# A new reference catalogue for the very metal-poor Universe: +150 OB stars in
Sextans A

M. Lorenzo,1,2 M. Garcia,1 F. Najarro1 A. Herrero,3,4 M. Cerviño,1 N. Castro,5
1Centro de Astrobiología (CAB), CSIC-INTA, Carretera de Ajalvir km 4, 28850
Torrejón de Ardoz, Madrid, Spain
2Departamento de Física Teórica, Universidad Autónoma de Madrid (UAM), Campus
de Cantoblanco, E-28049 Madrid, Spain
3Instituto de Astrofísica de Canarias, La Laguna, Tenerife, Spain
4Departamento de Astrofísica, Universidad de La Laguna, Tenerife, Spain
5Leibniz-Institut für Astrophysik Potsdam (AIP), An der Sternwarte 16, 14482,
Potsdam, Germany E-mail<EMAIL_ADDRESS>
(Accepted XXX. Received YYY; in original form ZZZ)

###### Abstract

Local Group (LG) very metal-poor massive stars are the best proxy for the
First Stars of the Universe and fundamental to modelling the evolution of
early galaxies. These stars may follow new evolutionary pathways restricted to
very low metallicities, such as chemically homogeneous evolution (CHE).
However, given the great distance leap needed to reach very metal-poor
galaxies of the LG and vicinity, no comprehensive spectroscopic studies have
been carried out at metallicities lower than the Small Magellanic Cloud (SMC,
Z = 1/5 Z⊙) until now. After five observing campaigns at the 10.4-m Gran
Telescopio Canarias, we have assembled a low-resolution (R $\sim$ 1000)
spectroscopic collection of more than 150 OB stars in the 1/10 Z⊙ galaxy
Sextans A, increasing by an order of magnitude the number of massive stars
known in this galaxy. The catalogue includes 38 BA-type supergiants, 4 red
supergiants, and the first candidate 1/10 Z⊙ binary systems, CHE sources and
systems hosting stripped stars. The sample massive stars mainly overlap the
higher concentrations of neutral gas of Sextans A. However, we find some
sources in low H i column-density regions. The colour-magnitude diagram of the
galaxy presents large dispersion, which suggests uneven, internal extinction
in Sextans A. This is the largest catalogue of OB-type stars ever produced at
sub-SMC metallicities. This sample constitutes a fundamental first step to
unveiling the evolutionary pathways and fates of very metal-poor massive
stars, analyzing the dependence of radiation-driven winds with metallicity,
and studying binary systems in an environment analogue to the early Universe.

###### keywords:

stars: massive – stars: early-type – stars: Population III – galaxies: stellar
content – galaxies: individual: Sextans A

††pubyear: 2022††pagerange: A new reference catalogue for the very metal-poor
Universe: +150 OB stars in Sextans A–C.5

##  1 Introduction

Massive stars ($M$ > 8 M⊙) are powerful drivers of galaxy evolution. Their
stellar winds, supernova explosions and mighty production of ionizing
radiation stir and impact the surrounding interstellar medium to such an
extent that they dominate the integrated light of star-forming regions.
Massive star properties and spectra are thus essential to interpret evidence
from galaxies at all redshifts, while they are also pivotal ingredients to
model the chemodynamical evolution of galaxies with time.

Yet, very little is known about the evolution of very metal-poor massive stars
($Z$ $\leq$ 1/10 Z⊙), key to interpreting the processes and events of the
early epochs of the Universe. The increasing number of open questions in the
field is demanding a large census of resolved massive stars in environments
with very low metallicity. Such a sample can be used, after quantitative
spectroscopic analysis, to study galaxies at the peak of the cosmic star
formation ($z$ $\sim$ 2, Madau & Dickinson, 2014) or earlier. Among these
questions are the evolutionary pathways of massive stars at metallicities
lower than the Small Magellanic Cloud (Groh et al., 2019; Szécsi et al.,
2022), the incidence of chemically homogeneous evolution (CHE, Szécsi et al.,
2015), CHE stars as possible ionizing sources of the extreme He ii 4686
nebular lines detected in some very metal-poor galaxies (Kehrig et al., 2015;
Sobral et al., 2019; Wofford et al., 2021), and the relative role of radiation
driven winds and binary mass exchange in shaping the evolution of very metal-
poor massive stars (Shenar et al., 2020). In addition, the evolutionary
channels that lead to $\sim$ 30 M⊙ stellar-size black holes are also most
likely related to the extremely low metallicity regimes (Mandel & de Mink,
2016; Marchant et al., 2016; van Son et al., 2020).

The spectroscopic studies of very metal-poor massive stars are greatly
hampered by sensitivity issues. For decades, the Small Magellanic Cloud (SMC,
$Z$ = 1/5 Z⊙) has represented both a metallicity and a distance frontier since
metal-poorer environments are usually found in galaxies at the outer rims of
the Local Group (Garcia et al., 2021). This situation is about to change with
the advent of the Extremely Large Telescope (ELT) and its multi-object
spectrographs HARMONI (Thatte et al., 2016, 2020) and MOSAIC (Hammer et al.,
2014; Evans et al., 2015). They will enable exquisite quantitative
spectroscopic analyses of extragalactic massive stars that so far have been
reserved to the Magellanic Clouds (e.g. Mokiem et al., 2007; Ramírez-Agudelo
et al., 2017; Sabín-Sanjulián et al., 2017). However, the design of both
instruments hinders their use as discoverers of massive stars in nearby very
metal-poor galaxies. The typical diameter of these galaxies is a few
arcminutes, while the field of view of HARMONI is very limited (9” $\times$ 6”
at most) and MOSAIC’s spacing of the robot positioners is too broad (23”
diameter). As a consequence, the construction of catalogues of very metal-poor
massive stars is still restricted to contemporary observing facilities.

In this work, we present the first major effort to mine the population of
massive stars in the galaxy Sextans A (10h11m00.s8 \- 04°41′34′′, dIrr, aka
DDO 75, McConnachie, 2012). Located in the outer edge of the Local Group (LG),
1.34 Mpc away, and presenting low foreground extinction
($\mathrm{E}(B-V)_{\mathrm{fg}}$ = 0.044, Tammann et al., 2011), Sextans A is
a contender to become the next benchmark for low metallicity regimes. Its
oxygen abundance derived from H ii regions ranges 12 + log(O/H) = 7.49–7.71
(1/15-1/10 O⊙, Skillman et al., 1989; Pilyugin, 2001; Kniazev et al., 2005;
Berg et al., 2012) and the analysis of blue supergiant revealed the lowest
abundance of Fe-group elements in the LG ($\langle$[(Fe ii, Cr ii)/H]$\rangle$
= -0.99 $\pm$ 0.06, Kaufer et al., 2004, K04). Its poor Fe content has been
confirmed by Garcia et al. (2017) with HST-COS spectroscopy of a selection of
O stars. In addition, Kaufer et al. (2004) found that the [$\alpha$/Fe]
abundance ratio of Sextans A is consistent with the solar ratio.

With such promising low values for metallicity, this galaxy has sparked the
interest of the community. Camacho et al. (2016, C16) were the first to
confirm the presence of OB-type stars in Sextans A. Garcia et al. (2019, G19)
reported four young O stars on its outskirts, signalling massive star
formation in low-density environments. In addition, Britavskiy et al. (2014);
Britavskiy et al. (2015) unveiled the first red supergiants in the galaxy with
optical spectroscopy of Spitzer/IRAC candidates. Nevertheless, these works
provide a deficient coverage of the most prominent star-forming regions of
Sextans A.

In this work, we present the results of four new observing campaigns at the
10.4-m Gran Telescopio Canarias (GTC) designed to cover all the UV bright
sources of the galaxy and the massive stellar content of its H ii shells.
Together with the sources from our first observing run (Camacho et al., 2016),
the resulting sample will yield first-order information on the evolution of
massive stars with sub-SMC metallicity and provide an extensive collection of
spectra to characterise the physical properties of massive star populations at
these low-$Z$ regimes.

The paper is structured as follows. The spectroscopic observations and adopted
target selection criteria are presented in Sect. 2. In Sect. 3, we describe
the data reduction process. The spectral classification, membership assessment
and radial velocity measurement are described in Sect. 4, along with an
evaluation of the $Q$ pseudo-colour as a photometric method to select OB
candidates. In Sect. 5, we offer an overview of the entire sample and discuss
the colour-magnitude diagram of Sextans A, exploring the implications for its
recent star formation history and stellar evolution at very low metallicity
regimes. Finally, a summary and general conclusions are provided in Sect. 6.

##  2 Description of observations

Observations consisted of two mask multi-object spectroscopy (MOS) observing
runs, GTC3-14AGOS (MOS1, PI A. Herrero) and GTC118-17B (MOS2, PI M. Garcia),
and two long-slit (LSS) observations, GTC43-19B (LSS2, PI M. Garcia) and
GTC58-20B (LSS3, PI M. Lorenzo). All of them were obtained with the Optical
System for Imaging and low-Intermediate-Resolution Integrated Spectroscopy
(OSIRIS) installed at the 10.4-m Gran Telescopio Canarias (GTC). Our data were
complemented with long-slit spectroscopy published in Camacho et al. (2016),
denoted as LSS1. The observing logs are provided in Appendix A.

The MOS1 programme consisted of 13 hours of dark and grey sky and seeing lower
than 1.2”. It was broken into one-hour-long observing blocks distributed in
2014 and 2015, and taken in service mode (see Table 5). In this data set,
targets were selected from Massey et al. (2007)’s photometric catalogue, with
19.0 $\leq$ $V$ $\leq$ 21.0 and a reddening-free pseudo-colour
$Q=(U-B)-0.72(B-V)\leq-0.8$ (see Section 4.4), following the criteria
described in Garcia & Herrero (2013) and Camacho et al. (2016). The
observations of MOS2 consisted of 7 hours divided into two dark nights carried
out in visitor mode (see Table 6). The primary targets of this proposal were
UV sources from Bianchi et al. (2012) with $m_{\mathrm{F170W}}$ $\leq$ 17.6
(HST-WFPC2’s filter F170W, $\lambda_{\mathrm{c}}$ = 1730 Å) and without
previous optical coverage. The long-slit LSS2 observations were designed to
collect all the UV bright sources not covered by the MOS2 run. A total of 4
hours in 2020 via service mode was devoted to the run. Lastly, the LSS3
programme was granted 42 hours of dark sky in visitor mode to mine the
population of massive stars in the giant H ii shells of the southeast of the
galaxy. However, due to bad weather conditions, only 26.2 hours were observed
over the first semester of 2021. Targets with 18.0 $\leq$ $V$ $\leq$ 21.4 and
$Q\leq-0.6$ were included.

Both the MOS and the LSS2 runs used the R2000B VPH filter, 2$\times$2 binning
and 1.2” slits, which granted a resolving power of $R$ =
$\lambda/\Delta\lambda\sim$ 1000 in the 4000 - 5500 Å range. For MOS data
sets, the exact wavelength coverage depends on the position of the slit in the
mask. Some slits (5, 8 and 13 from MOS1, and 2, 20 and 24 from MOS2) were
tilted (80°, 49°, -25°, 5°, 7.1° and 8° from the North-South axis,
respectively) in order to accommodate a larger amount of stars in the sample.
This inclination decreased the spatial and spectral resolution of their
corresponding spectra. Observing run LSS3 was carried out in remote-visitor
mode, and different slit widths could be used to match the seeing conditions,
ranging from 0.8” to 1.5”. This resulted in different spectral resolutions
within the sample. Both the tilt and the width of the slits are noted for the
corresponding stars in Tables LABEL:tab:cat_OB_highQ-LABEL:tab:cat_late.

##  3 Data reduction

The MOS and the LSS data sets were reduced using different procedures. For the
long-slit observations, we followed the routine described in Garcia & Herrero
(2013). The reduction of the two MOS data sets was carried out using the
gtcmos package, an iraf111iraf is distributed by the National Optical
Astronomy Observatory, operated by the Association of Universities for
Research in Astronomy (AURA) under agreement with the National Science
Foundation. pipeline developed by Dr. Divakara Mayya to reduce OSIRIS-MOS
spectra (Gómez-González et al., 2016). The pipeline delivers a 2D wavelength
calibrated image where we applied iraf’s apall task to extract the spectra and
perform background subtraction. Finally, we used our own IDL routines to
perform heliocentric correction, align the spectra in velocity space, coadd
the selected observing blocks, and generate the final spectrum. The spectra
were then normalized with iraf’s continuum task. Further details on relevant
aspects of the reduction steps are given below.

In order to maximise the final quality of the spectra obtained with MOS
observations and to define a reduction protocol for future MOS runs, we tested
three different approaches in the MOS1 run: (i) reducing the data sets
obtained in each observing block (OB) individually; (ii) creating super-
observing blocks (sOB), i.e. combining exposures that were carried out the
same night and shared the calibration images (sOB03 = OB03 + OB04 and sOB13 =
OB12 + OB13 + OB14); and (iii) coadding all raw images prior to reduction. The
latter idea was dismissed due to the spatial shift between the location of the
stars in the data sets of OBs taken in different years, hence coadding all
data would result in severely decreased spatial resolution. A similar shift
between OBs can be detected within the same night if the mask is not
frequently re-aligned. In Table 9, we list the OBs and sOBs used to generate
the final spectrum of each star observed in MOS1. We find that both procedures
return very similar results, although reducing the spectra individually prior
to coaddition provided slightly better results for most of the stars. For the
MOS2 data set, we coadded all exposures of Night-1 and Night-2 separately.

We did not apply Flat Field correction to MOS runs since the corresponding
gtcmos routine would only offer a solution in the common wavelength intervals
covered by all the slits of a given MOS data set. Hence its implementation
would lead to losing the extremes of the spectral range. Our tests showed
little signal-to-noise ratio (S/N) improvement if the Flat Field correction
was applied.

