generation of photometric metallicities using the
Pristine survey model on Gaia XP based CaHK magnitudes. Over the course of
this work, the photometric metallicities were updated as a result of
improvement in the methods. Most of the selected stars still have a V/EMP
metallicity inference, also in the published catalogue (MS23). In the rest of
the paper, we only ever list the final MS23 metallicities.

  * •

Percentages of Monte Carlo iterations used to determine [Fe/H] uncertainties
inside the grid is greater than 80% (mcfrac_CaHKsyn¿0.8)

  * •

84th percentile value of the probability distribution function (PDF) of the
photometric metallicity is greater than -3.999 (FeH_CaHKsyn_84th¿-3.999)

  * •

Photometric metallicity uncertainty less than 0.5 dex (0.5*(FeH_CaHKsyn_84th
\- FeH_CaHKsyn_16th)¡0.5 dex)

  * •

Extinction on B-V magnitude is less than 0.5 (E(B-V) ¡0.5)

  * •

Brightness cut on measured Gaia G magnitude (phot_g_mean_mag<15.5) to be able
to observe with INT/IDS with short exposure times.

  * •

Photometric quality cut that is defined as $C^{*}$¡3$\sigma_{C^{*}}$. Cstar is
Gaia DR3 corrected flux excess, C∗, as defined in equation 6 of Riello et al.
2021 and Cstar_1sigma is normalized standard deviation of C∗ for the G
magnitude of this source, as defined in equation 18 of Riello et al. 2021
(abs(Cstar)¡3*Cstar_1sigma)

  * •

Parallax cuts:

    * –

Good parallax giants: parallax_over_error>5

    * –

Bad parallax giants: Poorly constrainted small parallax (very likely a red
giant) - parallax_over_error¡5 and abs(parallax)¡0.2.

  * •

No star within radius rmax from the centre of globular clusters, with rmax
being the rough estimate of the size of a cluster defined by Vasiliev &
Baumgardt (2021)

  * •

Shifted VMP old isochrone (PARSEC isochrone111stev.oapd.inaf.it/cgi-bin/cmd of
age 13 Gyr and metallicity -2.2 dex, Bressan et al., 2012; Pastorelli et al.,
2020) selection to remove horizontal branch stars and stars away from RGB
(used to reliably select RGB stars as targets) - this selection is also
limited to absolute magnitude of 2 mag and below. The magnitude-limited
shifted isochrone selection is shown as the filled blue region in Figure 1.

  * •

No overlap with literature high-resolution follow-up of VMP stars (Hansen et
al., 2018; Sakari et al., 2018; Li et al., 2018; Li et al., 2022b) and the
Pristine survey training sample that also includes Pristine survey’s own
spectroscopic follow-up programs as described in MS23. We end up having a
handful of stars (14) that overlap with the SAGA database of metal-poor stars
(Suda et al., 2008) (mostly from Li et al. (2022b) follow-up to avoid
uncertainties due to many different follow-up methods) which we use to define
the systematic uncertainty in the spectroscopic metallicity inferred in this
work.

All the stars that pass the above selection criteria (using good parallax
subsample) with inverted parallax as an approximation for distance
(=1/parallax) are shown in gray in Figure 1. The reliable parallax stars that
were spectroscopically followed-up are shown as purple star symbols and the
ones that were selected based on small parallax with large errors are shown as
green triangle symbols.

####  2.2.1 Distance estimates for poorly constrained small parallax selection

Numerous studies have provided compelling evidence that relying solely on
inverted parallax measurements to estimate distances can yield inaccurate
results especially if the parallax measurements are imprecise ($\varpi$¡0”
and/or $\varpi/\sigma_{\varpi}$¿0.2). Incorporating additional priors and/or
photometry has proven to enhance distance estimation (e.g., Bailer-Jones et
al., 2018, 2021; Queiroz et al., 2018; Anders et al., 2022). To address these
challenges, we adopt a Bayesian approach to infer distances for the stars with
bad parallax in our sample. We employ a methodology (see Sestito et al. (2019)
for a full description of the method), that enables the estimation of the PDF
or posterior on the distance estimates. In summary, the likelihood function is
formulated as the product of Gaussian distributions from the parallax and Gaia
photometry. Our prior incorporates a power-law stellar distribution for the
halo that we believe to be more suited for VMP stars than the prior used in
more general distance methods. Additionally, we utilize VMP ([M/H] = -2.5)
MESA/MIST isochrones222waps.cfa.harvard.edu/MIST (Dotter, 2016; Choi et al.,
2016) to account for the age and mass characteristics of VMP stars (11-13.8
Gyr, ¡1 M⊙) and their distribution according to an IMF-based luminosity
function in the Colour absolute magnitude diagram (CaMD).

This Bayesian framework, extensively employed in chemodynamical investigations
of VMP stars (e.g., Sestito et al., 2019, 2020, 2023; Venn et al., 2020),
yields distance estimates with low uncertainties even when faced with
significant parallax uncertainties. Typical distance uncertainties for these
stars in the sample is as low as $\sim$8%. This is achieved because the MIST
isochrones restrict the potential distance solutions for a star with a given
colour to two possibilities: a dwarf or a giant solution. The parallax
measurement typically favors one of the two solutions, or in the case of a
very poor parallax measurement, assigns different probabilities to the two
peaks. This approach eliminates the possibility of intermediate distance
solutions. With this method all our targets that were selected based on poorly
constrained small parallax ended up having a PDF that favored a giant
solution. These are shown as green star symbols along a rough isochrone-like
positioning on the CaMD in Figure 1. All the Gaia broadband photometry in
Figure 1 and other photometry used in target selection criteria are corrected
for extinction based on the iterative method described in detail in MS23 and
available in their public catalogue. The on-sky distribution of the stars
selected and observed with the telescope facility is shown as purple stars in
Figure 2 while the gray stars in the background are all stars that satisfy the
quality cuts described above (except the observability cut). The effect of
Gaia’s scanning pattern, which has resulted in certain regions being more
frequently observed, is clearly evident in the distribution of gray points in
Figure 2. This pattern becomes particularly prominent as we push the limits of
our observations to detect fainter stars, approaching Gaia XP spectra’s
limiting apparent magnitude. For a more detailed visualization of Gaia’s
scanning pattern in the analysis of sources with XP spectra, see De Angeli et
al. (2023).

##  3 INT/IDS spectroscopic follow-up of V/EMP candidates

Table 1: Technical information of the facility used in this analysis Telescope | Instrument | Detector | Grating | Filter | Slit | Central $\lambda$ | Range | Dispersion | Resolving power ($\Delta\lambda$)
---|---|---|---|---|---|---|---|---|---
2.5m INT | IDS | Red+2 | R1200R | GG495 | 1.37” | 8500 $\AA$ | 7829.5 - 9163.7 $\AA$ | 0.51 $\AA$ px-1 | 8019 (0.51 $\AA$)

The spectroscopic follow-up presented in this paper stemmed from an ongoing
long-term program at the Isaac Newton Telescope (INT) in La Palma (PI: A.
Viswanathan, semesters 22B, 23A, 23B, PI: E. Starkenburg, A. Arentsen,
semester 22A). We followed-up 215 stars over the course of 17 nights (along
with a few other member candidates from substructures in the Milky Way). The
observations include the following gray-to-bright nights: June 17 to June 21
2022 (clear weather, seeing between 0.5 and 0.8), July 19 to July 21 2022
(clear weather, seeing $\sim$0.8), 8 September to 11 September 2022 (partly
cloudy, seeing $\sim$0.9) and 30 December 2022 to 3 January 2023 (partly
cloudy, seeing $\sim$1.4). Our observing strategy was to go for very bright
targets when the seeing is relatively bad, stay 40∘ away from the moon and
away from the clouds. The results from semesters 23A and 23B will be published
in a forthcoming paper. We used the Intermediate Dispersion Spectrograph (IDS)
equipped with the RED+2 CCD. Our configuration included the R1200R grating, a
1.37” slit width, and the GG495 order-sorting filter. This instrumental setup
allowed us to achieve an effective spectral range of approximately 7850 to
9150 $\AA$ with a low-to-medium resolving power of $\sim$8000 (0.51 Å) over 2
pixels at the detector, specifically at 8500 $\AA$, covering the calcium
triplet region. This choice was made due to the follow-up of stellar streams
(the original proposal that was changed to follow-up of EMPs in the later
semesters with the use of a blue CCD).