For the reduction of all the observing runs, sky subtraction was carried out
at the step of apall spectral extraction. In this step, we tried to minimize
the nebular contribution in the extracted stellar spectrum to avoid
compromising the spectral classification. To this end, we estimated the
background level from regions as close as possible to the source. This method
was preferred since it preserves spatial information of the nebular emission
at the vicinity of the stars. Yet, nebular subtraction is not always complete
as evinced by the presence of $\rm H\beta$ and [O iii] lines in emission in
the final spectra of the sample stars (Figs. 10-35). In other cases, the
nebular contribution is overcorrected, and the spectra show these lines in
absorption, or in some rare occasions, $\rm H\beta$ in strong absorption and
[O iii] in emission. We note that it is not possible to perform a complete
correction of the nebular emission since it would require to model the
variation of the nebular emissivity along the slit in each case. Moreover, the
spatial resolution of our observations is likely insufficient to adequately
sample the ionization structure of the H ii shells, which translates into an
incomplete nebular subtraction. During the reduction of the long-slit runs, we
tried to improve the sky subtraction by selecting the sky boxes at different
distances from the targets at apall, but the results were again
unsatisfactory.

In MOS observations, we detected dark columns on the right part of some slits
after applying the gtcmos omreduce task. Omreduce trims the portions of the
image corresponding to individual slits, applies the wavelength calibration
and finally returns a reconstructed MOS image with the slits aligned in
wavelength. The dark columns arise when the slits are too close to each other
and the extraction regions assigned by the pipeline overlap. The overlap can
sometimes create an artificial emission that may be mistaken for a stellar
source. At the step of apall extraction, we selected the stellar apertures and
defined their sky boxes avoiding these artefacts. A future recommendation to
prevent this issue is designing the MOS masks with enough space between slits.
We recommend leaving a projected distance larger than 5” in the horizontal
direction if gtcmos will be used for reduction. We also suggest using slits
with a length seven times greater than the expected seeing in the observations
to allow a proper selection of the sky boxes during sky subtraction.

Once the data of all individual observing blocks or super-observing blocks had
been reduced and extracted with apall, all the spectra extracted for a given
star were co-added with a semi-automatic IDL script. This script sets the
spectra to the heliocentric standard of rest and combines the exposures
weighing by their standard deviation. Observing blocks with too poor spectral
quality were discarded.

Some stars were observed in more than one programme and even more than once
within an LSS programme. To improve the S/N of the data, we combined the
observations after matching their spectral resolution and correcting for
radial velocity variations when necessary. Those stars suffering from radial
velocity shifts between exposures are marked in Tables LABEL:tab:cat_OB_highQ-
LABEL:tab:cat_late as "SB1", as they suggest a possible binary nature.

The final reduced spectra are shown in Figs. 10-35 in Appendix B, arranged by
spectral type. Identification numbers, coordinates and photometric
measurements are collected in Tables LABEL:tab:cat_OB_highQ-
LABEL:tab:cat_late, along with the spectral types and radial velocities
derived in this work.

##  4 Analysis

###  4.1 Spectral classification

The spectral classification of the OB-type stars was carried out following the
criteria outlined in Castro et al. (2008) for metal-poor regimes ($Z$ $\sim$
0.5 Z⊙), adapted to the metallicity of Sextans A by Camacho (2017). We also
used the guidelines given by Sota et al. (2011) for late O-types and the ones
from Walborn et al. (2002) for stars earlier than O5. The main diagnostics for
O stars are based on the relative intensity of He i and He ii lines. The
luminosity class of O-type stars was assigned based on the width of the Balmer
lines and on whether the He ii 4686 line was in absorption or emission. We
note the latter diagnostic may be affected by the weak winds of the targets.
The spectral types of B stars are indicated by the presence and strength of Si
transitions in different ionization stages, Mg ii 4481 and He i lines. Their
luminosity class was assigned from the width of the Balmer lines.

We also used the absolute magnitude $M_{\mathrm{V}}$ as an additional
indicator of luminosity class for some sources. We estimated $M_{\mathrm{V}}$
using the distance modulus of the galaxy ($\mu_{0}$ = 25.63, Tammann et al.,
2011) and the intrinsic colour $(B-V)_{\rm{0}}$ calculated from the pseudo-
colour $Q$ ($(B-V)_{0}=-0.005+0.317Q$, Massey et al., 2000).

Later spectral types (A-types and later) were classified following Gray &
Corbally (2009). The spectral classification of these late stars may be
affected by the low metallicity of the environment since metallic lines and
molecular bands are the main diagnostics.

Spectral types are compiled in Tables LABEL:tab:cat_OB_highQ-
LABEL:tab:cat_late. In addition, we provide brief notes on the spectral
classification of targets in Appendix C.

Table 1: Blue type stars (from O to F types) confirmed by spectroscopy in this work. Identification codes (ID), spectral types (SpT), radial velocities ($v_{\mathrm{rad}}$) and signal-to-noise ratio per element of resolution (S/N, calculated in a region free of spectral lines around 4800 Å) are provided. Stars whose radial velocity was determined manually are marked with (∗). Identification numbers, coordinates and photometry by Massey et al. (2007) are also included when available, otherwise we provide the ones from Bianchi et al. (2012)'s or Holtzman et al. (2006)'s catalogues, and mark the sources with (†) or (‡), respectively. We include additional information in column Notes. ID | Alt. ID | RA (J2000.0) | DEC (J2000.0) | SpT | $V$ | $B-V$ | $Q$ |  $v_{\mathrm{rad}}$ (km s-1) | S/N | Notes
---|---|---|---|---|---|---|---|---|---|---
s001*‡ | - | 10:11:05.29 | -04:42:41.7 | O3.5 V | F555W = 21.851 | - | - |  460 $\pm$ 100 | 55 | 1.0” slit. SB2?
s002 | J101105.28-044238.3 | 10:11:05.28 | -04:42:38.3 | O4 Vz | 20.369 | -0.279 | -0.937 |  275 $\pm$ 30 | 111 | 1.0” slit.
s003* | J101058.59-044328.9 | 10:10:58.59 | -04:43:28.9 | O3-O5 Vz | 20.804 | -0.095 | -1.120 |  425 $\pm$ 60 | 83 |  G19-s2
s004 | J101057.89-044310.2 | 10:10:57.89 | -04:43:10.2 | O5 III | 20.917 | -0.277 | -1.006 |  310 $\pm$ 30 | 91 |  G19-s4. Bin?
s005 | J101053.95-044057.4 | 10:10:53.95 | -04:40:57.4 | O5 V((fc)) | 21.751 | -0.364 | -0.800 |  375 $\pm$ 40 | 30 | 0.8” slit.
s006* | J101104.74-044206.1 | 10:11:04.74 | -04:42:06.1 | O6 I | 20.721 | -0.068 | -0.416 |  500 $\pm$ 50 | 29 |
s007 | J101106.48-044237.1 | 10:11:06.48 | -04:42:37.1 | O6 V | 20.515 | -0.280 | -0.965 |  395 $\pm$ 35 | 105 | SB2?
s008* | J101106.65-044213.0 | 10:11:06.65 | -04:42:13.0 | O6.5 V | 21.071 | -0.199 | -0.985 |  300 $\pm$ 50 | 71 | 1.2” slit. Blend.
s009 | J101108.48-044149.2 | 10:11:08.48 | -04:41:49.2 | O7 V | 20.659 | -0.285 | -0.948 |  320 $\pm$ 35 | 66 |
s010* | J101059.19-044308.2 | 10:10:59.19 | -04:43:08.2 | O7 V | 21.003 | -0.134 | -0.549 |  250 $\pm$ 50 | 30 | 1.5” slit.
s011 | J101103.70-044234.2 | 10:11:03.70 | -04:42:34.2 | O7.5 III + B0 I | 20.191 | -0.220 | -1.019 |  315 $\pm$ 30 | 80 | 1.0” slit. Blend.
s012 | J101054.23-044115.3 | 10:10:54.23 | -04:41:15.3 | O7.5 III | 20.115 | -0.181 | -0.835 |  290 $\pm$ 155 | 52 | SB2 or blend?
s013* | J101100.09-044336.0 | 10:11:00.09 | -04:43:36.0 | O7.5 III | 21.532 | -0.295 | -0.914 |  370 $\pm$ 50 | 43 |
s014 | J101053.81-044113.0 | 10:10:53.81 | -04:41:13.0 | O7.5 III((f)) | 20.688 | -0.265 | -1.015 |  330 $\pm$ 35 | 49 |  C16-OB326. Off-slit.
s015* | J101105.47-044236.8 | 10:11:05.47 | -04:42:36.8 | O7.5 V + B? | 22.259 | -0.410 | -1.010 |  225 $\pm$ 50 | 35 | 1.0” slit. SB2?
s016 | J101056.28-044253.0 | 10:10:56.28 | -04:42:53.0 | O8 II | 20.639 | -0.192 | -1.065 |  285 $\pm$ 95 | 72 |
s017* | J101106.50-044139.0 | 10:11:06.50 | -04:41:39.0 | O8 II | 21.399 | 0.198 | -1.281 |  330 $\pm$ 50 | 18 | Off-slit.
s018* | J101053.61-044246.3 | 10:10:53.61 | -04:42:46.3 | O8 III | 21.092 | -0.321 | -0.894 |  210 $\pm$ 50 | 41 |
s019† | 3485 | 10:11:05.84 | -04:42:12.3 | O8 III | 21.930 | -0.250 | -0.850 |  270 $\pm$ 35 | 32 | 0.8” slit. Blend.
s020 | J101107.33-044205.6 | 10:11:07.33 | -04:42:05.6 | O8 V | 20.544 | -0.196 | -0.968 |  315 $\pm$ 35 | 125 | 1.5” slit. SB2 or blend?
s021 | J101104.79-044220.9 | 10:11:04.79 | -04:42:20.9 | O8 V | 20.598 | -0.468 | -0.616 |  335 $\pm$ 35 | 95 | 0.63” slit. Tilted 3°. Blend.
s022 | J101105.38-044240.1 | 10:11:05.38 | -04:42:40.1 | O8 V | 19.459 | -0.259 | -0.950 |  370 $\pm$ 30 | 95 |  C16-OB521. Blend.
s023 | J101105.21-044240.5 | 10:11:05.21 | -04:42:40.5 | O8 V | 20.288 | -0.299 | -0.927 |  235 $\pm$ 35 | 82 | 1.0” slit.
s024 | J101105.90-044210.3 | 10:11:05.90 | -04:42:10.3 | O8 V | 20.810 | -0.162 | -1.044 |  280 $\pm$ 35 | 77 |
s025 | J101110.14-044129.7 | 10:11:10.14 | -04:41:29.7 | O8 V | 21.060 | -0.265 | -0.976 |  350 $\pm$ 45 | 52 |
s026† | 3935 | 10:11:07.43 | -04:42:04.5 | O8 V | 22.500 | -0.060 | -1.007 |  325 $\pm$ 110 | 51 | 0.8” slit. Blend.
s027† | 2398 | 10:11:03.39 | -04:42:37.1 | O8 V | 21.960 | -0.220 | -0.852 |  280 $\pm$ 35 | 20 | 1.5” slit.
s028 | J101107.34-044231.7 | 10:11:07.34 | -04:42:31.7 | O8.5 III | 20.568 | -0.228 | -1.033 |  300 $\pm$ 35 | 83 |
s029 | J101058.19-044318.4 | 10:10:58.19 | -04:43:18.4 | O8.5 III | 20.803 | -0.247 | -1.005 |  285 $\pm$ 35 | 74 |  G19-s3. SB1.
s030 | J101109.39-044055.5 | 10:11:09.39 | -04:40:55.5 | O8.5 III | 21.008 | -0.322 | -0.848 |  380 $\pm$ 35 | 56 | Tilted 7.1°.
s031* | J101057.33-044016.6 | 10:10:57.33 | -04:40:16.6 | O8.5 V | 21.229 | -0.181 | -1.069 |  475 $\pm$ 50 | 115 | Tilted 8°.
s032 | J101109.35-044057.8 | 10:11:09.35 | -04:40:57.8 | O8.5 V | 21.216 | -0.326 | -0.776 |  365 $\pm$ 30 | 67 | Tilted 7.1°. Blend.
s033 | J101105.64-044215.6 | 10:11:05.64 | -04:42:15.6 | O8.5 V | 21.654 | -0.271 | -0.960 |  295 $\pm$ 75 | 40 | 0.8” slit.
s034 | J101058.05-044021.8 | 10:10:58.05 | -04:40:21.8 | O8.5 V | 20.923 | -0.168 | -0.548 |  250 $\pm$ 30 | 26 | 1.5” slit.
s035* | J101106.50-044212.4 | 10:11:06.50 | -04:42:12.4 | O8.5 V | 20.683 | -0.089 | -0.498 |  170 $\pm$ 100 | 24 | Blend.
s036 | J101105.07-044214.6 | 10:11:05.07 | -04:42:14.6 | O9 I | 19.738 | -0.237 | -0.931 |  360 $\pm$ 30 | 200 |
s037 | J101104.78-044224.1 | 10:11:04.78 | -04:42:24.1 | O9 I | 20.681 | -0.234 | -1.001 |  345 $\pm$ 30 | 110 |  C16-OB623.
s038 | J101106.05-044211.4 | 10:11:06.05 | -04:42:11.4 | O9 I((f)) | 19.492 | -0.231 | -1.013 |  450 $\pm$ 35 | 91 |  C16-OB523.
s039 | J101059.96-044332.9 | 10:10:59.96 | -04:43:32.9 | O9 III | 21.084 | -0.171 | -0.502 |  340 $\pm$ 35 | 49 | 1.0” slit.
s040* | J101106.56-044217.1 | 10:11:06.56 | -04:42:17.1 | O9 III + mid-B? | 21.285 | -0.311 | -0.900 |  250 $\pm$ 30 | 29 | SB2?
s041 | J101058.53-044414.4 | 10:10:58.53 | -04:44:14.4 | O9.5 II | 20.877 | -0.029 | -1.071 |  290 $\pm$ 30 | 73 |  G19-s1
s042 | J101105.44-044221.4 | 10:11:05.44 | -04:42:21.4 | O9.5 III | 20.642 | -0.226 | -0.939 |  275 $\pm$ 30 | 108 | Blend.
s043 | J101105.03-044224.8 | 10:11:05.03 | -04:42:24.8 | O9.5 V | 20.863 | 0.250 | -0.607 |  305 $\pm$ 30 | 92 | 1.5” slit. Blend.
s044 | J101110.45-044132.9 | 10:11:10.45 | -04:41:32.9 | O9.5 V | 20.800 | -0.255 | -0.966 |  270 $\pm$ 95 | 60 | SB2?
s045 | J101105.93-044208.0 | 10:11:05.93 | -04:42:08.0 | O9.5 V | 20.511 | -0.169 | -0.430 |  250 $\pm$ 35 | 60 | 0.8” slit.
s046* | J101104.87-044216.0 | 10:11:04.87 | -04:42:16.0 | O9.5 V | 20.748 | -0.214 | -0.872 |  325 $\pm$ 30 | 45 | 0.8” slit. Blend.
s047 | J101104.58-044213.0 | 10:11:04.58 | -04:42:13.0 | O9.5 V | 20.668 | 0.267 | -0.274 |  335 $\pm$ 100 | 31 |  C16-OB625.
s048*† | 3691 | 10:11:06.34 | -04:42:28.2 | O9.5 Vn | 22.180 | -0.240 | -0.797 |  200 $\pm$ 60 | 20 | 1.0” slit.
s049 | J101105.30-044210.1 | 10:11:05.30 | -04:42:10.1 | O9.7 I | 19.976 | -0.268 | -0.929 |  350 $\pm$ 30 | 91 | SB1? Blend.
s050 | J101100.66-044044.3 | 10:11:00.66 | -04:40:44.3 | O9.7 I | 19.609 | -0.248 | -0.960 |  295 $\pm$ 30 | 90 |  C16-OB321. 0.63” slit.
s051 | J101058.27-044230.6 | 10:10:58.27 | -04:42:30.6 | O9.7 I | 20.774 | -0.277 | -0.875 |  335 $\pm$ 30 | 70 |
s052 | J101106.46-044207.3 | 10:11:06.46 | -04:42:07.3 | O9.7 I | 20.320 | -0.166 | -0.690 |  345 $\pm$ 30 | 62 | SB2?
s053* | J101106.13-044208.3 | 10:11:06.13 | -04:42:08.3 | O9.7 I | 21.255 | -0.269 | -0.904 |  325 $\pm$ 50 | 60 |
s054 | J101104.53-044219.3 | 10:11:04.53 | -04:42:19.3 | O9.7 III | 20.511 | -0.202 | -0.951 |  330 $\pm$ 30 | 125 |
s055 | J101053.90-044111.0 | 10:10:53.90 | -04:41:11.0 | O9.7 III | 20.323 | -0.103 | -0.920 |  300 $\pm$ 30 | 103 | Blend.
s056 | J101100.94-044003.2 | 10:11:00.94 | -04:40:03.2 | O9.7 V | 20.914 | -0.147 | -0.922 |  335 $\pm$ 30 | 98 | Oe?
s057 | J101059.58-044135.4 | 10:10:59.58 | -04:41:35.4 | O9.7 V | 20.697 | -0.110 | -0.393 |  355 $\pm$ 35 | 59 |
s058* | J101104.05-044226.1 | 10:11:04.05 | -04:42:26.1 | O9.7 V + B? | 20.957 | -0.204 | -0.932 |  350 $\pm$ 50 | 40 | SB2?
s059 | - | 10:10:57.09 | -04:44:46.3 | O9.7 Vn | - | - | - |  355 $\pm$ 30 | 36 |
s060* | J101107.27-044032.4 | 10:11:07.27 | -04:40:32.4 | O9.7 V | 21.222 | -0.065 | -1.121 |  400 $\pm$ 100 | 26 |  Blend or bin? Broad He i.
s061 | J101057.32-044055.4 | 10:10:57.32 | -04:40:55.4 | B0 I | 19.915 | -0.247 | -0.899 |  290 $\pm$ 30 | 159 |
s062 | J101103.80-044212.8 | 10:11:03.80 | -04:42:12.8 | B0 I | 20.358 | -0.269 | -0.872 |  315 $\pm$ 30 | 76 |
s063 | J101107.43-044025.4 | 10:11:07.43 | -04:40:25.4 | B0 III | 20.727 | -0.295 | -0.885 |  305 $\pm$ 35 | 133 |
s064* | J101108.49-043905.9 | 10:11:08.49 | -04:39:05.9 | B0 III | 21.223 | -0.268 | -0.879 |  325 $\pm$ 50 | 44 | SB2?
s065* | J101107.31-044205.3 | 10:11:07.31 | -04:42:05.3 | B0 V | 20.889 | -0.601 | -0.595 |  100 $\pm$ 50 | 63 | SB2? Be?
s066 | J101058.01-044027.8 | 10:10:58.01 | -04:40:27.8 | B0.5 I | 20.041 | -0.225 | -0.837 |  310 $\pm$ 40 | 124 |
s067 | J101059.29-044343.8 | 10:10:59.29 | -04:43:43.8 | B0.5 I | 20.342 | -0.137 | -0.464 |  360 $\pm$ 40 | 82 |
s068* | J101106.48-044210.5 | 10:11:06.48 | -04:42:10.5 | B0.5 I | 20.521 | 0.272 | -0.344 |  0 $\pm$ 50 | 69 | = s093. 0.8” slit. Bin. Blend.
s069 | J101107.59-044159.9 | 10:11:07.59 | -04:41:59.9 | B0.5 II + O? | 19.977 | -0.214 | -0.941 |  390 $\pm$ 40 | 109 | SB2?
s070 | J101106.03-044209.1 | 10:11:06.03 | -04:42:09.1 | B0.5 III | 20.834 | -0.272 | -0.907 |  270 $\pm$ 45 | 82 |  C16-OB524
s071 | J101105.69-044213.6 | 10:11:05.69 | -04:42:13.6 | B1 I | 19.699 | -0.236 | -0.903 |  265 $\pm$ 35 | 256 | Tilted 49°. SB1?
s072 | J101059.21-043948.1 | 10:10:59.21 | -04:39:48.1 | B1 I | 19.329 | -0.192 | -0.880 |  340 $\pm$ 30 | 211 |  C16-OB421. SB1?
s073 | J101103.98-044232.5 | 10:11:03.98 | -04:42:32.5 | B1 I | 20.069 | -0.199 | -0.790 |  285 $\pm$ 30 | 101 | 1.5” slit.
s074 | J101054.62-044103.0 | 10:10:54.62 | -04:41:03.0 | B1 I | 19.739 | -0.218 | -0.942 |  330 $\pm$ 30 | 42 |  C16-OB422.
s075* | J101104.02-044222.9 | 10:11:04.02 | -04:42:22.9 | B1 I + O? | 21.220 | -0.023 | -0.107 |  200 $\pm$ 50 | 34 | 0.8” slit. SB2?
s076† | 2990 | 10:11:04.66 | -04:42:23.6 | B1 I | 21.940 | -0.070 | -0.920 |  335 $\pm$ 40 | 29 |
s077 | J101054.08-044111.5 | 10:10:54.08 | -04:41:11.5 | B1 III | 19.487 | -0.217 | -0.875 |  340 $\pm$ 30 | 92 |  C16-OB323. SB1. Blend.
s078 | J101102.18-044315.4 | 10:11:02.18 | -04:43:15.4 | B1 III | 21.290 | -0.137 | -0.915 |  290 $\pm$ 30 | 54 |
s079* | J101105.82-044213.0 | 10:11:05.82 | -04:42:13.0 | B1 III | 21.774 | -0.163 | -0.984 |  325 $\pm$ 50 | 50 | Blend.
s080* | J101054.97-044116.2 | 10:10:54.97 | -04:41:16.2 | B1 III | 21.471 | -0.292 | -0.983 |  250 $\pm$ 50 | 34 |
s081† | 3646 | 10:11:06.21 | -04:42:25.7 | B1 V | 23.250 | 0.000 | -1.260 |  320 $\pm$ 40 | 44 | Blend.
s082 | J101105.65-044038.9 | 10:11:05.65 | -04:40:38.9 | B1.5 I | 19.896 | -0.186 | -0.822 |  400 $\pm$ 30 | 160 |
s083 | J101055.35-044106.1 | 10:10:55.35 | -04:41:06.1 | B2 I | 19.067 | -0.171 | -1.007 |  280 $\pm$ 30 | 183 |  C16-OB222. SB1.
s084 | J101104.33-044220.8 | 10:11:04.33 | -04:42:20.8 | B2 I | 20.552 | 0.061 | -1.017 |  265 $\pm$ 35 | 72 | 0.8” slit.
s085 | J101104.69-044228.8 | 10:11:04.69 | -04:42:28.8 | B2 I | 20.850 | 0.254 | - |  320 $\pm$ 30 | 62 | 1.5” slit. SB1?
s086* | J101104.97-044218.9 | 10:11:04.97 | -04:42:18.9 | B2 III | 20.509 | 0.107 | 0.014 |  100 $\pm$ 50 | 126 | SB1?