The combination of the INT and the IDS instrument provided the necessary
capabilities to achieve accurate metallicity information and radial velocities
for the observed V/EMP candidates (see Table 1 for further technical details).

###  3.1 Data reduction

The spectra were processed and reduced using the Image Reduction and Analysis
Facility (IRAF, Tody, 1986) software package. Standard reduction techniques
were applied, including bias sky subtraction, removal of telluric lines, flat
fielding, spectrum extraction, sky subtraction, wavelength calibration - using
CuNe+CuAr lamps and the ONEDSPEC package in IRAF, and heliocentric radial
velocity correction. For wavelength calibration, we use a 10s arc per target
and use the 10s and 30s arc from the beginning of the night for the
calibration itself. The fringing effect, which can sometimes be a mild issue
for the RED+2 CCD on the
INT333https://www.ing.iac.es/astronomy/instruments/ids/ids_redplus2.html#4.,
was found to have an amplitude of 2% or less for the redmost wavelengths below
9000 $\AA$. Since the wavelength range concerned by this study is covered by
the Calcium II Triplet from approximately 8400 $\AA$ to 8750 $\AA$, no
fringing correction was necessary.

Figure 3: INT/IDS spectra of six V/EMP subsample stars selected using
photometric metallicities from the Pristine-Gaia synthetic catalogue, centred
on the calcium triplet lines. The stars are presented and sorted based on
their S/N ratio. The normalised spectra are shown with solid green lines while
the fits, derived from our pipeline for Voigt line profiles detailed in
section 3.3, are shown with dashed purple lines. Their S/N ratio, photometric
metallicities used to select them and spectroscopic metallicities derived are
also mentioned near the corresponding spectra. For clarity, the normalised
flux of each star is shifted by 0.5 and the spectra between the second and the
third calcium triplet line is cut-off. The calcium triplet lines are
highlighted in blue and central line is indicated by gray dashed lines. Note
that we focus on the three calcium triplet lines and remove the core between
the second and third calcium triplet lines only for clarity of the spectra
presented in this figure and for no scientific purpose.

Figure 3 illustrates spectra of six examples with varying S/N ratios between
and 10 and 140. The S/N values for each spectra are shown above each of the
spectra. The S/N values were computed in the calcium triplet region between
line 2 (8542 Å) and line 3 (8662 Å). Due to the brightness range of the
targets, (see left panel of Figure 4) we have achieved a mean S/N ratio of
34.2 with a minimum S/N of 9.7 with an average exposure time of 10 minutes
(the exposure time varies per target based on the brightness of the target and
seeing conditions).

###  3.2 Radial velocities

To derive the radial velocities from Calcium II Triplet lines, we first cut
the spectra within the wavelength range $\lambda$ = [8400, 8750] $\AA$, remove
the cores of the Calcium Triplet (only to calculate the continuum) using
dispcor, to find the continuum, and normalise the spectra using
onedspec.continuum routine. We then run fxcor to obtain radial velocities
through the cross-correlation of an observed standard star444Radial velocity
standards are chosen from this table: obswww.unige.ch/ udry/std/stdnew.dat
from the night with the observed V/EMP target. This is done by relative radial
velocity between the observed target and the standard star and then
subtracting the radial velocity of the standard star from the literature
within fxcor module. All the routines used here are available in PyRAF and
IRAF. At this stage of the analysis, the spectra has undergone several
preprocessing steps including normalisation, radial velocity shift, removal of
cosmic rays, subtraction of sky background, and correction for telluric
absorption. Resulting spectra are shown in green with the flux of each
spectrum shifted by 0.5 for better visualisation in Figure 3. Around 2 $\AA$
around each of the three calcium triplet regions are indicated within the blue
filled regions with the central wavelength as black dashed lines.

###  3.3 Metallicities

To compute the (very small) additional radial velocity offset that may have
been caused due to the difference in standard stars’ radial velocity and
equivalent widths, we use a modified version of the pipeline defined and used
by Longeard et al. (2022) and Viswanathan et al. (2024). It is noteworthy that
this radial velocity correction is negligible. The first step in our analysis
is to estimate the initial radial velocity of the star. We begin by creating a
smoothed spectrum for each star using a Gaussian kernel with a width
corresponding to 4 elements of resolution. This smoothing helps highlight the
calcium triplet (CaT) lines. We model each calcium triplet (CaT) line using a
Voigt profile and determine their positions by minimizing the square
difference between a simulated spectrum containing only the modeled CaT lines
and the observed spectrum. The Voigt profiles consider parameters such as the
amplitude of the Lorentzian component, the standard deviations of both the
Lorentzian (half-width at half-maximum) and Gaussian components, and the
Doppler shift in wavelength linked to the radial velocity offset based on the
definition in McLean et al. (1994).

To obtain an initial estimate of the radial velocity, we perform a cross-
correlation between the simulated spectrum and the observed spectrum. This
initial estimate is typically close to zero (mean of $\sim$1 km/s) due to
previous alignment to the rest frame using a standard star. The derived radial
velocity is then used as a starting point for further analysis and to improve
the accuracy of the radial velocity measurement from the fxcor pipeline in
IRAF. The mean statistical uncertainty in the inferred radial velocity is
$\sim$6.9 km/s. The radial velocity measured from the INT spectra show good
agreement with the radial velocity from Gaia’s Radial Velocity Spectrometer
(RVS) derived radial velocity for a large majority of our sample agreeing
within 25 km/s with Gaia RVS radial velocity, with an exception of $\sim$6% of
the stars with more than 2$\sigma$ difference between the two measurements.
For 2 stars, this difference is as large as 5$\sigma$. Such large difference
in radial velocity measurements suggests the possibility of these stars being
part of a binary system. We refrain from using these members (2$\sigma$ and
above) in the dynamical analysis of the paper. It is noteworthy that we do not
find any systematic offset between the two radial velocities. We use the Gaia
RVS radial velocity (where available) for the dynamical analysis of this paper
with 6 stars added using INT radial velocities for which Gaia radial
velocities are not available.

Next, we employ a Monte Carlo Markov Chain (MCMC) algorithm with a million
iterations per spectrum to fit the observed spectra. Using the simulated
spectrum that includes only the CaT lines shifted according to the initial
radial velocity guess and the initial amplitude guess (which is set to start
lower for EMP stars than normal metal-rich population in the Galaxy), we
derive the final radial velocity and equivalent widths (EWs) of the lines. The
step size defined to explore the parameter space to achieve the optimal
acceptance ratio is set based on the S/N of the spectrum. The fitting process
involves simultaneously optimising the central wavelengths, normalised fluxes,
and standard deviations of each line. Constraints are applied to ensure the
relative depths and widths of the lines agree with respect to each other. For
example, the second CaT line is constrained to be deeper than the third one
which in turn is set to be deeper than the first one. We also ensure that the
first line is not narrower than the other two lines, and the third line is
narrower than the second line. These constraints help maintain the expected
order and shape of the CaT lines during the analysis. The MCMC analysis is
performed for each star, and the parameters that maximize the likelihood are
selected as the best-fit values. The central wavelengths (and therefore the
radial velocities, v), the normalised fluxes of each line a1, a2, and a3 as
well as their standard deviations are fitted by minimizing the likelihood that
is defined as follows:

$\mathcal{L}_{i}=\frac{1}{\sigma_{i}\sqrt{2\pi}}\exp{\left(-0.5\frac{(s_{obs,i}-s_{sim,i})^{2}}{\sigma_{i}^{2}}\right)}$
(1)

where $\sigma_{i}$ is the flux uncertainty on star i, $s_{obs,i}$ is the
observed spectrum and $s_{sim,i}$ is the simulated spectrum of star i, all of
which are vectors. To calculate the EWs, we integrate each CaT line within a
15 $\AA$ window centred on the line. The integration is performed on the best-
fit simulated spectrum, and the resulting integrated values represent the EWs
for each line. The best-fit Voigt profile for each of the six example stars of
varying S/N is shown as purple dashed lines in Figure 3.