s087 | J101053.94-044110.1 | 10:10:53.94 | -04:41:10.1 | B2 III + O? | 20.526 | -0.230 | -0.921 |  255 $\pm$ 35 | 79 | Blend.
s088* | J101110.46-044131.5 | 10:11:10.46 | -04:41:31.5 | B2 III + O? | 21.265 | -0.265 | -0.435 |  325 $\pm$ 50 | 29 | 0.8” slit. SB2?
s089 | J101102.38-044014.6 | 10:11:02.38 | -04:40:14.6 | B2.5 I | 19.581 | -0.099 | -0.683 |  340 $\pm$ 30 | 150 |  C16-OB622. Tilted -25°. Bin?.
s090 | J101106.11-044229.8 | 10:11:06.11 | -04:42:29.8 | B2.5 I | 20.358 | -0.133 | -0.507 |  295 $\pm$ 30 | 87 | 1.0” slit. Blend.
s091 | J101058.27-044257.8 | 10:10:58.27 | -04:42:57.8 | B2.5 I | 19.176 | -0.074 | -0.706 |  330 $\pm$ 30 | 83 |  C16-OB221.
s092 | J101105.50-044128.8 | 10:11:05.50 | -04:41:28.8 | B2.5 I | 20.155 | -0.141 | -0.671 |  325 $\pm$ 30 | 77 |
s093* | J101106.48-044210.5 | 10:11:06.48 | -04:42:10.5 | B2.5 I | 20.521 | 0.272 | -0.344 |  450 $\pm$ 50 | 51 | = s068. 0.8” slit. Bin. Blend.
s094* | J101058.58-044423.5 | 10:10:58.58 | -04:44:23.5 | B2.5 I + O? | 21.621 | 0.205 | -0.966 |  335 $\pm$ 50 | 37 | Off-slit. Tilted 5°. SB2?
s095 | J101105.96-044221.4 | 10:11:05.96 | -04:42:21.4 | B2.5 III | 21.672 | -0.188 | -0.730 |  320 $\pm$ 30 | 47 |
s096* | J101056.64-044102.1 | 10:10:56.64 | -04:41:02.1 | B2.5 III | 21.062 | 0.263 | -0.056 |  200 $\pm$ 50 | 34 | 0.8” slit.
s097 | J101107.62-044203.6 | 10:11:07.62 | -04:42:03.6 | B3 I | 20.085 | -0.175 | -0.662 |  300 $\pm$ 30 | 158 | O.8” slit.
s098 | J101106.12-044219.9 | 10:11:06.12 | -04:42:19.9 | B3 I | 19.782 | -0.163 | -0.789 |  335 $\pm$ 30 | 136 |
s099 | J101102.75-044047.6 | 10:11:02.75 | -04:40:47.6 | B3 I | 20.201 | -0.141 | -0.657 |  275 $\pm$ 40 | 121 |
s100 | J101051.62-044326.4 | 10:10:51.62 | -04:43:26.4 | B3 III | 20.997 | -0.263 | -0.966 |  320 $\pm$ 35 | 62 |
s101 | J101054.45-044114.1 | 10:10:54.45 | -04:41:14.1 | B3 III | 21.254 | -0.328 | -0.897 |  255 $\pm$ 35 | 43 |
s102 | J101104.59-044222.2 | 10:11:04.59 | -04:42:22.2 | B7 I | 20.189 | -0.074 | -0.463 |  280 $\pm$ 35 | 94 |
s103* | J101103.93-044214.0 | 10:11:03.93 | -04:42:14.0 | B7 I | 20.437 | -0.069 | -0.386 |  400 $\pm$ 40 | 93 | Blend.
s104 | J101105.41-044210.9 | 10:11:05.41 | -04:42:10.9 | B7 I + O? | 21.153 | -0.206 | -0.751 |  250 $\pm$ 40 | 90 | 1.5” slit. Blend?
s105 | J101104.42-044232.4 | 10:11:04.42 | -04:42:32.4 | B7 I | 19.227 | -0.075 | -0.554 |  270 $\pm$ 35 | 86 |
s106 | J101106.02-044214.2 | 10:11:06.02 | -04:42:14.2 | B7 II | 18.771 | 0.054 | -0.277 |  250 $\pm$ 30 | 102 |  C16-OB121. SB1. Blend.
s107* | J101059.81-044305.5 | 10:10:59.81 | -04:43:05.5 | B7 III | 20.575 | 0.126 | -0.011 |  200 $\pm$ 100 | 77 | 1.5” slit.
s108 | J101105.86-044232.0 | 10:11:05.86 | -04:42:32.0 | B7 III + B? | 21.560 | -0.119 | -0.464 |  320 $\pm$ 55 | 44 | 1.0” slit. SB2?
s109* | J101107.14-044131.4 | 10:11:07.14 | -04:41:31.4 | B7 III + O? | 21.538 | -0.268 | -0.873 |  350 $\pm$ 50 | 28 | SB2?
s110 | J101105.98-044217.9 | 10:11:05.98 | -04:42:17.9 | B8 I | 20.008 | -0.048 | -0.358 |  260 $\pm$ 55 | 111 | 1.5” slit.
s111 | J101059.13-044051.0 | 10:10:59.13 | -04:40:51.0 | B8 II | 20.486 | -0.069 | -0.412 |  340 $\pm$ 35 | 54 |  C16-OB324. Off-slit.
s112* | J101104.08-044224.2 | 10:11:04.08 | -04:42:24.2 | B8 III | 21.104 | 0.038 | -0.144 |  270 $\pm$ 30 | 67 | 0.8” slit.
s113 | J101059.15-044215.1 | 10:10:59.15 | -04:42:15.1 | B8 III | 21.514 | -0.111 | -0.281 |  265 $\pm$ 35 | 38 | 1.0” slit.
s114 | J101056.30-044105.7 | 10:10:56.30 | -04:41:5.70 | B8 III | 21.478 | -0.228 | -0.719 |  340 $\pm$ 50 | 23 | Blend.
s115 | J101105.17-044236.0 | 10:11:05.17 | -04:42:36.0 | B9 I | 19.360 | -0.044 | -0.427 |  295 $\pm$ 250 | 95 | 0.8” slit.
s116 | J101054.89-044112.1 | 10:10:54.89 | -04:41:12.1 | B9 I | 18.480 | -0.010 | -0.529 |  375 $\pm$ 30 | 75 |  C16-OB122. Blend.
s117 | J101058.86-044126.2 | 10:10:58.86 | -04:41:26.2 | B9 III | 20.169 | 0.020 | -0.040 |  325 $\pm$ 45 | 46 | 0.8” slit.
s118*† | 3713 | 10:11:06.39 | -04:42:16.4 | B9 V | 21.980 | -0.120 | -0.184 |  300 $\pm$ 100 | 35 |
s119* | J101106.38-044155.9 | 10:11:06.38 | -04:41:55.9 | A0 II | 20.519 | -0.043 | -0.155 |  350 $\pm$ 40 | 41 |  C16-OB525.
s120 | J101107.24-044219.7 | 10:11:07.24 | -04:42:19.7 | A2 I | 18.873 | 0.081 | -0.009 |  370 $\pm$ 30 | 114 |
s121* | J101054.71-044115.6 | 10:10:54.71 | -04:41:15.6 | A2 I | 19.211 | 0.129 | -0.019 |  270 $\pm$ 30 | 83 |
s122 | J101057.50-044208.9 | 10:10:57.50 | -04:42:08.9 | A2 I | 19.499 | -0.011 | -0.259 |  275 $\pm$ 40 | 60 | 1.0” slit.
s123 | J101056.97-044206.6 | 10:10:56.97 | -04:42:06.6 | A2 III | 20.449 | 0.172 | -0.017 |  320 $\pm$ 40 | 32 | 1.0” slit.
s124* | J101104.97-044233.8 | 10:11:04.97 | -04:42:33.8 | F5 I | 17.477 | 0.174 | -0.188 |  350 $\pm$ 40 | 132 |  C16-OB621. K04-SexA-513.
s125*† | J10111837-0440480 | 10:11:18.37 | -04:40:48.1 | F8 V | - | - | - |  0 $\pm$ 30 | 74 | Foreground.
Table 1: continued. Table 2: Blue type stars for which a spectral subtype could not be assigned. The description of columns is the same as Table LABEL:tab:cat_OB_highQ. ID | Alt. ID | RA (J2000.0) | DEC (J2000.0) | SpT | $V$ | $B-V$ | $Q$ |  $v_{\mathrm{rad}}$ (km s-1) | S/N | Notes
---|---|---|---|---|---|---|---|---|---|---
s126*† | 3545 | 10:11:06.00 | -04:42:28.4 | O V | B = 23.04 | - | - |  230 $\pm$ 50 | 48 | Blend.
s127* | J101100.30-044330.7 | 10:11:00.30 | -04:43:30.7 | O V + gal | 21.398 | -0.071 | -0.263 |  250 $\pm$ 100 | 84 |
s128* | - | 10:11:06.64 | -04:41:33.5 | O V | - | - | - |  235 $\pm$ 100 | 8 |
s129 | J101104.28-044238.2 | 10:11:04.28 | -04:42:38.2 | O V + neb | 21.237 | -0.310 | -1.000 |  330 $\pm$ 30 | 15 |
s130 | J101106.30-044140.8 | 10:11:06.30 | -04:41:40.8 | O | 20.934 | -0.075 | -0.473 |  380 $\pm$ 45 | 15 | Off-slit.
s131 | J101106.32-043914.0 | 10:11:06.32 | -04:39:14.0 | O + neb | 21.277 | -0.140 | -1.003 |  390 $\pm$ 35 | 31 |
s132* | J101101.33-044150.9 | 10:11:01.33 | -04:41:50.9 | O V + neb | 21.347 | -0.093 | -1.017 |  490 $\pm$ 70 | 23 | Off-slit.
s133 | J101100.17-043939.2 | 10:11:00.17 | -04:39:39.2 | O + neb | 21.219 | -0.098 | -1.038 |  225 $\pm$ 30 | 74 |
s134* | J101101.20-044337.2 | 10:11:01.20 | -04:43:37.2 | O V | 21.097 | -0.056 | -0.447 |  325 $\pm$ 50 | 33 |
s135* | J101105.28-044223.4 | 10:11:05.28 | -04:42:23.4 | O V | 20.689 | 0.011 | -0.239 |  325 $\pm$ 50 | 50 | 1.5” slit.
s136 | J101109.27-044147.7 | 10:11:09.27 | -04:41:47.7 | O + neb | 20.639 | -0.062 | -1.111 |  435 $\pm$ 80 | 31 |
s137* | J101108.58-044125.7 | 10:11:08.58 | -04:41:25.7 | Be + O? | 21.007 | -0.096 | -1.108 |  325 $\pm$ 50 | 63 | SB2?
s138* | J101100.56-044327.6 | 10:11:00.56 | -04:43:27.6 | B III + gal | 20.406 | 0.026 | -0.513 |  175 $\pm$ 50 | 69 | 1.0” slit. SB1?
s139*† | 3988 | 10:11:07.24 | -04:41:50.3 | B + O? | 22.190 | -0.180 | -0.820 |  200 $\pm$ 50 | 18 | SB2?
s140* | J101105.48-044202.9 | 10:11:05.48 | -04:42:2.90 | B III | 20.907 | -0.266 | -0.837 |  250 $\pm$ 50 | 15 | Off-slit.
s141 | J101059.21-044348.8 | 10:10:59.21 | -04:43:48.8 | B | 21.159 | -0.138 | -0.576 |  255 $\pm$ 85 | 16 | Off-slit.
s142 | J101055.14-044117.0 | 10:10:55.14 | -04:41:17.0 | B | 21.422 | 0.230 | -0.914 |  315 $\pm$ 50 | 26 | Blend.
s143 | - | 10:11:06.17 | -04:42:13.7 | B | - | - | - |  270 $\pm$ 40 | 27 | 0.8” slit.
s144 | J101105.90-044155.8 | 10:11:05.90 | -04:41:55.8 | B | 21.535 | 0.043 | 0.069 |  305 $\pm$ 40 | 14 | Off-slit.
s145*† | 2432 | 10:11:03.50 | -04:42:27.4 | B | 22.380 | 1.020 | - |  325 $\pm$ 50 | 24 | 0.8” slit.
s146 | J101101.76-044314.0 | 10:11:01.76 | -04:43:14.0 | B | 21.678 | -0.006 | -0.306 |  345 $\pm$ 35 | 18 |
s147* | J101101.90-044241.1 | 10:11:01.90 | -04:42:41.1 | B | 21.001 | -0.095 | -0.393 |  150 $\pm$ 50 | 23 | 0.8” slit.
s148* | J101101.93-044147.1 | 10:11:01.93 | -04:41:47.1 | OB | 20.904 | 0.338 | -0.216 |  325 $\pm$ 100 | 40 | 0.8” slit.
s149* | J101103.95-044225.0 | 10:11:03.95 | -04:42:25.0 | OB | 21.163 | 0.684 | - |  325 $\pm$ 100 | 36 | 0.8” slit.
s150*† | 3416 | 10:11:05.71 | -04:42:10.6 | OB | 22.190 | -0.260 | -0.793 |  400 $\pm$ 100 | 32 | 0.8” slit.
s151 | J101105.55-044221.9 | 10:11:05.55 | -04:42:21.9 | OB | 20.376 | -0.002 | -0.062 |  405 $\pm$ 85 | 29 | Blend.
s152*† | 2392 | 10:11:03.37 | -04:42:28.9 | OB | 21.590 | 1.190 | - |  325 $\pm$ 50 | 28 | 0.8” slit.
s153* | J101103.34-044156.2 | 10:11:03.34 | -04:41:56.2 | OB | 21.589 | -0.164 | -0.440 |  245 $\pm$ 100 | 26 | 0.8” slit.
s154* | J101100.55-044352.8 | 10:11:00.55 | -04:43:52.8 | OB | 21.314 | -0.115 | -0.234 |  450 $\pm$ 50 | 22 |
s155* | J101104.54-044229.7 | 10:11:04.54 | -04:42:29.7 | OB | 22.006 | -0.212 | -0.142 |  325 $\pm$ 50 | 22 | 0.8” slit.
s156*† | 3509 | 10:11:05.95 | -04:41:54.0 | OB | 22.180 | -0.170 | -0.848 |  205 $\pm$ 50 | 16 |
s157* | - | 10:11:07.12 | -04:42:09.6 | OB | - | - | - |  260 $\pm$ 100 | 13 | 1.5” slit.
s158* | J101057.38-044019.0 | 10:10:57.38 | -04:40:19.0 | OBA | 21.217 | -0.075 | -0.284 |  460 $\pm$ 50 | 34 | Tilted 8°.
s159* | J101105.64-044224.0 | 10:11:05.64 | -04:42:24.0 | OBA | 21.023 | -0.204 | -0.528 |  325 $\pm$ 100 | 32 | Blend.
s160* | J101101.83-044117.6 | 10:11:01.83 | -04:41:17.6 | OBA | 20.706 | -0.032 | 0.039 |  160 $\pm$ 50 | 27 | 1.5” slit.
s161 | J101108.21-044153.5 | 10:11:08.21 | -04:41:53.5 | OBA | 21.899 | -0.153 | -1.048 |  440 $\pm$ 55 | 25 |
s162 | J101105.15-044212.8 | 10:11:05.15 | -04:42:12.8 | OBA | 21.676 | -0.080 | -0.357 |  475 $\pm$ 230 | 24 |
s163* | J101103.49-044157.2 | 10:11:03.49 | -04:41:57.2 | OBA | 21.063 | 0.263 | -0.123 |  290 $\pm$ 100 | 20 | 0.8” slit.
s164* | J101106.92-044217.8 | 10:11:06.92 | -04:42:17.8 | OBA | 22.253 | -0.436 | -0.616 |  325 $\pm$ 150 | 19 |
s165* | J101057.66-044130.2 | 10:10:57.66 | -04:41:30.2 | OBA | 21.497 | -0.107 | -0.130 |  325 $\pm$ 100 | 19 |
s166* | J101104.51-044225.9 | 10:11:04.51 | -04:42:25.9 | OBA | 20.934 | -0.201 | -0.925 |  350 $\pm$ 100 | 17 |
Table 2: continued. Table 3: Late type stars observed in this work. Identification codes (ID), spectral types (SpT) and radial velocities ($v_{\mathrm{rad}}$) are included. Stars whose radial velocity was determined manually are marked with (∗). Identifications numbers and coordinates by Massey et al. (2007) are provided when available, otherwhise we give the ones from the 2MASS catalogue (Cutri et al., 2003) and mark the sources with (†). Column Memb. flags the membership to Sextans A. ID | Alt. ID | RA (J2000.0) | DEC (J2000.0) | SpT | $V$ | $B-V$ |  $v_{\mathrm{rad}}$ (km s-1) | Memb. | Notes
---|---|---|---|---|---|---|---|---|---
s167* | J101106.39-044218.1 | 10:11:06.39 | -04:42:18.1 | G5 I | 19.588 | 1.446 |  300 $\pm$ 30 | yes |  BR19-Sextans A 6.
s168* | J101106.00-044212.5 | 10:11:06.00 | -04:42:12.5 | G8 I + B7 I | 19.453 | 1.689 |  350 $\pm$ 30 | yes | 0.8” slit. Blend.
s169* | J101103.95-044228.3 | 10:11:03.95 | -04:42:28.3 | G8 I | 20.031 | 1.768 |  250 $\pm$ 30 | yes |  BR19-Sextans A 4.
s170* | J101102.21-044149.4 | 10:11:02.21 | -04:41:49.4 | K0 V | 16.785 | 0.747 |  40 $\pm$ 30 | no | 0.8” slit.
s171* | J101116.62-044014.4 | 10:11:11.78 | -04:41:18.6 | K2 V | 20.144 | 1.349 |  0 $\pm$ 30 | no |
s172* | J101104.78-044238.8 | 10:11:04.78 | -04:42:38.8 | K5 III | 20.266 | 1.423 |  300 $\pm$ 40 | yes | 1.0” slit.
s173*† | J10110861-0440473 | 10:11:08.60 | -04:40:47.0 | K7 V | - | - |  0 $\pm$ 30 | no | Off-slit.
s174* | J101115.72-044441.0 | 10:11:15.72 | -04:44:41.0 | M3 V | 19.797 | 1.714 |  0 $\pm$ 30 | no | 1.5” slit.
Table 3: continued.