Figure 4: Distribution of Gaia G magnitudes and the derived metallicities from
Ca II Triplet region for the full follow-up spectroscopic sample of 220
Pristine-Gaia synthetic V/EMP stars. Median uncertainties in bins of 0.5 dex
in spectroscopic [Fe/H] are shown on the right.

To convert the equivalent widths (EWs) of the calcium triplet lines into
metallicity measurements, we employ the calibration provided by Carrera et al.
(2013) for Voigt profile fits. This calibration is known to reliably estimate
metallicities down to a value of -4.0. The inputs for this conversion include
magnitudes, calcium triplet equivalent widths, and distances (inverted
parallax) or height above or below the horizontal branch. Magnitudes adjusted
for extinction are taken from the input photometric metallicity catalogues.
Converting from Gaia G magnitude to Johnson-Cousins V or I magnitudes is
achieved through the conversion prescribed by Riello et al. (2021). To
determine the uncertainties associated with the metallicity measurements, we
utilize a Monte Carlo procedure. In each iteration, we randomly draw values of
the EWs from their PDF. We then calculate the spectroscopic metallicity for
each iteration, taking into account the individual uncertainties in photometry
and distances. Additionally, we incorporate the uncertainties associated with
the calibration relation itself by considering the proposed uncertainties on
the coefficients stated by Carrera et al. (2013). By performing this Monte
Carlo process, we construct a PDF that captures the uncertainty in the
metallicity determination for each star, considering all the relevant
parameters involved in the spectroscopic metallicity derivation. The standard
deviation (using a Gaussian approximation) of this distribution is taken as
the uncertainty on the metallicities. The calibration for metallicities
derived from the CaT equivalent width is grounded by Carrera et al. (2013) in
empirical observations of 55 metal-poor field stars. These stars were examined
at high-resolution, R¿20,000, measuring Fe I and Fe II spectral lines, which
are less influenced by non-local thermodynamic equilibrium (NLTE) effects
compared to the Ca II triplet lines. Consequently, the metallicities estimated
through this calibration closely align with NLTE equivalents.

The brightness range probed by the sample and the spectroscopic metallicities
obtained from this analysis are shown as 1D histograms in Figure 4. Four stars
that are either too hot (spectra with no distinguishable calcium triplet lines
indicating blue horizontal branch contamination) or have cosmic ray on the
calcium triplet lines are removed from the analysis. Futhermore, we also
remove stars that have an absolute magnitude in Johnson-Cousins I-band outside
the range, $-4<I<1$, within which the established EW to metallicity relation
works reliably. Stars that pass all these quality criteria in spectra and
magnitudes are referred to as ”quality spec” in Figure 5. The median
uncertainties on bins of 0.5 dex in metallicities is shown at the top of the
histogram in the right panel of Figure 4. We recommend to add to the reported
measurement uncertainty (only $\sim$0.08 in the median), a systematic
uncertainty of 0.15 dex derived from comparison with high-resolution analyses
of VMP stars from the SAGA database. This is based on the 14 stars in overlap
between our follow-up and the SAGA database. The remaining $\sim$200 stars are
newly discovered V/EMP stars from this work, ideal for high-resolution follow-
up. The on-sky distribution of the stars colour-coded by their inferred
metallicities is shown in Appendix A.

Figure 5: Comparison of photometric and spectroscopic metallicities for the
full followed-up sample of stars. The left panel shows the photometric
metallicities obtained using the Pristine-Gaia synthetic (XP spectra) versus
the spectroscopic metallicities derived from Calcium Triplet equivalent widths
for the entire sample. A small subset of the sample also has photometric
metallicities derived from higher S/N Pristine CaHK magnitudes, which are
shown in the right panel versus the spectroscopic follow-up metallicities.
Median uncertainties in bins of 0.5 dex in photometric [Fe/H] is shown in both
panels. The uncertainties increase with decreasing metallicities in general,
but the first and last bins are affected by low number statistics and do not
follow this trend.

##  4 Results

In this section, we present the findings of our spectroscopic analysis, which
encompasses both the dynamical properties and the metallicity measurements of
the observed stars.

Table 2: The number of stars with photometric metallicities from the Pristine-Gaia synthetic catalogue below –2.5 and –3.0, the number of stars confirmed spectroscopically below those metallicity thresholds, and the corresponding success rates are provided for three distinct datasets: the entire spectroscopic sample after basic quality cuts, the sample post-application of the stricter selection criteria to increase the 1:1 agreement in photometric and spectroscopic metallicities, and the subset of stars with photometric metallicities below –2.5 | Observed with recommended quality cuts | Stricter quality cuts | [Fe/H]phot ¡ –2.5
---|---|---|---
Total observed | 215 | 173 | 168
$\text{[Fe/H]}_{phot}$ ¡ –2.5 | 210 | 168 | 168
$\text{[Fe/H]}_{phot}$ ¡ –3.0 | 60 | 38 | 38
$\text{[Fe/H]}_{spec}$ ¡ –2.5 | 163 | 130 | 129
$\text{[Fe/H]}_{spec}$ ¡ –3.0 | 32 | 21 | 21
$\text{[Fe/H]}_{spec}$ ¿ –2.0 (outliers) | 9 (4%) | 6 (3%) | -
success [Fe/H] ¡ –2.0 | 206 (96%) | 166 (97%) | -
success [Fe/H] ¡ –2.5 | 162 (77%) | 129 (77%) | -
success [Fe/H] ¡ –3.0 | 23 (38%) | 14 (37%) | -

###  4.1 Performance of the Pristine-Gaia synthetic pre-selection

In the subsequent discussion of the results, we employ the terms ”Pristine
metallicities” and ”photometric metallicities” interchangeably to denote the
metallicity values obtained from the narrow-band synthetic CaHK derived from
Gaia XP spectra combined with Gaia broadband BP-RP data. Conversely, we use
the terms ”CaT metallicities” and ”spectroscopic metallicities” to refer to
the metallicities derived from the analysis of spectra acquired using the
INT/IDS facility in the calcium triplet region. Figure 5 shows the one-on-one
comparison between the spectroscopic metallicities of the V/EMP stars observed
and analysed in this work versus their photometric metallicities from the
Pristine-Gaia synthetic catalogue in the left panel and a subsample (69) of
stars that have photometric metallicities from the Pristine DR1 bright
catalogue (higher S/N CaHK narrowband values) in the right panel. The 1:1 line
and median uncertainties in photometric metallicities in bins of metallicities
are shown in both panels. The difference between photometric and spectroscopic
metallicities have a median offset of -0.08 dex and 1$\sigma$ deviation of
0.41 dex for the Pristine-Gaia synthetic catalogue and a median offset of
+0.11 dex and 1$\sigma$ deviation of 0.23 dex for the Pristine survey data
release 1 catalogue. The 1$\sigma$ deviation is reduced by about half with the
higher S/N CaHK narrowband measurements from the Pristine survey DR1 data. As
mentioned in section 2.2, some of the quality cuts defined by MS23 came into
effect after the start of our follow-up program. We investigate how these cuts
and some other stricter cuts can improve the one-to-one agreement of
photometric and spectroscopic metallicities. For the stricter quality cuts, we
implement the probability of photometric variability to be less than 30% (Pvar
¡ 0.3, which was a quality cut established while the follow-up program was
already ongoing), stricter photometric quality cut (abs(Cstar ¡
Cstar_1sigma)), astrometric quality cut (Renormalised Unit Weight Error RUWE ¡
1.4, only three stars observed have ruwe greater than 1.4 and have been
removed in the dynamical analysis section of this paper because they can be
non-single or otherwise problematic astrometric solutions), photometric
metallicity uncertainty less than 0.3 dex (0.5*(FeH_CaHKsyn_84th \-
FeH_CaHKsyn_16th) ¡ 0.3 dex), percentage of Monte Carlo iterations for [Fe/H]
uncertainties equal to a 100% (mcfrac_CaHKsyn == 1.0), 16th percentile value
of photometric metallicity greater than -3.999 (FeH_CaHKsyn_16th¿-3.999), and
CASU photometric data reduction flag (flag = -1, denoting very likely point-
sources - only for sources with Pristine DR1 measurements). Stars that pass
these extra quality cuts are referred to as ”quality spec+phot” in Figure 5.
The median offset and 1$\sigma$ deviation in the difference between
spectroscopic and photometric metallicities reduces from -0.09 dex to -0.07
and 0.41 dex to 0.35 dex respectively with the stricter quality cuts in the
input Pristine-Gaia synthetic catalogues, at the expense of losing 20% of the
VMP stars followed-up. The 1$\sigma$ deviation between photometric and
spectroscopic metallicities are lower ($\sim$0.23 dex) for Pristine DR1
photometric metallicities due to their high S/N CaHK measurements. It is
noteworthy that regardless of the stricter quality cut, the agreement is quite
good between the spectroscopic and photometric metallicities, and while the
stricter quality cuts remove few outliers, they also remove few interesting
VMP stars. For this reason, our conclusion is that the stricter quality cuts
do not improve the selection

Table 2 shows the number of stars present in various V/EMP photometric and
spectroscopic metallicity bins for three sub-samples from this spectroscopic
follow-up program: A total sample of 215 stars observed with this program
selected based on the recommended quality cuts and other cuts mentioned in
subsection 2.2, stricter quality cuts to improve the 1:1 agreement between
spectroscopic and photometric metallicities as described above, and all stars
with photometric metallicity estimates below –2.5 dex (a subsample 168 stars).
Firstly, the table outlines the number of stars within each sample that were
estimated by the Pristine-Gaia synthetic photometric metallicities to have
[Fe/H] values equal to or less than -2.5 or -3.0, respectively. It also
presents analogous counts based on spectroscopic [Fe/H] values and the
corresponding success rates. Success rates denote the fraction of stars
anticipated to have [Fe/H] values below a specified threshold by photometric
metallicities, which were indeed observed to possess spectroscopic [Fe/H]
values below that threshold. Outliers are stars with spectroscopic
[Fe/H]¿–2.0. A similar exercise was performed by Youakim et al. (2017) and
Aguado et al. (2019) using the medium spectroscopic follow-up of CFHT-based
Pristine survey photometric metallicities. Comparing the success rates, we
find that the success rates at finding stars with [Fe/H] ¡ -2.5 and [Fe/H] ¡
-3.0 has increased from 56% and 23% to 77% and 38% respectively and the
outliers ([Fe/H] ¿ -2.0) have gone down from 12% to 3%, with no catastrophic
outliers ([Fe/H] ¿ -1). We refrain from making estimates for the Pristine DR1
catalogue due to limitations in the number of stars but these numbers are
expected to be higher due to the higher S/N CaHK measurements and deeper
magnitude ranges probed by the Pristine survey. The improvement in success
rates are mainly due to the well-defined selection criteria weeding out
catastrophic outliers due to photometry, astrometry and limitations in the
photometric metallicity model. Given these new and improved success rates, we
predict that the photometric metallicities from the Pristine survey and the
Pristine-Gaia synthetic catalogues will bring more than 10,000-20,000
homogenously analysed EMP stars from the WEAVE low-resolution (R$\sim$5000)
follow-up of Pristine EMP stars (Jin et al., 2024) and allow for unprecedented
statistical analyses of the metal-poor Galaxy.

###  4.2 Chemokinematics of the V/EMP stars

To calculate the orbital properties of V/EMP stars from our follow-up program,
we use right ascension, and declination in J2016 format, proper motion in
right ascension and declination from Gaia DR3 data, with distances based on
parallax (from Gaia DR3) or photometry as described in section 2.2 and radial
velocities from Gaia RVS from DR3 when available and INT spectra based radial
velocities otherwise (for 6 stars). We adjust the stars for solar motion,
considering (U⊙, V⊙, W⊙) = (11.1, 12.24, 7.25) km/s according to Schönrich et
al. (2010), and for the motion of the local standard of rest (LSR) using vLSR
= 232.8 km/s as per McMillan (2017). Both the Galactocentric cartesian and
cylindrical positions and velocities of the stars are computed assuming R⊙ =
8.2 kpc (McMillan, 2017) and z⊙ = 0.014 kpc (Binney et al., 1997). Our
coordinate system is oriented such that x points towards the Galactic centre,
y aligns with the direction of motion of the disc, and positive (negative) z
signifies the height above (below) the disc. Angular momenta, Lz and L⊥ =
$\sqrt{\text{L}_{x}^{2}+\text{L}_{y}^{2}}$, are calculated for the stars with
the sign of Lz inverted to ensure it being positive for prograde orbits.
Energy, E, is determined using AGAMA (Vasiliev, 2019) and the Milky Way
potential from McMillan (2017). This potential, which is axisymmetric,
encompasses a stellar thin and thick disc, an HI gas disc, a molecular gas
disc, a bulge, and an NFW halo, which defaults to a spherical shape. The
integration time is set to 1 Gyr. The output orbital parameters are the
Galactocentric Cartesian coordinates (x, y, z), the maximum distance from the
Milky Way plane Zmax, the apocentric and pericentric distances, the
eccentricity, the energy E, and the spherical actions coordinates (Jϕ, Jr ,
Jz).

In the left top panel of Figure 6, we see the total energy (E) versus the
z-component of angular momentum (Lz). We see four stars (all below -2.5 in
[Fe/H]) that are unbound in our chosen gravitational potential of the Milky
Way, along with 7 stars that are not associated with any known substructure
region and have high apocentres. These unbound stars have a velocity of about
700 km/s at a distance of about 15 kpc which is too small to classify them as
hyper velocity stars. From the literature, we see that most of the hyper
velocity stars are coming from the Galactic disc. These 4 high energy stars
also have confined orbits (Zmax ¡ 3.5 kpc). However, given that their proper
motions and velocities are small, it is safe to say that these stars are most
likely unbound due to the measurement uncertainties. In the future, after a
careful target selection of high proper motion/velocity stars, a large sample
of stars in high energy orbits can be used to infer the escape velocity of the
Milky Way. Note that the high energy stars are not shown on the left panels
due to their very high apocentre and Zmax.

It is anticipated that stars originating from the same physical structure will
exhibit clustering in integrals of motion, energy and angular momentum
(especially Lz) within an axisymmetric potential, despite being widely
dispersed in configuration space (Helmi et al., 2000). Since the second data
release DR2 of the Gaia mission, phase mixed streams have been identified in
kinematically selected halo stars. These structures, including Gaia-Enceladus-
Sausage, Sequoia, Thamnos, Helmi streams, LMS-1/Wukong, Heracles/Kraken/Koala,
Cetus-Palca as discussed in Koppelman et al. (2019); Myeong et al. (2019);
Yuan et al. (2020); Naidu et al. (2020); Lövdal et al. (2022); Ruiz-Lara et
al. (2022); Malhan et al. (2022); Thomas & Battaglia (2022); Horta et al.
(2023); Dodd et al. (2023), are believed to have an accreted origin primarily
based on their inferred orbits, often supported further by chemical abundances
where available. Nevertheless, their chemical composition, essential for
determining the types of galaxies from which they originated, remains
inadequately characterized, particularly at the lowest metallicities.
Therefore, we associate our V/EMP stars to known substructures with ellipses
as detected using red giant stars from the H3 survey (Naidu et al., 2020) in
the integral-of-motion (IOM) space (defined using the substructures in
McMillan (2017) potential in Figure 23 from the same work). This is shown in
the top centre panel of Figure 6). We associate stars to a particular
accretion event based on the selection criteria presented in Naidu et al.
(2020). Most significantly we can associate stars with Sequoia+Arjuna+I’itoi
group (23 stars with a huge spread in distance between 2-32 kpc, see also the
small clump of the 3 VMP and EMP stars clustered tightly in this region and in
left middle panel), GES (total of 81 stars with $\sim$97% within 20 kpc, note
the high eccentricity stars clumped around L${}_{z}\sim$ 0 in the left bottom
panel) and the metal-poor LMS-1/Wukong (14 stars with medium eccentricity, the
slightly prograde clump of stars in left bottom panel). We associate a handful
of EMP stars with the Helmi streams (6), Thamnos (11), and Sagittarius (4),
with both Thamnos and the Helmi streams populating distances less than 10 kpc
while Sagittarius stars all lie above 7 kpc. There are also few stars in the
thick disc and the in-situ halo regions. We also find small clumps of stars
with very similar metallicities that could originate from low-mass merger
events or globular cluster disruptions. These groups are indicated with a
group number in Table 3, which will be released electronically with this work.
Due to the low number statistics and complex selection function, we refrain
from defining any new substructures from this work. In the future, a more
complete and large sample of V/EMP stars with dynamics and chemistry will
allow us to also identify the smallest and earliest galaxies that merged into
the Milky Way (Yuan et al., 2020).