####  4.1.1 Spectral content of the catalogue

Although our observations were designed to find OB stars in Sextans A, we also
observed other types of sources. Here we briefly describe the objects classes
unveiled by our observations.

O stars. There are 71 O-type stars in the sample. Among them, we find the
stars with the earliest spectral types observed at metallicities lower than
1/5 Z⊙. Some of the sources classified as O-types are candidates to undergo
chemically homogeneous evolution or to host a stripped star. This will be
discussed in more detail in Section 5.2.

BA supergiants. Our sample contains 38 B- and A-type supergiants. This
collection will allow us to better constrain the current chemical composition
of Sextans A and to study the [$\alpha$/Fe] abundance ratio. The quantitative
analysis of the spectra of 22 early-B supergiants will yield the abundances of
C, N, O and Si. Strikingly, the silicon lines in the observed spectra appear
weaker than predicted for their corresponding types and the metallicity of the
galaxy. The abundance of magnesium and iron-group elements, can be derived
from analyses of late-B and A supergiants. These analyses will be carried out
as future work.

Red stars. There are 4 evolved stars from Sextans A in our catalogue, two of
them already reported by Britavskiy et al. (2019, B19). These two are
reclassified from K1-3 to G types. This update reinforces the observed trend
of the average spectral types of red supergiants (RSGs) shifting towards
earlier types at lower host galaxy metallicity (Elias et al., 1985; Massey &
Olsen, 2003; Levesque & Massey, 2012). A plausible explanation for this
observed pattern is the dependence of the molecular bands, main indicators of
spectral type in evolved stars, on the metallicity of the environment.
However, this trend may also reflect the shift of the Hayashi limit to higher
effective temperatures as the metallicity decreases (Elias et al., 1985).

Foreground stars. 5 foreground stars landed in our programmed slits. Their
identification as non-members of Sextans A is justified in Section 4.2.

Binary candidates. A total of 33 stars show some signatures of binarity in
their spectra. Some targets present radial velocity variations between
different observations. These are marked as "SB1" in the Notes column of
Tables LABEL:tab:cat_OB_highQ-LABEL:tab:cat_late, along with a question mark
when the variation was smaller than 50 km s-1 or the S/N of the data was poor.
Those sources whose spectra show doubled-peaked lines or exhibit lines of a
second component are marked as "SB2". The ones that met other binary
diagnostics, such as variation in their spectral features, are labelled as
"Bin". All of these cases are discussed individually in Appendix. C. We note
that this constitutes the first sample of candidate binary systems ever
reported in Sextans A.

###  4.2 Membership

To assess the membership of the targets to Sextans A, we inspect whether they
have registered parallaxes or proper motions in the Gaia Early Data Release 3
(Gaia EDR3, Gaia Collaboration et al., 2021). Given the distance to the
galaxy, the majority of the catalogue stars are too faint to have an entry in
Gaia. For those that are included, we expect null proper motions and
parallaxes. Most of our sample stars follow this trend.

We find 12 stars that deviate from the expected parameters. Two of them had
already been identified as foreground based on their radial velocity (Sect.
4.3). The rest of the sources lie within 2$\sigma$ of the mean of the weighted
normal distribution of the Gaia entries for Sextans A, for both parallax and
proper motions. We thus consider that they follow the general trend observed
in the galaxy and discard them as outliers.

We also checked that the absolute magnitudes $M_{V}$ of all catalogue stars
(calculated as described in Section 4.1) are consistent with their spectral
type at the distance of Sextans A. There are only small discrepancies with
respect to the calibrated values, which are discussed in Sect. 5.2.

Finally, we used radial velocity ($v_{\text{rad}}$) as an additional
membership check for all of our sample stars. A group of 5 late type stars
with $v_{\mathrm{rad}}$ $\leq$ 40 km s-1 was identified as foreground. For the
rest of our survey, the majority presents radial velocities consistent with
the systemic velocity and radial velocity curve of Sextans A, as discussed in
Section 4.3, which further supports their membership to the galaxy.

###  4.3 Radial velocity

Figure 1: Comparison of the radial velocities obtained for the sample stars
and the H i rotation curve of Sextans A derived by Skillman et al. (1989). We
excluded from the figure stars identified as foreground. The shaded area
represents the 1$\sigma$, 2$\sigma$ and 3$\sigma$ uncertainty of the measured
rotation curve, being $\sigma$ = 30 km s-1. Stars whose IDs are provided
depart more than 3$\sigma$ from the radial velocity curve, after error bars
have been considered. We also highlight in orange our list of binary
candidates identified from spectral features. In the lower-right corner of the
figure, we draw a representative error bar illustrating the average radial
velocity error of $\pm$ 50 km s-1 for a star observed with resolution R =
1000.

We estimated the radial velocity of each star by measuring the Doppler shift
experienced by the hydrogen and helium lines. Using an automatic procedure, we
determined the centroid of the lines by fitting their profile to a Gaussian
function. To avoid biases introduced by nebular oversubtraction that may
affect the core and introduce profile asymmetries, we imposed a higher weight
on the line wings, defined as the points located between the continuum and a
value set by the program. The height of this threshold was selected according
to the strength of the lines compared to the local continuum. We calculated
the depth of the lines as the difference between the flux of their cores
($f_{\mathrm{min}}$) and the continuum, and positioned the threshold at 15%
and 30% of the estimated depth for strong (1 - $f_{\mathrm{min}}$ > 0.4) and
weak (1 - $f_{\mathrm{min}}$ $\leq$ 0.4) Balmer lines, respectively. These
values stem from the fact that contrasts higher than 1 - $f_{\mathrm{min}}$ >
0.4 are likely caused by nebular oversubtraction. For He lines, where nebular
oversubtraction is less prominent, we set the threshold at 50% of the line
depth for weak (1 - $f_{\mathrm{min}}$ < 0.1) lines and 80% for strong (1 -
$f_{\mathrm{min}}$ $\geq$ 0.1) lines.

After the program fits all lines automatically to Gaussian profiles, it then
rejects results from some lines in two steps according to: (i) their depth
with respect to the noise of the spectrum (S/N values are compiled in Tables
LABEL:tab:cat_OB_highQ-LABEL:tab:cat_late), and (ii) their line width,
rejecting lines when the width exceeded twice the median of all fitted lines
of the spectrum. We also imposed a limit of 1000 km s-1 to the resulting
values of the radial velocity for individual lines and applied a
$\sigma$-clipping algorithm using the median as a reference and rejecting by
2$\sigma$. Finally, we calculated the weighted average of the estimated
$v_{\text{rad}}$ of the selected lines.

The uncertainty was calculated by averaging the error of the Gaussian fit to
the selected lines and then dividing by the square root of the number of used
lines. This result depends not only on the spectral S/N, but also on the
number of lines suitable to measure the radial velocity.

The radial velocity of some stars in the sample could not be determined by our
automatic measurement due to poor spectral S/N. In such cases, we defined the
central wavelength of lines manually and estimated their corresponding Doppler
shift. The assigned uncertainty was also determined manually. These stars are
marked with (∗) in Tables LABEL:tab:cat_OB_highQ-LABEL:tab:cat_late.

In Fig. 1, we compare our results with the rotation curve of Sextans A,
derived by Skillman et al. (1989), to assess the membership of our sample
stars to the galaxy. Overall, our targets are consistent with Sextans A’s
radial velocity. However, there are ten sources (identified in the Figure)
whose radial velocity (and error bars) departs from the H i rotation curve by
more than 3$\sigma$, being $\sigma$ the estimated error of the curve ($\sigma$
= 30 km s-1). The $v_{\mathrm{rad}}$ departures of the ten outliers could be
explained by a possible runaway nature or their belonging to a multiple
system. Yet, spectroscopic follow-up will be needed to confirm any of these
hypotheses. We highlight in orange the stars identified as binary candidates
from their spectra in Figure 1.

###  4.4 The $Q$ pseudo-colour as an indicator of OB type stars

Figure 2: $U-B$ vs $Q$ diagram of Sextans A, showing Massey et al. (2007)’s
catalogue in grey. The main criterion to select candidate O- and early B-type
stars for our observing runs was $Q$ $\leq$ -0.8. The observed stars,
excluding foreground sources, are colour-coded by spectral type and marked
based on their luminosity classes (see legend). Stars without a defined
spectral subtype are drawn with empty symbols. Most of the O-type stars
confirmed by this work are found within the black box, which further supports
the efficiency of the $Q$ pseudo-colour as a selection criterion of massive
stars.

Due to their degenerate colours in the optical range and the effect of
extinction, O- and early B-type stars cannot be identified using photometry
only. However, Garcia et al. (2009) found that the OB stars of the galaxy IC
1613 lied in a particular locus of the $U-B$ vs. $Q$ diagram (see their Fig.
5). This is a consequence of the biunivocal relation between the reddening-
free pseudo-colour $Q=(U-B)-0.72(B-V)$ and spectral type that holds down to A0
types. We have used this as the main criterion to select OB candidates in our
subsequent discovery spectroscopic runs. Garcia & Herrero (2013) used $Q$ <
-0.8 together with UV photometry, achieving a success rate of 70% in the
detection of O-type stars in IC 1613. With the same criteria, Camacho et al.
(2016) had a 50% rate in Sextans A.

The stars in our more extended catalogue are distributed throughout a broad
range of $Q$ values, which enables us to reassess the efficiency of this
candidate selection criterion.

Roughly half of our catalogue, 79 stars, are found in the $Q<-0.8$ box (see
Figure 2). Of these, we find 50 O-type stars, including the earliest of the
sample, 26 B stars, the majority earlier than B1.5, and 3 OB stars. The latter
are not considered in the forthcoming statistics since their poor spectral S/N
prevented us from assigning more definite spectral types. The ratio of
confirmed O-type stars and total observed stars within the $Q<-0.8$ box is
66%. However, if we push this criterion further and reduce the candidate box
to values of $Q<-1.0$, the number increases to 84%, with 21 O stars out of the
25 stars with confirmed spectral type below this limit. Therefore, in the hunt
for the high-mass stars in Sextans A and other galaxies, we advise selecting
candidates with very low $Q$ values.

Outside the $Q<-0.8$ box, we found 17 O stars. The reason for their locus in
the $U-B$ vs. $Q$ diagram is unclear. For instance, in the case of the O star
s127, its colours are unreliable as the star overlaps with a foreground galaxy
(see individual notes in Appendix C). Therefore, if we aim to seek
completeness in the O star population of the galaxy, we must observe beyond
the $Q<-0.8$ box. This highlights the importance of untargeted surveys of
galaxies to fully account for their population of massive stars.

##  5 Discussion

We have obtained spectroscopy of 174 sources, of which 159 are OB stars
(including 41 whose spectral subtype could not be assigned due to poor
spectral S/N), 6 are A- and F-type stars, 4 are red giants/supergiants members
of Sextans A, and 5 are identified as foreground stars. In addition, we
provide the first list of binary candidates with metallicities of 1/10 Z⊙ and
a list of candidates experiencing chemically homogeneous evolution or systems
hosting a stripped star (see Sect. 5.2).

Few massive stars were previously known in very low metallicity environments,
as they are located at far distances (Garcia et al., 2021). The existing
catalogues in IC 1613, WLM and NGC 3109, with 1/7 O⊙ and 1/5 Fe⊙, do not
exceed seventy stars (e.g. Bresolin et al., 2006, 2007; Evans et al., 2007;
Garcia & Herrero, 2013). In our target galaxy Sextans A, with 1/10 Z⊙, only 19
massive stars had been confirmed so far (Camacho et al., 2016; Garcia et al.,
2019). At lower metallicities, only 2 OB stars in SagDIG ($\sim$ 1/20 O⊙,
Garcia, 2018) and one confirmed O-type star in Leo P ($\sim$ 1/30 O⊙, Evans et
al., 2019) are known.

The sample of 159 OB stars assembled in this work constitutes the largest
census of massive stars at sub-SMC metallicities to date. This catalogue also
enables us to study the recent star formation history of Sextans A and
provides a first glimpse into the evolution of 1/10 Z⊙ massive stars. We
discuss the two topics in this section.

###  5.1 The recent star formation history of Sextans A

Figure 3: Sextans A RGB composite image made with $\rm H\alpha$– (red) and
$V$–bands (green) from Massey et al. (2007), and GALEX FUV (blue). The LITTLE
THINGS neutral hydrogen map (Hunter et al., 2012) is overlaid in white.
Regions with confirmed ongoing star formation are marked in the figure. We
show the catalogue stars, excluding identified foreground stars, colour-coded
according to their spectral type and with different symbols based on their
luminosity class. Stars with an undefined spectral subtype are drawn with
empty markers. Figure 4: Same as Fig. 3. Each panel shows the sample OB stars
in different intervals of spectral type. From left to right, we plotted
targets with spectral types earlier than O8, stars with types between O8 and
O9.7 (both included), and sources classified as B-type stars. There is an
apparent stratified pattern of spectral type in region–B, where the early
O-type stars overlap with the high gas density areas, and later types are
located at the inner rims of the H i cloud.

After a quiescent period of 7.5 Gyr, star formation resumed 2.5 Gyr ago in
Sextans A and has been growing continuously, with an increase in the last 0.6
Gyr (Dolphin et al., 2003). The galaxy hosts three major regions with
confirmed ongoing star formation, qualitatively defined by Dohm-Palmer et al.
(2002) and named region–A, –B and –C by Camacho et al. (2016, see also Fig.
3). Region–A (in the north) has been forming stars for the past 400 Myr and it
is about to exhaust its local gas content. The other two, regions–B (in the
southeast) and –C (in the northwest), have been active for 200 Myr and 20 Myr,
respectively, and overlap with the higher concentrations of ionized and
neutral hydrogen that frame the visually bright part of the galaxy (van Dyk et
al., 1998; Dohm-Palmer et al., 2002). Recently, Garcia et al. (2019) found
young massive stars in the southern tip of Sextans A and coined this area as
region–D.

With typical ages under 50 Myr, massive stars signal ongoing star formation.
In Fig. 3, we show the location of the sample stars in Sextans A and use our
159 OB sources as star formation tracers.