In Figure 6, the three panels on the right illustrate the pericentric
distance, eccentricity, and maximum height from the plane relative to the
apocentric distance. Meanwhile, the right-hand three panels depict the energy
versus the rotational component of the action (top), rotational versus
vertical action (middle) and rotational action versus eccentricity (bottom).
In the right middle panel we see stars separated into prograde and retrograde
orbits. We clearly see that our sample has more prograde than retrograde
stars. The sample is classified into four distinct groups based on their Zmax
and apocentre, akin to the categorization outlined in Sestito et al. (2023).
These categories are described as follows:

  * •

Bulge group: Stars within this group remain confined within a sphere with a
radius of 3.5 kpc from the Galactic centre, i.e., apocentre ¡ 3.5 kpc. No
spectroscopic follow-up has been conducted in this region.

  * •

Confined group (purple stars): Stars within this group exhibit Zmax ¡ 3.5 kpc
and apocentre ¿ 3.5 kpc. They are primarily restricted near the Milky Way
disc, and this group comprises 41 (20%) stars that have undergone
spectroscopic follow-up and presented in this work

  * •

Inner halo group (green stars): Comprising 101 (46%) stars, this group
consists of stars with Zmax ¿ 3.5 kpc and apocentre ¡ 20 kpc.

  * •

Outer halo group (blue stars): This group comprises 73 (34%) stars
characterized by Zmax ¿ 3.5 kpc and apocentre ¿ 20 kpc.

As depicted in Figure 6, all stars, regardless of their group, show a range of
eccentricities. It is noteworthy that stars in the confined group do not
necessarily exhibit planar orbits. If we define planar orbits as stars with a
ratio of Zmax/apocentre ¡ 0.2, these constitute half of the stars in our
confined group. Of these stars, about 80% are in prograde disc-like orbits.
These findings are discussed in detail in the next section.

In the right middle panel of Figure 6, we see that about 71% of the stars that
are confined have small pericentre comparable to the length of the Galactic
bar (R${}_{b}\sim$ 3.5 kpc, Wegg et al., 2015; Lucey et al., 2023). Dillamore
et al. (2023) suggest that a possible co-rotation resonance of a centrally
concentrated halo population with the rotating bar could explain why we see
planar prograde stars in the V/EMP end in the solar neighbourhood. From our
sample of confined stars (of which about half of them are prograde planar),
indeed the bar can play a role in the orbit of this population. However, we do
see a population of confined stars (small purple symbols in right middle and
bottom panels of Figure 6) that have higher pericentres and low eccentricity,
more like disc-orbits. The on-sky distribution of the stars colour-coded by
their categories (confined, inner halo, outer halo) are shown in Appendix A.
Other dynamical spaces such as position distribution, velocity distribution,
and IOM space for these categories of stars are shown in Appendix B.

Figure 6: Chemokinematics of the V/EMP stars from this work. Top centre:
Energy E versus angular momentum in z direction Lz. The substructure regions
from Naidu et al. (2020) identified using RGB stars in the H3 survey are shown
as coloured ellipses to associate V/EMP stars from this work. Left: (Top) Same
as top centre but without the substructure ellipses. Stars with energy greater
than 0 are high energy/unbound stars in the Milky Way potential used. (Middle)
Rotational component of the action Jϕ versus vertical action Jz. (Bottom)
Rotational action Jϕ versus eccentricity. Right: (Top) The INT sample is
divided into 3 dynamical groups according to their apocentre and Zmax. The
grey shaded area denotes the forbidden region in which Zmax ¿ apocentre.
Vertical and horizontal lines separate the groups. (Middle) The groups are
shown in apocentre versus pericentre view. The grey shaded area denotes the
forbidden region in which pericentre ¿ apocentre. (Bottom) The groups in
apocentre versus eccentricity view. All EMP stars observable by the
spectroscopic follow-up program are shown in gray.

##  5 Discussion

In this section, we discuss in detail the implications of planar metal-poor
stars in the sample and present and discuss new members of the most metal-poor
stellar stream C-19.

###  5.1 Planar metal-poor stars

The oldest and most metal-poor stars are usually found in the bulge and halo
because they are relics from the era of the smallest, earliest galaxies that
merged into the Milky Way halo. However, there have been recent discoveries of
very, extremely and ultra metal-poor stars in prograde, planar, and disc-like
orbits confined to the plane (z ¡ 3 kpc, Sestito et al., 2019). The mechanisms
for producing such stars beyond just the halo has been highly debated. Some of
the proposed scenarios are as follows: (i) minor mergers that sink onto the
plane before disrupting due to dynamical friction (Abadi et al., 2003; Sestito
et al., 2019), (ii) stars from the building blocks of the proto-Milky Way or
filaments creating a halo that is slowly rotating prograde (Sestito et al.,
2019, 2020; Belokurov & Kravtsov, 2022, 2023; Rix et al., 2022; Zhang et al.,
2023), (iii) the in-situ formation of this component of the disc at early
times (Sestito et al., 2019, 2020; Di Matteo et al., 2020; Fernández-Alvar et
al., 2021, 2024; Nepal et al., 2024). Regarding the feasibility of the latter
scenario, in the NIHAO-UHD and FIRE simulations these stars are formed within
2-3 Gyr from the Big Bang, while the disc appears later (Sestito et al., 2021;
Santistevan et al., 2021). However, disc formation in simulations cannot be
evidence enough to rule out this scenario. (iv) A possible co-rotation
resonance of a halo population with the bar (Dillamore et al., 2023). Yuan et
al. (2023) illustrated that as a result of a rapidly decelerating bar, certain
bulge stars acquire rotational motion by becoming ensnared in co-rotating
regions and subsequently migrating outward along prograde planar orbits.
However, the proportion of stars influenced by this mechanism is insufficient
to explain all of the metal-poor rotators identified. Consequently, their
origin remains ambiguous.

Sestito et al. (2019) conducted an analysis of the kinematics and dynamics of
all ultra metal-poor (UMP, [Fe/H] ¡ -4.0) stars from the literature available
at the time. Their findings revealed that 11 out of 42 ($\sim$26%) known UMP
stars remain confined within 3 kpc of the Milky Way plane throughout their
orbital lifetime. Moreover, 10 out of these 11 UMP stars exhibit prograde
orbits (vϕ ¿ 0). Building on this work, Sestito et al. (2020) arrived at a
similar conclusion using a significantly larger dataset comprising 1027 V/EMP
stars. This expanded dataset incorporated spectroscopic follow-up data from
the Pristine survey (Aguado et al., 2019) and LAMOST spectroscopy (Li et al.,
2018). Their analysis revealed that approximately 31% of stars with —z— ¡ 3
kpc observed today never venture beyond 3 kpc of the disc plane. Additionally,
they examined their sample in the Jϕ-Jz projection of the action space.
Sestito et al. (2020) observed that the number of stars in prograde disc-like
orbits (characterized by high Jϕ and low Jz) is significantly greater than the
number of stars in retrograde disc-like orbits, with a statistical
significance of 5$\sigma$.

We conducted a similar orbital analysis to Sestito et al. (2019, 2020) and
investigated the behavior of our V/EMP stars in the action space, as depicted
in Figure 7. This figure can be directly compared to the results from Sestito
et al. (2020). We see an overdensity of prograde V/EMP stars confined to the
plane up to and even higher than the solar azimuthal action with low to
intermediate eccentricities. Our findings indicate that the fractions of VMP
and [Fe/H] ¡ -2.5 (chosen instead of EMP due to low number statistics in EMP)
stars with present-day —z— ¡ 3 kpc and Zmax ¡ 3 kpc are 31% and 21%,
respectively, which align very closely with the results of Sestito et al.
(2020). In the rotation versus vertical action plane, we identified stars
within the black dashed boxes representing disc-like prograde and retrograde
orbits. The region corresponding to low-Zmax corotating stars is denser by
4$\sigma$ (entire sample), 3.5$\sigma$ (VMP) and 3.8$\sigma$ ([Fe/H] ¡ -2.5)
compared to its retrograde counterpart (using Poisson statistics), slightly
lower than the findings reported by Sestito et al. (2020), but still
significantly more prograde than retrograde. It is essential to note, however,
that the methodologies, sample selections, and selection functions employed in
these studies vary significantly. The action space distribution of these
confined, confined planar (from section 4.2) and prograde planar (disc-like)
stars are shown in Appendix C.

Zhang et al. (2023) investigated this with XGBOOST derived metallicities from
XP spectrophotometry and found a similar (if not higher significance) prograde
to retrograde ratio of stars in the metallicity range -1.6 ¡ [M/H] ¡ -3. They
suggest that this asymmetry can be explained by a combination of a stationary
and a prograde halo (proto-MW) component in the VMP end. In a recent work by
Ardern-Arentsen et al. (2024), that investigates orbits of VMP stars in the
inner Galaxy (R ¡ 3.5 kpc), they report that VMP stars are likely consisting
of a mix of a prograde and a stationary halo component. The apocentre
distribution shows that VMP stars are very centrally concentrated, with a
steep drop towards higher apocentre (and not many with apocentre ¿ 10 kpc).
The apocentre and Zmax distribution of our planar stars are shown in Appendix
C.

The chemical properties of these planar stars indicate that a variety of
different systems contribute to the formation of this population. Dovgal et
al. (2024) shows that the large scatter in $\alpha$-elements of the planar
star might be indicative that the prograde planar stars is composed of
multiple systems. Sestito et al. (2024) shows that a planar prograde star with
high eccentricity has the same chemical signatures of ultra-faint dwarf stars.
Spectroscopic follow-up of some planar stars (retrograde+prograde) show that
these observed stars formed in the same formation site, which look like a
classical dwarf galaxy (Sestito, Fernandez-Alvar et al., in prep). Further
investigation, incorporating age distributions and/or more extensive chemical
abundance measurements along with extensive orbital analysis in the V/EMP end,
is essential to definitively determine the exact metallicity threshold at
which the early Milky Way disc originated (Viswanathan et al., in prep.).
Stars from this spectroscopic follow-up and the ongoing follow-up will add
more information in light of this problem and can be potential targets for
high-resolution follow-up in the future.

Figure 7: Distribution of low-metallicity stars from the INT sample in the
Jϕ–Jz action space. The action quantities are scaled by the solar values
(i.e., Jϕ⊙ = 2009.92 kpc km/s, Jz⊙ = 0.35 kpc km/s). Two black boxes represent
the prograde and the retrograde disc-like orbits showing an overdensity in the
prograde direction. All EMP stars observable by the spectroscopic follow-up
program at INT are shown in gray.

###  5.2 C-19 stellar stream new members

Figure 8: C19 stellar stream members from this work. Declination versus radial
velocity distribution (top) and declination versus right ascension
distribution (bottom) of the INT sample of V/EMP stars colour-coded by their
inferred spectroscopic metallicities with C-19 orbit from the STREAMFINDER
algorithm and Yuan et al. (2022) identified C-19 members overlaid. The
tentative C-19 members identified in this work is shown as large star symbols.
All EMP stars observable by the spectroscopic follow-up program at INT are
shown in gray. Figure 9: Right ascension versus distance distribution of V/EMP
stars from this work colour-coded by their spectroscopic metallicities. C-19
members identified are shown with black-dashed boxes around them. All EMP
stars observable by the spectroscopic follow-up program at INT are shown in
gray.

The discovery of the most metal-poor stream C-19 by Ibata et al. (2021);
Martin et al. (2022a, b) has sparked considerable interest, with its
interpretation as a disrupted Globular Cluster (GC) based on several key
observations. C-19 stream is the first and foremost exceptionally low-
metallicity structure in the universe. Notably, its remarkably narrow
metallicity dispersion and a range in Na abundances are consistent with
observations of globular clusters (Martin et al., 2022b; Yuan et al., 2022),
and contrasting with observations in dwarf galaxies or among the Galactic halo
field stars. What sets C-19 apart is its mean metallicity ([Fe/H] = -3.4,
EMP), which surpasses that of even the most metal-poor GC (Vasiliev &
Baumgardt, 2021) and represents the lowest metallicity structure ever found.
The absence of Galactic Globular Clusters observed below a metallicity of -2.5
has led to interpretations suggesting a metallicity floor beyond which GCs
cannot form. The Phoenix stream (mean [Fe/H] = -2.7, Balbinot et al., 2016;
Wan et al., 2020) and ED-2 stream (mean [Fe/H] = -2.6, Balbinot et al., 2023,
2024) are other streams that challenge the GC metallicity floor. However,
despite chemical evidence suggesting a Globular Cluster origin, the dynamical
width and velocity dispersion of C-19 resemble those of a dwarf galaxy. One
hypothesis suggests that the stream may have originated from a dark matter-
dominated dwarf galaxy, with the star cluster comprising most of the mass in
the system, although the remnants of such a dwarf galaxy have yet to be
discovered (Errani et al., 2022). Alternatively, Helmi & White (1999)
demonstrated through numerical simulations that the velocity dispersion of
streams can increase significantly at the turning point of their orbit,
potentially explaining C-19’s high velocity dispersion. However, determining
whether the stream is indeed at its turning point poses a challenge due to its
sensitivity to the assumed potential. Thereby, finding more members on C-19’s
orbit and brightest members (better suitable for high-resolution follow-up)
that belong to the streams will help us constrain the origin of this ancient
structure.

Figure 8 shows the right ascension versus declination in the bottom panel and
radial velocity versus declination in the top panel for all the stars analysed
in this work, colour-coded by the inferred spectroscopic metallicity. Overlaid
are the C-19 members from Yuan et al. (2022) in green, with one radial
velocity member discovered 30∘ away from the main body of the stream. From the
metallicities and the STREAMFINDER inferred C-19 orbit, we find five (three
new, two rediscovered) C-19 members (inferred metallicities between -3.24 and
-3.6, all consistent with C-19 mean metallicity within 2$\sigma$ considering
the measurement and systematic errors), two of which are the brightest members
of the stream discovered yet at the tip of the red giant branch (RGB). One of
the members discovered is 50∘ away from the main stream counterpart, above the
disc plane at positive height above the plane. This is the first C-19 member
discovered above the disc plane. This member increases the span of the stream
on the sky. The overdensity of stars in one part of the stream could be due to
the coverage in the region of the Pristine survey as it probes deeper to find
C-19 members with photometric metallicities. The mean metallicity of the
stream with the new members is -3.33$\pm$0.07, the stream width is
0.39${}^{\circ}\pm$0.07∘ and the velocity dispersion (excluding the member
above the disc plane) is 4.11$\pm$2.05 km/s consistent with the literature
values. Note that C-19 member stars are not associated with any of the
accretion events in Figure 6, according to the literature selections.
Although, Malhan et al. (2022) suggests that C-19 is mostly likely a stream
associated with the LMS-1/Wukong structure. C-19 stream and LMS-1/Wukong have
different orbits. Even though they are both polar, the orbital planes are
separated by a large angle. A more comprehensive look at the members of C-19
all over the sky with chemistry and dynamics will be presented in the upcoming
work (Yuan et al., in prep.) which will provide more insight on the progenitor
and formation of the C-19 stream.