Most of the OB stars are located in regions–B and –C, overlapping with the
high column density areas of neutral gas and producing large structures and
shells of ionized hydrogen. Other preferred locations for hot stars are the
northeastern and southern edges of the galaxy, where they ionize small H ii
bubbles and arcs.

The stars holding the earliest types (see Fig. 3) are located independently of
the stellar density and star-forming rate of their host areas. Regions–B, –C,
and –D correspond to regions–1, –3, and –4 from Shi et al. (2014), who
inferred star formation rates using FUV photometry from GALEX of -2.66, -2.32
and -3.19 log $\rm M_{\odot}yr^{-1}kpc^{-2}$, respectively. Finding such early
types in these three areas indicates that star formation has been active until
at least 5 Myr ago (Massey, 2003). No supernova remnant that could help us
further trace star formation is known.

Region–B presents more dispersion in spectral types, and we have found more
early-type stars than in the other star-forming regions of the galaxy. This is
in line with Camacho et al. (2016)’s suggestion that, although activated
longer ago, region–B is still undergoing star formation.

The distribution of massive stars in region–B appears to be layered by
spectral type and correlated with the column density of the neutral gas (see
Fig. 4). The early O-type stars overlap with the highest density of H i (left
panel of Fig. 4), while the later types are located at at the inner edges of
the H i cloud where the column density is lower (middle and right panels of
Fig. 4). The stratification might indicate a temporal sequence of star
formation, which would proceed from the inner rim outwards, following the H i
column-density gradient. Region–C could show a similar layered pattern,
although the catalogue coverage is more sparse in this area due to the way the
observing runs were planned.

Figure 5: Left: LITTLE THINGS neutral hydrogen map (Hunter et al., 2012)
overlaid on a $V$-band image of Sextans A (Massey et al., 2007). We show the
population of OB stars in blue, and identify and highlight in orange the
isolated OB stars. Right: Same RGB image as Fig. 3. We again highlight the
isolated OB stars in orange and show the distance these stars would have
travelled from their nearest OB population if they were runaway stars (white
lines). A projected angle of 1’ corresponding to a distance of 390 pc in
Sextans A is plotted as reference.

A significant fraction of OB stars in our census are found on the outskirts of
Sextans A, at the edges of the neutral gas reservoir and with no associated H
ii structures (see left panel of Fig. 5). Garcia et al. (2019) already pointed
out these unexpected star-forming sites with the detection of star s041 (their
s1) in the southern region of the galaxy. Similar examples exist in other
galaxies near and far. Garcia et al. (2010) spotted OB associations and UV
sources outside the brightest centre of IC 1613. The detection of extended UV-
disc galaxies (Gil de Paz et al., 2005) already suggested the existence of
star formation in low-gas density environments on the galactic outskirts. Only
the spectroscopic confirmation of OB stars at these sites has allowed us to
directly associate some of the UV sources with massive stars. In addition, we
also find OB stars in the central part of the galaxy, where star formation was
considered already finished. This recalls the massive star V39 detected by
Herrero et al. (2010) in an inconspicuous region of IC 1613, and resembles the
observed star formation in the outer fringes of the galaxy. Examples of star
formation at low H i column density also occur in low-mass dwarf galaxies
outside the Local Group (e.g. Teich et al., 2016).

The existence of an underlying population supporting these apparently isolated
stars cannot be discarded since no sufficiently deep IR observations of the
galaxy exist. Yet, Garcia et al. (2019) examined archival WFC3-IR observations
of region–D and found no additional stars around s1. Moreover, Herschel
observations do not support the existence of large masses of dust at the
location of most of these isolated stars.

Could these isolated stars be runaways? Given the distance to Sextans A, we
cannot infer any information from proper motions. A star crossing the entire
galaxy in 10 Myr would exhibit proper motions of $\sim$ 0.01 mas/yr,
significantly lower than the typical uncertainties of Gaia for the proper
motions of stars with $G$ = 20 mag (0.5 mas/yr, Lindegren et al., 2021).

We have manually selected sources that are far from region–A, –B, –C and –D
and highlighted them in Fig. 5. The distances between the isolated sources and
the approximate centre of their nearest star-forming regions (see right panel
of Fig. 5) range between 250 and 825 pc. Assuming lifetimes of 10 Myr and
ejection at early stages of evolution, covering these distances would
correspond to projected velocities of 25-80 km s-1. These values could be
consistent with the high-end tail of the probability distribution of runaway
stars produced by binary interaction found by Renzo et al. (2019). At the
metallicity of Sextans A, their calculations suggest that the velocity of the
ejected stars would be $\sim$ 10 km s-1 higher on average. This is still
significantly lower than the velocity of the alleged runaways. Moreover, the
fraction of expected unbound binaries would be lower at low metallicity. Given
the low expected probability of these cases ($<$ 0.01%, Renzo et al., 2019),
we consider unlikely that these are runaway stars.

We therefore conclude that these sources were likely formed in situ, perhaps
supported by compact, CO-dark gas clouds not signalled by H i maps. An
alternative explanation would be that they are the products of mergers of two
or more lower-mass sources, although this possibility does not solve the
conundrum of the missing underlying population.

###  5.2 Evolutionary status

Figure 6: CMD of Sextans A, showing Massey et al. (2007)’s catalogue in grey
and our sample sources, excluding foreground stars, colour-coded according to
their spectral type. Those without a defined spectral subtype are drawn as
empty symbols. Their luminosity classes are indicated by different symbols
(see legends), and binary candidates are marked with empty brown squares. We
include Kaufer et al. (2004) and Britavskiy et al. (2019) supergiants,
labelled with distinct markers. Error bars for the photometry are shown; we
note that errors for the $V$–band magnitudes are usually smaller than the
symbol size. The computed photometry for $Z$ = 0.10 Z⊙, $v\sin i$ = 100 km s-1
evolutionary tracks by Szécsi et al. (2022) are also provided (see Sect. 5.2).
Dashed lines mark the points where the interpolation of the theoretical
atmosphere models to calculate the photometric magnitudes was incomplete. The
tracks have been shifted to account for Sextans A’s distance modulus and
foreground extinction.

Fig. 6 shows Sextans A’s colour-magnitude diagram (CMD) built with Massey et
al. (2007)’s photometric catalogue, highlighting our sample stars. We also
incorporate blue and red supergiants from previous studies by Kaufer et al.
(2004) and Britavskiy et al. (2019).

We include evolutionary tracks by Szécsi et al. (2022) with Sextans A’s
metallicity ($Z$ = 0.10 Z⊙= 0.00105) and an initial rotational velocity of
$v\sin i$ = 100 km s-1, consistent with the average rotational velocity of OB
stars in other low metallicity regimes (e.g. Sabín-Sanjulián et al., 2017;
Ramírez-Agudelo et al., 2017; Ramachandran et al., 2019). To compare
observations with the evolutionary predictions, we computed the synthetic
photometry of the evolutionary tracks using a grid of tlusty atmosphere models
with an iron abundance of  $[\mathrm{Fe}/\mathrm{H}]=-1.0$ dex. For each point
of the tracks, defined by effective temperature ($T_{\mathrm{eff}}$) and
surface gravity ($\log g$), we assigned an atmosphere model by using a four-
point linear interpolation. Some regions of the evolutionary paths lay outside
the tlusty grid coverage, and only 2 or 3 points could be used to interpolate,
leading to perhaps non-realistic atmospheres. These regions are shown as
dashed lines in Fig. 6. Finally, the tracks are shifted to account for the
distance modulus ($\mu_{0}$ = 25.63, Tammann et al., 2011) and foreground
extinction ($\mathrm{E}(B-V)_{\mathrm{fg}}$= 0.044, Skillman et al., 1989) of
Sextans A. No correction is applied to the photometry of the catalogue stars.

The OB stars of our catalogue are spread across the blue plume. The observed
dispersion in both spectral types and luminosity classes impact the masses and
ages that can be derived from the evolutionary tracks. Although most of the
earliest stars concentrate in the bluest regions of the plume, others overlap
with stars of later spectral types. The spread can be partly explained by the
intrinsic errors of the assigned spectral type, but it is mostly caused by
internal interstellar extinction. The same effect has been observed in other
LG dIrr galaxies like IC 1613 (Garcia et al., 2009) and SagDIG (Garcia, 2018).
A thorough characterisation of the internal extinction of Sextans A will be
published in a subsequent paper of this series.

The late-type stars of Sextans A are mainly found at the red ends of the
evolutionary tracks. They are concentrated in two different groups, a less
massive one between 12 and 19 M⊙ and a second population with higher initial
evolutionary masses ranging between 30 and 40 M⊙. Recent works (Beasor et al.,
2019; Britavskiy et al., 2019) speculate that the more luminous and massive
group of RSGs are the evolved analogues of the massive blue straggler stars
and might be products of mergers or binary mass transfer.

Figure 7: Same as Fig. 6, zooming into the main sequence and bluer colours. We
shade in blue the area of bluer colours than the evolutionary tracks, with a
lower boundary in $B-V$ defined by the intrinsic colour of a blackbody at 200
kK. This region contains the list of blue outliers listed in Table 4. In the
lower-left corner, we provide the reddening vector, adopting $R_{V}$ = 3.1.

We detect that 18 OB-type stars exhibit bluer colours than the evolutionary
tracks (shaded region in Fig. 7). This number may increase after internal
reddening is corrected. Slowly rotating single-star evolutionary tracks do not
predict OB stars in this locus, as shown in Fig. 6. We imposed a lower limit
on $B-V$ of -0.36 mag corresponding to the colour of a blackbody with
$T_{\mathrm{eff}}$ = 200 kK (the highest effective temperature observed in
Wolf-Rayet stars, Sander et al., 2012; Tramper et al., 2015). Photometric
measurements bluer than this value are discarded as spurious. They might be
due to strong emission lines contributing to the total spectral energy
distribution and thus the measured colours.

A star might reach extreme hot temperatures, hence blue colours, if it evolves
chemically homogeneously. Rapid rotation may lead a massive star to efficient
mixing, dragging helium into its envelope as it is produced, and thus never
establishing a chemical gradient (Maeder & Meynet, 2000). With hydrogen still
burning in its core, the star will evolve to colours and effective
temperatures hotter than the ZAMS, entering a Wolf-Rayet-like stage and
remaining extremely hot throughout its lifetime (Yoon & Langer, 2005). This
evolutionary path is more expected among very low-Z massive stars (Brott et
al., 2011), since the reduced mass- and angular momentum-loss via winds may
allow for rapid rotation rates. However, there is still no observational
confirmation for the existence of CHE stars.

The photometry and the assigned spectral types of some of the stars in the
blue locus are consistent with them following chemically homogeneous
evolution. However, this needs to be confirmed with measurements of high
helium surface abundances and large rotational velocities ($v\sin i$ > 250 km
s-1). Unfortunately, the S/N and the resolution ($\sim$ 300 km s-1) of the
current data set do not allow such analysis.

Alternatively, these sources could be post-binary interaction products, in
particular, binary systems whose primary star has been removed of its
hydrogen-rich envelope through the interaction with its companion. This loss
produces very compact, hot helium stars (e.g. Kippenhahn, 1969; Podsiadlowski
et al., 1992), known as stripped stars. Theoretical atmosphere models predict
that the spectra of these objects at their high-mass end show He and N lines
(Götberg et al., 2017), similar to the observed in nitrogen-rich Wolf-Rayet
stars. Such spectral features have not been detected within our sample.
However, at $Z$ = 1/5 Z⊙ and intermediate or low mass values, their spectra
present typical features of early O-type stars with nitrogen enrichment
(Götberg et al., 2018). Some of our blue outliers exhibit this spectral
morphology. On the other hand, if the companion of the stripped star were to
dominate the optical part of the spectrum, we would observe a typical OB
spectrum, possible signals of multiplicity and UV and blue photometric excess.
The star would also experience higher rotational velocities after the mass and
angular momentum transfer. We detect 4 binary candidates within the 18 sources
found at bluer colours than the evolutionary tracks.

We have assessed the different features that both CHE stars and systems
hosting a stripped star would exhibit. According to the theoretical
predictions, CHE stars are hot and underluminous for their spectral type
(Kubátová et al., 2019). They also hold high projected rotational velocities
and helium abundances. On the other hand, systems hosting a stripped star
would show signs of multiplicity and UV excess. If the stripped star dominates
the optical spectrum, it would present a lower luminosity than a star with the
same spectral classification (Götberg et al., 2018).

Figure 8: UV CMD of Sextans A, displaying Bianchi et al. (2012)’s catalogue in
grey and our sample OB stars colour-coded according to their spectral type.
Binary candidates are marked with brown squares and the blue outliers with
black diamonds.

With this in mind, we have compiled a selection of properties for the list of
blue outliers in Table 4. We have analysed the location of these stars in the
ultraviolet CMD using photometry by Bianchi et al. (2012) (see Fig. 8), noting
whether or not they are also blue outliers in this diagram, which would flag
the existence of a UV excess. In addition, we have derived the luminosity of
these objects using Sextans A’s distance modulus $\mu_{0}$ = 25.63 (Tammann et
al., 2011), $(B-V)_{0}$ colours estimated from the pseudo-colour $Q$ (Massey
et al., 2000) and calibrations to estimate effective temperatures and
bolometric corrections (Balona, 1994; Martins et al., 2005; Markova & Puls,
2008). A luminosity lower than the calibrated values by Martins et al. (2005);
Ramachandran et al. (2019) for their spectral type would qualify the objects
as candidates for CHE or stripped stars. Nine blue outliers are underluminous
compared to the calibrations by a difference of $\geq$ 0.2 dex.

Finally, we have also studied whether the spectra are consistent with high
projected rotational velocities to support their designation as CHE candidates
or companions to stripped stars. We have used FASTWIND models compatible with
the assigned spectral classification of the candidates (Lorenzo et al. in
prep.) and helium abundance of $\epsilon=0.09$, and convolved to the
resolution of the respective observations. Since the S/N and $R$ $\sim$ 1000
of our observations prevent us from discerning $v\sin i$ in the range of
75-300 km s-1, we have checked whether the line profiles of the observed
spectra were broader than the synthetic lines with $v\sin i$ = 300 km s-1 (see
Fig. 9 for an example of the followed procedure). Results of the evaluation
are collected in Table 4. Six objects in our list of blue outliers present
line profiles broader than fast-rotating models. We caution, however, that the
broadening may also be due to blends or unresolved binary systems.