Figure 9 shows the distance versus right ascension plane of all the stars
observed and analysed in this work, colour-coded by metallicity. Tentative
C-19 members from this work are highlighted with a dashed black box and the
members from the main body of the stream are clustered at around 21.5 kpc,
while the member above the disc plane lies at 17 kpc which indicates a
distance gradient across the stream. The mean distance of the stream is
slightly higher than what has been reported in the literature. Ibata et al.
(2021) estimated the distance to C-19 using the STREAMFINDER algorithm, which
determines a most likely heliocentric distance, based on the local
distribution of stars on the sky, in proper motion space, and in CMD space
which is not individually precise, but for a stream as a whole, it is a good
first guess. They found a distance between 16 and 22 kpc with no distance
gradient along the stream. Martin et al. (2022b) estimates the distance to
C-19 using 7 blue horizontal branch stars (17.5 kpc) at a favoured distance of
18 kpc to minimize the difference between the FeI and FeII iron abundances,
and a distance estimate based on Gaia parallaxes and the best fit C-19 orbit
(20.9 kpc). All these stars have bad parallax, which means that the PDF of the
distance inference depends highly on the photometry, extinction, systematic
offset in the method and the choice of isochrone which is most likely a bad
match for [Fe/H]¡-3 stars. Therefore, we do not expect these stars to have
distances as good as standard candles. However, Bonifacio et al., in prep.
followed-up and analysed fainter sub-giant branch stars from the C-19 stream
and found that an EMP isochrone fits better with 20.9 kpc distance in the
subgiant branch that would match well with our estimate. Based on our
analysis, we find a mean distance of 21.5 kpc, which is a slightly larger
distance than the literature-reported distances.

In Figure 9, the INT observable stars are in gray. Because these are good
parallax stars, we can see that the gray points end at about 10 kpc while due
to the bad parallax selected giants, our sample goes up to 35 kpc in the EMP
bright end, much farther than what is possible using only good parallax
selections. Given the success demonstrated in this follow-up paper, we pave
the way for the use of bright and distant metal-poor stars and streams from
these photometric studies as invaluable guides to the early universe with a
more homogeneous selection of giant stars (Viswanathan, Byström et al., in
prep.)

##  6 Conclusions and Outlook

With the recent Gaia DR3, we have spectro-photometric information from the XP
spectra for $\sim$219 million stars. Using these data, MS23 calculated
synthetic, narrowband, metallicity-sensitive CaHK magnitudes that mimic the
observations of the Pristine survey, and infer photometric metallicities using
the narrowband CaHK and broadband Gaia magnitudes trained on the survey’s
training sample. In this work, we present the first dedicated spectroscopic
follow-up of V/EMP stars from this all-sky catalogue of photometric
metallicities. The target selection is based on the quality cuts recommended
by the catalogue paper, MS23, in the bright end (Figure 1 and 2). We use the
INT/IDS instrument in La Palma for the low-to-medium-resolution spectroscopic
follow-up in the calcium triplet region, which is widely used to study VMP
stars (Table 1). We infer radial velocities and spectroscopic metallicities
using the pipeline based on Viswanathan et al. (2024, Figure 3). Due to the
available 6D phase space information, we also take a deeper look into the
dynamics of these stars in the context of Galactic archaeology. Our main
results are presented as follows:

Table 3: Description of the columns of the INT follow-up spectra V/EMP stars catalogue made available publicly in this work. Note that the reduced 1D spectra from this work is also made available publicly. All these data will be released electronically after acceptance and before upon reasonable request. Column | Description | Unit | Type
---|---|---|---
source_id |  Gaia DR3 Source ID | unitless | longint
ra |  Gaia DR3 right ascension in ICRS (J2016) format | degrees | float
dec |  Gaia DR3 declination in ICRS (J2016) format | degrees | float
G_0 | de-reddened Gaia G magnitude | unitless | float
BP_0 | de-reddened Gaia GBP magnitude | unitless | float
RP_0 | de-reddened Gaia GRP magnitude | unitless | float
distance | Photometric distance or inverted parallax based distance derived in this work | mas | float
distance_error | Uncertainty on the photometric distance | mas | float
snr | Signal-to-noise ratio of the normalised INT CaT spectra | unitless | float
v_r | Radial velocity measured from the INT CaT spectra | km/s | float
dv_r | Uncertainty in the measured radial velocity | km/s | float
ew1 | Equivalent width of the first calcium triplet line around 8498 Å | Å | float
dew1 | Uncertainty on the equivalent width of the first calcium triplet line around 8498 Å | Å | float
ew2 | Equivalent width of the second calcium triplet line around 8542 Å | Å | float
dew2 | Uncertainty on the equivalent width of the second calcium triplet line around 8542 Å | Å | float
ew3 | Equivalent width of the third calcium triplet line around 8542 Å | Å | float
dew3 | Uncertainty on the equivalent width of the third calcium triplet line around 8662 Å | Å | float
feh | Spectroscopic metallicity derived in this work using INT CaT spectra | unitless | float
dfeh | Measurement uncertainty associated with the spectroscopic metallicity derived | unitless | float
E | Total energy | kpc km2/s2 | float
Jphi | rotational action vector | kpc km/s | float
Jz | vertical action vector | kpc km/s | float
Jr | radial action vector | kpc km/s | float
zmax | maximum height above the plane of the orbit | kpc | float
rapo | apocentre of the orbit | kpc | float
rperi | pericentre of the orbit | kpc | float
ecc | eccentricity of the orbit | unitless | float
group_number | Group number for stars that are clustered in IOM space | unitless | int

  * •

The spectroscopic metallicities are inferred at a precision of 0.08 in
measurement added in quadrature with 0.15 dex in systematic uncertainty down
to $\sim$-3.6 in metallicity (Figure 4).

  * •

The photometric metallicities from the Pristine-Gaia synthetic catalogues
agree well with the inferred spectroscopic metallicities from our follow-up
with $\sim$77% and $\sim$38% success rates of having [Fe/H] ¡ -2.5 and -3.0
stars respectively. The outliers ([Fe/H ¿ -2.0) are reduced down to 3% with no
catastrophic outliers ([Fe/H ¿ -1.0). This is a huge improvement over the
existing methods that search for EMP stars (Figure 5 and Table 2). With the
enhanced success rates observed, we anticipate that the photometric
metallicities obtained from the Pristine survey and the Pristine-Gaia
synthetic catalogues will facilitate the identification of more than
10,000-20,000 uniformly analyzed EMP stars through the WEAVE low-resolution
(R$\sim$5000) follow-up of Pristine EMP stars (Jin et al., 2024). This will
enable unprecedented statistical analyses of the metal-poor Galaxy.

  * •

From the kinematics of V/EMP stars from this work, we find 20% of them in
confined orbits, 46% in the inner halo and 34% in the outer halo orbits (see
left panel of Figure 6).

  * •

We associate V/EMP stars with several known accretion events such as GES,
LMS-1/Wukong, Thamnos, Sequoia, Sagittarius and Helmi streams which would
define the EMP end of their metallicity distribution. Some of the small and
tight clumps in the E-Lz space could belong to low-mass mergers in the distant
past (Figure 6).

  * •

We find that 31% of our VMP stars with —z— ¡ 3 kpc do not venture beyond 3 kpc
in their orbital history. Along this line, the prograde region with low
vertical action (similar to disc-like orbits) is 4$\sigma$ overdense compared
to the retrograde counterpart, suggesting that an important fraction of V/EMP
stars reside in the disc plane (Figure 7).

  * •

We associate 5 stars from our follow-up to the most metal-poor stellar stream
C-19, 3 of which are new. One of these stars is 50∘ away from the main body of
the stream, and is the first star belonging to C-19 with positive height above
the disc plane. Adding these bright members of C-19 helps to improve the
constraints on its progenitor and orbital history. We suggest that the stream
might lie farther than what has been reported in the literature (mean distance
of about 21.5 kpc, compared to $\sim$18 kpc reported in Martin et al. 2022b)
while the mean metallicity, velocity dispersion and dynamical width are
consistent with the literature values (Figure 8 and 9).

With this follow-up, we characterise the success rates of the Pristine-Gaia
synthetic catalogue of photometric metallicities and the implications of V/EMP
stars from a chemokinematics point-of-view. Future medium-to-high-resolution
follow-up with chemical abundances and large multi-object spectrographs such
as WEAVE (Jin et al., 2024) and 4MOST (de Jong et al., 2019) results will take
the full chemodynamical analysis of V/EMP stars further, allowing us to study
some of the smallest and earliest galaxies that merged into the Milky Way.