We venture a classification for the blue outliers in Table 4, based on their
spectral type, luminosity, UV excess, tentative $v\sin i$ and possible
companion. Nonetheless, we remark that follow-up observations with higher
resolution and S/N are strongly needed to measure rotational velocities and He
abundances that may support any of these scenarios.

Figure 9: Comparison of the spectrum of the blue outlier s023 (O8 Vz), in
black and smoothed for clarity, with FASTWIND models compatible with O8-type
dwarfs (Lorenzo in prep.), convolved to the resolution of the spectrum (R
$\sim$ 1300) and with the projected rotational velocity $v\sin i$ = 300 km
s-1(in orange).  Table 4: List of blue outliers in the CMD of Sextans A.
Identification codes (ID), spectral types (SpT) and optical photometry from
Massey et al. (2007) are provided. We also include the UV photometry from
Bianchi et al. (2012) when available and flag the objects that are outliers in
the UV CMD. The derived luminosities (calculated as indicated in the text) and
the calibrated values for the corresponding spectral type are included. We
also mark whether or not the source presents signatures of multiplicity in its
spectra and if its line profiles are broader than models with $v\sin i$ > 300
km s-1. Armed with this information, we tentatively assign the objects into a
category in column Suggested type.

|  | Optical phot. | UV phot. |  |  |  |  |
---|---|---|---|---|---|---|---|---
ID | SpT | $V$ | $B-V$ |  | UV
---
detection
F170W | F170W - F225W |  | UV
---
outlier
| log (L/L⊙)
---
This work
| log (L/L⊙)
---
Calibrated
| Binary
---
cand.
| Broad
---
lines
| Suggested
---
type
s002 | O4 Vz | 20.369 | -0.279 | no info | - | - | - | 5.60 | 5.67 | no | no | ?
s004 | O5 III | 20.917 | -0.277 | no info | - | - | - | 5.34 | 5.73 | yes | no | Stripped star
s005 | O5 V((fc)) | 21.751 | -0.364 | yes | 18.47 | -0.64 | yes | 4.83 | 5.49 | no | no | Stripped star
s007 | O6 V | 20.515 | -0.280 | yes | 17.45 | -0.43 | yes | 5.44 | 5.32 | yes | yes | CHE or OB + stripped star
s009 | O7 V | 20.659 | -0.285 | yes | 17.82 | -0.96 | yes | 5.30 | 5.14 | no | no | OB + stripped star
s013 | O7.5 III | 21.532 | -0.295 | no info | - | - | - | 4.87 | 5.42 | no | yes | CHE
s015 | O7.5 V + B? | 22.259 | -0.410 | yes | 18.64 | -0.18 | no | 4.50 | 5.05 | yes | no | Stripped star
s018 | O8 III | 21.092 | -0.321 | no info | - | - | - | 4.97 | 5.35 | no | yes | CHE
s023 | O8 V | 20.288 | -0.299 | no info | - | - | - | 5.36 | 4.96 | no | yes | OB + stripped star
s030 | O8.5 III | 21.008 | -0.322 | no info | - | - | - | 4.94 | 5.28 | no | yes | CHE
s032 | O8.5 V | 21.216 | -0.326 | no info | - | - | - | 4.86 | 4.86 | no | yes | CHE
s033 | O8.5 V | 21.654 | -0.271 | yes | 18.92 | -0.41 | no | 4.83 | 4.86 | no | no | ?
s040 | O9 III + mid-B? | 21.285 | -0.311 | yes | 18.22 | -0.35 | no | 4.83 | 5.21 | yes | ? | ?
s063 | B0 III | 20.727 | -0.295 | yes | - | - | - | 4.95 | 4.98 | no | no | ?
s080 | B1 III | 21.471 | -0.292 | yes | 18.38 | -0.69 | yes | 4.50 | 4.72 | no | no | ?
s101 | B3 III | 21.254 | -0.328 | yes | 18.28 | -0.69 | yes | 4.18 | 3.90 | no | no | ?
s129 | O V + neb | 21.237 | -0.310 | yes | 17.81 | -0.75 | yes | 3.85 | 4.68 | no | ? | Stripped star
s164 | OBA | 22.253 | -0.436 | yes | 18.99 | -0.24 | no | 3.13 | - | no | ? | ?

##  6 Summary and conclusions

Very low metallicity massive stars hold the key to interpreting processes from
the earliest times of the Universe, such as the re-ionization epoch and early
cosmic chemical enrichment. Their low metal content may lead to significant
changes in their evolution, some of them with an extreme impact on their
ionizing fluxes like the chemically homogeneous evolution channel. In contrast
with the extensive studies in the Magellanic Clouds (e.g. Ramírez-Agudelo et
al., 2017; Sabín-Sanjulián et al., 2017; Ramachandran et al., 2019), current
references for the metal-poor Universe, massive stars at sub-SMC metallicities
were very poorly studied. Building an extensive census was a fundamental first
step to understanding the physics and evolution of these objects, and to
ultimately extrapolating the properties of the first, metal-free stars.

In this work, we present the results of a spectroscopic survey of massive
stars in the 1/10 Z⊙ galaxy Sextans A built with observations from the 10.4-m
telescope GTC. Our 159 OB star catalogue is the largest census of massive
stars at metallicities lower than the SMC to date. We also provide the first
list of candidate stars undergoing chemically homogeneous evolution or systems
hosting a stripped star. Confirming even one single case of these scenarios
will have critical implications for the stellar evolution theories and the
interpretation of high-redshift galaxies. In addition, we identify 33 sources
within the sample displaying signs of multiplicity in their spectra.

Most of the sample OB stars are located in rich H i density regions producing
large or small bubbles of ionized gas. However, we also find massive stars in
inconspicuous areas of the galaxy where the column density of the neutral gas
is low. This was already pointed out by Garcia et al. (2019) in the south of
Sextans A, but we also detect them in the gas-void galactic centre and other
external areas of the galaxy. Our extensive spectroscopic sample and future
detailed maps of gas will enable disentangling the roles of neutral and
molecular gas in star formation at very low metallicity and identify the
mechanisms that have triggered star formation in unexpected regions.

We find high scatter in the colour-magnitude diagram, hinting that the
internal extinction of Sextans A is significant and non-uniform. Mapping the
local reddening and estimating the stellar parameters and abundances of the
massive stars of our catalogue by quantitative spectroscopic analyses is left
for future work.

We have also assembled a sample of 38 BA supergiants that, with further
quantitative analyses, will yield the abundances of $\alpha$\- and Fe-group
elements. Studying the chemical composition of blue stars will help us
investigate the possible existence of chemical abundance variations across the
galaxy and will shed light on the chemodynamical evolution of Sextans A.

In addition, we observed 4 late-type stars in the galaxy. Enhancing the
statistics of evolved massive stars in LG galaxies allows us to quantify the
role of metallicity in evolution.

This work presents the first extensive sample of sub-SMC metallicity massive
stars. This catalogue is a fundamental first step for high-quality, high-
resolution spectroscopic observations that will enable a thorough
characterisation of stellar and wind properties of very metal-poor massive
stars. These, in turn, will constrain stellar evolutionary models leading to
more realistic predictions of stellar feedback for these stars, which is
fundamental to interpreting observations of high-redshift galaxies.

## Acknowledgements

M. Lorenzo, M. Garcia and F. Najarro gratefully acknowledge support by grants
PID2019-105552RB-C41 and MDM-2017-0737 Unidad de Excelencia "María de
Maeztu"-Centro de Astrobiología (CSIC-INTA), funded by
MCIN/AEI/10.13039/501100011033. M. Lorenzo ackowledges funding from grant
PRE2019-087988 under project MDM-2017-0737-19-3, and by “ESF Investing in your
future". A. Herrero acknowledges support by the Spanish MCI through grant
PGC-2018-0913741-B-C22 and the Severo Ochoa Program through CEX2019-000920-S.
M. Cerviño acknowledges financial support from the Spanish MCI through grant
PID2019-107408GB-C41. N. Castro gratefully acknowledges funding from the
Deutsche Forschungsgemeinschaft (DFG) – CA 2551/1-1.

The work has made use of the gtcmos pipeline for the reduction of the
GTC/OSIRIS spectroscopic data for which we thank its author Divakara Mayya. It
also has made use of "Aladin sky atlas" developed at CDS, Strasbourg
Observatory, France (Bonnarel et al., 2000; Boch & Fernique, 2014); SIMBAD
database, operated at CDS, Strasbourg, France (Wenger et al., 2000); and
NASA’s Astrophysics Data System Bibliographic Services.

Finally, we want to thank our colleagues at the IAU Symposium 361: Massive
Stars Near and Far for their insightful questions and comments.

## Data Availability

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

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## Appendix

##  Appendix A Observation logs

This section collects the logs of our four observation campaigns, including
the date and weather conditions and total exposure times (Tables 5-8). In
addition, we provide a list of the observing blocks and super-observing blocks
coadded to generate the final spectrum of the stars observed in the MOS1 run
(Table 9).

Table 5: MOS1 observing log. | Observing
---
block
Date | Seeing (”) | Transparency | Moon | Airmass | Exposure time (s)
OB01 (pre-image) | 2 Apr 2014 | 0.8 - 1.0 | Clear | Grey | 1.29 | -
OB02 | 16 Jan 2015 | 1.2 | Clear | Dark | 1.20 | 2950
OB03 | 3 May 2014 | 1.1 | Photometric | Dark | 1.24 | 2950
OB04 | 3 May 2014 | 1.2 | Photometric | Dark | 1.37 | 2950
OB05 | 4 May 2014 | 0.9 | Clear | Grey | 1.57 | 2950
OB06 | 5 May 2014 | 1.0 | Clear | Grey | 1.27 | 2950
OB07 | 5 May 2014 | 1.0 | Clear | Grey | 1.42 | 2950
OB08 | 5 May 2014 | 0.9 | Clouds | Grey | 1.72 | 2950
OB09 | 6 May 2014 | 1.0 | Clear | Grey | 1.29 | 2950
OB10 | 16 Jan 2015 | 1.0 | Clear | Dark | 1.26 | 2950
OB11 | 16 Jan 2015 | 0.9 | Clear | Dark | 1.39 | 2950
OB12 | 18 Jan 2015 | 1.1 | Clear | Dark | 1.20 | 2950
OB13 | 18 Jan 2015 | 0.8 | Clear | Dark | 1.25 | 2950
OB14 | 18 Jan 2015 | 0.8 | Clear | Dark | 1.39 | 2950
Table 6: MOS2 observing log. | Observing
---
block
Date | Seeing (”) | Transparency | Moon | Airmass | Exposure time (s)
Night-1 | 16 Feb 2018 | 0.9 | Clear | Dark | 1.20 | 12600
Night-2 | 17 Feb 2018 | 1.2 | Clear | Dark | 1.20 | 9000
Table 7: LSS2 observing log. | Observing
---
block
Date | Seeing (”) | Slit-width (”) | Transparency | Moon | Airmass | Exposure time (s)
OB01 | 18 Feb 2020 | 1.0 | 1.2 | Clear | Dark | 1.23 | 2806
OB02 | 18 Feb 2020 | 1.0 | 1.2 | Clear | Dark | 1.20 | 2806
OB03 | 18 Feb 2020 | 1.3 | 1.2 | Clear | Dark | 1.28 | 2806
OB04 | 18 Mar 2020 | 0.9 | 1.2 | Phot. | Dark | 1.68 | 2806
Table 8: LSS3 observing log.

| Observing
---
block
Date | Seeing (”) | Slit-width (”) | Transparency | Moon | Airmass | Exposure time (s)
OB01 | 16 Jan 2021 | 1.0 - 1.2 | 1.2 | Photometric | Dark | 1.25 | 3600
OB06 | 16 Jan 2021 | 1.0 - 1.2 | 1.2 | Photometric | Dark | 1.20 | 3600
OB04 | 16 Jan 2021 | 1.3 | 1.2 | Photometric | Dark | 1.25 | 3600
OB05 | 16 Jan 2021 | 1.4 | 1.5 | Photometric | Dark | 1.41 | 3600
OB03 | 17 Jan 2021 | 0.8 | 0.8 | Photometric | Dark | 1.30 | 3788
OB02 | 17 Jan 2021 | 0.75 | 0.8 | Photometric | Dark | 1.20 | 3788
OB13 | 17 Jan 2021 | 0.8 | 1.0 | Photometric | Dark | 1.25 | 6846
OB17 | 17 Jan 2021 | 0.8 | 0.8 | Photometric | Dark | 1.75 | 5055
OB18 | 18 Jan 2021 | 0.8 | 0.8 | Photometric | Dark | 1.48 | 3804
OB15 | 18 Jan 2021 | 0.8 | 0.8 | Photometric | Dark | 1.21 | 6846
OB11 | 18 Jan 2021 | 0.8-1.0 | 1.0 | Photometric | Dark | 1.29 | 6846
OB44 | 19 Jan 2021 | 1.2-1.5 | 1.5 | Photometric | Dark | 1.25 | 4388
OB31 | 19 Jan 2021 | 1.5 | 1.2 | Photometric | Dark | 1.20 | 6225
OB45∗ | 9 Mar 2021 | 1.3-1.4 | 1.2 | Photometric | Dark | 1.22 | 7125
OB46∗ | 10 Mar 2021 | 1.3-1.5 | 1.2 | Photometric | Dark | 1.50 | 7125
OB47∗ | 9 Mar 2021 | 1.3-1.5 | 1.2 | Photometric | Dark | 1.30 | 2883
OB49∗ | 10 Mar 2021 | 1.3-1.5 | 1.2 | Photometric | Dark | 1.21 | 7125
OB50∗ | 10 Mar 2021 | 1.3-1.5 | 1.2 | Photometric | Dark | 1.35 | 3704

  * •

Notes. (∗) observing blocks performed in service mode.