###### Acknowledgements.

AV thanks Ewoud Wempe and Emma Dodd for helpful discussions. ZY, AAA, NFM,
RAI, and SR gratefully acknowledge support from the French National Research
Agency (ANR) funded project ”Pristine” (ANR-18-CE31-0017) along with funding
from the European Research Council (ERC) under the European Unions Horizon
2020 research and innovation programme (grant agreement No. 834148). AAA
acknowledges support from the Herchel Smith Fellowship at the University of
Cambridge and a Fitzwilliam College research fellowship supported by the Isaac
Newton Trust. ES, MB, and MM acknowledge funding through VIDI grant ”Pushing
Galactic Archaeology to its limits” (with project number VI.Vidi.193.093)
which is funded by the Dutch Research Council (NWO). TM was supported by a
Gliese Fellowship at the Zentrum für Astronomie, University of Heidelberg,
Germany. Research by ADD is supported by NSF grant AST-1814208. JMA
acknowledges support from the Agencia Estatal de Investigación del Ministerio
de Ciencia en Innovación (AEIMICIN) and the European Social Fund (ESF+) under
grant PRE2021-100638. CAP acknowledge financial support from the Spanish
Ministry of Science and Innovation (MICINN) project PID2020-117493GB-I00. JMA
acknowledges support from the Agencia Estatal de Investigación del Ministerio
de Ciencia en Innovación (AEI-MICIN) and the European Regional Development
Fund (ERDF) under grant number AYA2017-89076-P, the AEI under grant number
CEX2019-000920-S and the AEI-MICIN under grant number
PID2020-118778GBI00/10.13039/501100011033. DA also acknowledges financial
support from the Spanish Ministry of Science and Innovation (MICINN) under the
2021 Ramón y Cajal program MICINN RYC2021-032609. Based on observations
obtained with MegaPrime/MegaCam, a joint project of CFHT and CEA/DAPNIA, at
the Canada-France-Hawaii Telescope (CFHT) which is operated by the National
Research Council (NRC) of Canada, the Institut National des Sciences de
l’Univers of the Centre National de la Recherche Scientifique of France, and
the University of Hawaii. 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. We are deeply grateful to
the INT support astronomers, Rosa Hoogenboom, and Clár-Bríd Tohill, for
carrying out observation in service mode during Covid-19 times. We also thank
the INT staff for answering all our questions both administrative and
scientific during the observing runs. The INT/IDS is operated on the island of
La Palma by the Isaac Newton Group of Telescopes in the Spanish Observatorio
del Roque de los Muchachos of the Instituto de Astrofísica de Canarias. This
research was supported by the International Space Science Institute (ISSI) in
Bern, through ISSI International Team project 540 (The Early Milky Way). This
research has been partially funded from a Spinoza award by NWO (SPI 78-411).
AV also thanks the availability of the following packages and tools that made
this work possible: vaex (Breddels & Veljanoski, 2018), pandas (Reback et al.,
2022), astropy (Astropy Collaboration et al., 2022), NumPy (Oliphant, 2006;
Van Der Walt et al., 2011), SciPy (Jones et al., 2001), matplotlib (Hunter,
2007), seaborn (Waskom et al., 2016), agama (Vasiliev, 2019), gala (Price-
Whelan, 2017), galpy (Kluyver et al., 2016), pyraf (de La Peña et al., 2001),
iraf (National Optical Astronomy Observatories, 1999), gaiadr3-zeropoint
(Lindegren et al., 2021), JupyterLab (Kluyver et al., 2016), and topcat
(Taylor, 2018).

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##  Appendix A Galactic distribution of V/EMP stars

Figure 10: Mollweide projection of the Galactic coordinates for the V/EMP
stars with [Fe/H]¡-2.5 in the Pristine-Gaia synthetic catalogue, displayed in
gray. The large scale patterns visible in the map are mainly attributed to the
scanning law of the Gaia satellite. The coloured stars represent the subset of
stars observed and analysed in this spectroscopic follow-up colour-coded by
their inferred spectroscopic metallicities (top). The different coloured stars
belong to confined, inner and outer halo stars as defined in section 4.2
(bottom).

Figure 10 top panel shows the sky distribution of stars spectroscopically
followed-up in this work colour-coded by their inferred spectroscopic
metallicities from this work. We do not see any striking feature or
substructure(s) in this view. The distribution of both VMP and EMP stars are
almost random on the sky showing that the Galactic position does not affect
our inferred metallicities. The bottom panel of Figure 10 shows the sky
distribution of confined to the disc plane, inner halo and outer halo stars as
defined in section 4.2 in the same colour scheme. We can see that the confined
stars are almost within 30∘ of the Galactic latitude, while the inner and
outer halo stars are distributed evenly over the sky.

##  Appendix B Additional chemokinematics of the V/EMP stars

Figure 11: (Top) Galactocentric radius versus absolute scale height
distribution of confined, inner and outer halo stars as defined in section
4.2. (Middle) Velocity distribution of these stars in vϕ versus vr space.
(Bottom) Energy versus angular momentum in z-direction for the classified
stars. All EMP stars observable by the spectroscopic follow-up program are
shown in gray.

In Figure 11, we additionally show Galactic distribution and kinematics view
of the confined to the plane, inner halo and outer halo group defined in
section 4.2. The top panel of Figure 11 shows the radius
($\sqrt{x^{2}+y^{2}}$) versus absolute scale height distribution of our
observed V/EMP stars. We see that most of the confined stars have a flattened
distribution over the z-plane and are solar-neighbourhood stars with more
stars towards the inner Galaxy, while the inner and outer halo stars go out to
15-20 kpc in radius and scale height as expected. The middle panel of Figure
11 shows the velocity distribution in azimuth versus radial direction. The
overall distribution looks quite isotropic with a small prograde rotation
(about 20 km/s). The confined stars peak at around 180 km/s with the high
energy stars going out to about 600 km/s. Inner halo stars are have a smaller
spread in their velocities compared to outer halo, as expected. The bottom
panel of Figure 11 shows energy versus angular momentum in the z-direction
(IOM space) for the confined, inner and outer halo stars. We see a clear
difference in the energy distribution between inner and outer halo stars with
inner halo stars sinked deeper into the potential than the outer halo stars.
Confined stars have a mix of prograde and retrograde orbits. One other
interesting note is that all our high energy (unbound) stars belong to the
confined plane suggesting a very high apocentre and small Zmax, which favours
the argument that they are less likely to be hypervelocity stars but their
extreme properties are more likely due to uncertainties in positions and
velocities.

##  Appendix C Planar metal-poor stars in other kinematic spaces

Figure 12: (Top) Action space diamond, the difference between the vertical and
radial component of the action, versus the rotational component of the action
with the axes normalised by Jtot = —Jϕ — + Jr \+ Jz for confined and confined
planar stars as defined in section 4.2 and prograde planar (disc-like) stars
as defined in section 5.1. All EMP stars observable by the spectroscopic
follow-up program are shown in gray (Bottom) Apocentre and Zmax distribution
of these stars shown in same colour.

In Figure 12, we show action diamond space (normalized to the sum of the three
actions of motion) in the top panel and apocentre and Zmax distribution of
confined (Zmax ¡ 3.5 kpc and apocentre ¿ 3.5 kpc) and confined planar
(confined + Zmax/apocentre ¡ 0.2) stars from section 4.2 and prograde planar
(disc-like, prograde box from Figure 7) stars from section 5.1 in the bottom
panel. We can see that the confined stars are close to the plane but not
always in the plane (few circular and radial stars, even) while confined
planar stars are in the plane (prograde or retrograde). We see that the
prograde planar (disc-like) stars is approximately around the same region of
overdensity seen with metal-poor stars moving close to the plane in the action
diamond (prograde and planar) from various spectroscopic surveys as shown in
Fernández-Alvar et al. (2024). The apocentre distribution is close to the
solar neighbourhood and falls off immediately after, favouring the scenario
that these stars could originate from a centrally concentrated, slightly
prograde proto-Galaxy as shown by Zhang et al. (2023); Ardern-Arentsen et al.
(2024). However, it is important to note that our sample size for this
population is too small to say anything conclusive.