Table 9: List of the observing blocks and super-observing blocks coadded to generate the final spectrum of each star observed in the MOS1 run. ID | Exposure time (s) | Exposures
---|---|---
s003 | 26550 | OB02, sOB03(=OB03+OB04), OB05, OB06, OB10, sOB13(= OB12+OB13+OB14)
s004 | 26550 | OB02, OB03, OB04, OB05, OB06, OB07, OB10, OB12, OB13
s028 | 26550 | OB02, OB03, OB05, OB06, OB07, OB10, OB12, OB13, OB14
s029 | 26550 | OB02, OB03, OB04, OB05, OB06, OB07, OB10, OB12, OB13
s041 | 26550 | OB02, OB03, OB04, OB05, OB06, OB07, OB10, OB12, OB13
s042 | 23600 | OB02, OB03, OB04, OB06, OB10, OB12, OB13, OB14
s061 | 32450 | OB02, OB03, OB04, OB05, OB06, OB07, OB10, OB11, OB12, OB13, OB14
s063 | 20650 | OB02, OB03, OB04, OB10, OB11, OB12, OB13
s071 | 29500 | OB02, OB03, OB04, OB05, OB06, OB07, OB10, OB12, OB13, OB14
s089 | 29500 | OB02, OB03, OB04, OB06, OB07, OB10, OB11, OB12, OB13, OB14
s092 | 32450 | OB02, OB03, OB04, OB05, OB06, OB07, OB09, OB10, OB12, OB13, OB14
s099 | 26650 | OB02, sOB03(= OB03+OB04), OB05, OB06,OB10, sOB13(= OB12+OB13+OB14)

##  Appendix B Spectra

In this section, we provide the GTC-OSIRIS spectra of all our targets. We
first show the spectra for the blue massive stars with a full classification
(corresponding to Table LABEL:tab:cat_OB_highQ), then for the blue massive
stars with no defined spectral subtype (Table LABEL:tab:cat_OB_lowQ) and
finally for the late type stars (Table LABEL:tab:cat_late).

Figure 10: GTC–OSIRIS spectra of the blue massive stars of our catalogue with
complete spectral classification. The data have been corrected by heliocentric
and radial velocity, and smoothed for clarity based on the S/N of the
spectrum. Figure 11: Same as Figure 10, continued. Figure 12: Same as Figure
10, continued. Figure 13: Same as Figure 10, continued. Figure 14: Same as
Figure 10, continued. Figure 15: Same as Figure 10, continued. Figure 16: Same
as Figure 10, continued. Figure 17: Same as Figure 10, continued. Figure 18:
Same as Figure 10, continued. Figure 19: Same as Figure 10, continued. Figure
20: Same as Figure 10, continued. Figure 21: Same as Figure 10, continued.
Figure 22: Same as Figure 10, continued. Figure 23: Same as Figure 10,
continued. Figure 24: Same as Figure 10, continued. Figure 25: Same as Figure
10, continued. Figure 26: Same as Figure 10, continued. Figure 27: Same as
Figure 10, continued. Figure 28: GTC–OSIRIS spectra of the sampled blue
massive stars with no defined spectral subtype. The data have been corrected
by heliocentric and radial velocity, and smoothed for clarity based on the S/N
of the spectrum. Figure 29: Same as Figure 28, continued. Figure 30: Same as
Figure 28, continued. Figure 31: Same as Figure 28, continued. Figure 32: Same
as Figure 28, continued. Figure 33: Same as Figure 28, continued. Figure 34:
GTC–OSIRIS spectra of the late type stars in the sample. The data have been
smoothed for clarity, and corrected by heliocentric and radial velocity.
Figure 35: Same as Figure 34, continued.

##  Appendix C Comments on targets

In this section, we provide information on individual stars regarding the
observation and reduction procedure when it is needed to understand our
spectral classification and the morphology of the given spectra. We also
detailed the process of spectral classification when it is problematic or when
the source presents interesting features.

###  C.1 Stars with a full classification

s001 (O3.5 V): The spectrum is dominated by nebular emission lines, although
it displays stellar wings for the Balmer series. He ii 4541 and He ii 4686 are
present, but we do not detect any He i lines. The star is classified as O3.5 V
based on the tentative detection of N v 4604 in absorption. Class V is
assigned based on the width of the Balmer lines. We also note that He ii lines
are weaker than expected, which may be due to the contamination from a nearby
star that we detected at the reduction stage. The radial velocity was measured
using He ii lines, and we note the resulting value is not consistent with the
core of the Balmer lines, hence we mark it as a possible binary.

s002 (O4 Vz): The spectrum displays high nebular contamination, He i 4471, He
i 4920 and He i 5015 lines in emission and strong emissions at [OIII] 4363 Å,
[OIII] 4959 Å and [OIII] 5007 Å, fitting the typical locus of the sky lines.
It also presents strong absorption of He ii lines. Since He i 4471 is hidden
by the nebular contamination, we classify this star by comparing the depths of
He ii 4200, He i 4144 and He i 4121. As He ii 4686 in absorption is stronger
than He ii 4541, we assigned Vz as its luminosity class.

s003 (O3-O5 Vz): This is star s2 in Garcia et al. (2019), who noted the broad
He ii lines and the absence of He i lines. The new reduction confirms the
initial O3-O5 Vz spectral classification, further supported by the tentative
detection of N iii 4634, 4640–4642 in emission. The updated spectrum also
shows $\rm H\beta$ and $\rm H\gamma$ in emission. However, a nebular or
stellar origin cannot be discerned at the current resolution.

s004 (O5 III): This star was observed in two observing runs, MOS1 and MOS2,
where no radial velocity variations were detected. The poorer S/N of the MOS2
observation prevented the co-addition of the two spectra. We reclassify the
star with respect to Garcia et al. (2019) since the new reduction reveals
stronger He ii 4541. The luminosity class is reassigned to III based on the
width of the Balmer lines, although He ii 4686 is in strong absorption. We
note the new classification makes the star even more underluminous for its
spectral type than reported in Garcia et al. (2019).
We observe a $v_{\mathrm{rad}}$ mismatch between He i 4471 and He ii 4541 in
different observing blocks, and He i 4471 is not detected in OB03-OB06, so we
mark the star as a possible binary.

s005 (O5 V((fc))): The spectrum presents an unidentified emission at 4375 Å, a
very broad He ii 4200 line and high nebular contamination shifted to the blue
in the Balmer lines with respect to their cores. By comparing the strong
absorption of He ii 4541 with the He i 4471 line, we classify this star as O5.
Nevertheless, this star might carry a later type since He i 4471 may be
affected by nebular contamination. The very broad Balmer lines and the
relative strength between He ii 4686 and He ii 4541 indicate a luminosity
class of V. This star presents a counterpart in the UV which additionally
supports its O type. Lastly, due to the tentative presence of N iii 4634 and C
iii 4647 in emission, we add the qualifier ((fc)).

s006 (O6 I): The star is classified as O6 due to the strong He ii 4200 and He
ii 5411 lines. The spectrum is severely affected by nebular oversubtraction,
which leads to a different radial velocity ($v_{\mathrm{rad}}$ = 300 km s-1)
for the Balmer lines and He i 4471 absorption. Since stellar Balmer wings are
dominated by nebular oversubtraction and He ii 4686 is not in absorption,
luminosity class I is assigned.

s007 (O6 V): This source is the central star of one of the southeast bubbles
of region-B. It was observed in a very short slit of MOS2, which made
unfeasible the selection of sky boxes for apall background subtraction. The
spectrum of one of the sky-slits was used instead, which results in a still
incomplete nebular subtraction (see Balmer lines and He i 4471) and the loss
of the bluest spectral range ($\lambda$ $\lesssim$ 4050 Å).
We detected a $v_{\mathrm{rad}}$ mismatch of $\sim$ 70 km s-1 between
different OBs, corrected prior to co-adding the spectra. The main He ii lines
and some He i lines seem double in the raw spectra of Night-1, suggesting it
could be a binary- However, we cannot discard that this may be due to central
nebular contamination by incomplete nebular subtraction at the current
resolution.
This star was also observed in run LSS2, which was used to recover the broader
$\rm H\delta$ profile, but it does not shed further information on the binary
hypothesis.

s008 (O6.5 V): This source is a blend of two unresolved stars identified in
Bianchi et al. (2012). The spectrum displays strong nebular contamination
affecting the Balmer series and He i 4471. The He ii lines are in strong
absorption and their relative strengths with the He i lines (He i 4387 and He
i 4144) suggest an O6.5 type. Since the Balmer series are broad, we assigned a
dwarf luminosity class. The spectrum may be contaminated by nearby stars.

s009 (O7 V): This star was observed in two observing runs, MOS2 and LSS3. We
did not detect any radial velocity variations, and we discarded the co-
addition of the spectra since there was no S/N improvement. The spectrum is
plotted with smooth equals to 3 for clarity purposes. Without it, the
component of nebular oversubtraction can be cleaned from the He i 4471 core.
This line is similar in strength to He ii 4541, rendering type O7. Likewise,
the feature around Mg ii 4481 is revealed as an extremely wide absorption,
unlikely to originate in a stellar spectrum, so it is not considered for
classification. Luminosity class V is assigned based on the strong absorption
of He ii 4686, although the width of the Balmer lines suggests class III.

s010 (O7 V): The spectrum is very noisy and shows normalization issues.
Spectral type of O7 is assigned based on the relative strength of He i 4471
and He ii 4541. The adopted radial velocity is inconsistent with the Balmer
series. He ii 4686 in absorption suggest class V.

s011 (O7.5 III + B0 I): The spectrum shows strong He ii lines, but also clear
Si iv 4089 and perhaps Si iii 4552 and Mg ii 4481 lines. The target was
observed in two different nights, but we have not detected changes in the
spectral morphology or a radial velocity variation, hence we conclude that it
is likely a blend of two stars. We assigned an O7.5 type to one of the stars
based on the strength of its He ii lines compared with FASTWIND models
(Lorenzo et al. in prep.) and considering that the strong He i lines have
contamination from the other star. Since He ii 4686 is present in absorption
and its strength is similar to He ii 4541, we assigned a giant luminosity
class. Spectral type B0 I is assigned to the second star due to the presence
of Si iv 4089 and its relative strength with respect to He i 4121.

s012 (O7.5 III): This star was observed in a short slit with little room for
the selection of sky boxes for sky subtraction. This problem is further
exacerbated by the black column artefact from the adjacent slit. As a
consequence, sky subtraction is very poor. We found that the best solution was
to use only the spectrum observed during the second night. He i lines are
wide, while He ii lines are narrow, and their $v_{\mathrm{rad}}$ is slightly
different. While it cannot be discarded that this is due to nebular
oversubtraction, we mark the star as a possible binary. The spectral type was
assigned on the premise that the He ii lines are stellar.

s014 (O7.5 III((f))): This star shows indications of increased He abundance
and strong winds (Camacho et al., 2016).

s015 (O7.5 V + B?): The spectrum displays strong He ii lines consistent with a
radial velocity of $v_{\mathrm{rad}}$ $\sim$ 225 km s-1 and an O7.5 type.
However, it also shows strong Si iii 4552. In addition, the He i lines suggest
a different radial velocity. Therefore, we mark the star as a possible SB2
with a possible B star companion.

s018 (O8 III): The core of $\rm H\beta$ and $\rm H\gamma$ is dominated by
oversubtracted nebular lines, which also affect the He i 4471 line.

s019 (O8 III): This source is a blend of three stars included in the Bianchi
et al. (2012)’s catalogue. We assign the photometry and ID of the brightest
one. We clearly detect the Balmer series, some He i lines and strong He ii
lines. However, we note He i 4471 is affected by nebular contamination. We
classify the star as an O8 type based on the ratio of He ii 4541 and He i
4387, the strength of He ii 4200 and the presence of He ii 5411. We assign
luminosity class III because of the similar strength of He ii 4541 and He ii
4686.

s020 (O8 V): Strong nebular contamination affecting the Balmer series, He i
4471 and some other He i lines. The spectrum displays He ii lines. The
relative strength of He ii 4541 with He i 4388 and He ii 4200 with He i 4144
points to an O8 type. Since He ii 4686 is present in absorption and the Balmer
lines are wide, we assign a dwarf luminosity class.
Two spectra were taken for this star in two consecutive nights. The spectrum
of the second night seems to lack the absorption of He ii 4200 and He ii 4541
(see Figure 36). Since the star is located in a crowded region we simply
marked it as an SB2 candidate or blend. However, this could be a short period
eclipsing binary. We do not detect radial velocity variations between
observations.

Figure 36: Observations of star s020 on two consecutive nights. We detect a
change in the spectra features of the source. On the first night (in orange),
He ii 4200 is detected in strong absorption, while on the second night (in
black) it appears extinct, suggesting a possible binary system.

s021 (O8 V): This star was observed in a narrow 0.63” slit tilted from the N-S
orientation by 3° in the MOS2 run. The slit was too short, compromising the
selection of sky subtraction with apall. We used the sky spectrum of one of
the sky-slits as it yielded better results. The star was also observed in the
observing run LSS3, no radial velocity variations were detected between the
two runs. We discarded the combination of the spectra since no improvement in
the S/N on the resulting spectrum was detected.
This star is a UV bright source in Bianchi et al. (2012)’s catalogue with a
nearby star undetected in the UV. Hubble Space Telescope resolves them fully
at F555W, but Massey’s ground-based imaging could not separate them and
provides one single detection for both. We included only the blue object in
the slitlet, but some contamination is to be expected. In fact, the whole
blend was included in the LSS3 run, and no significant difference is found
between the spectrum of the MOS2 run and the blend, supporting the hypothesis
that the source itself is also blended. There seems to be a mismatch between
the radial velocities of He i and He ii lines of $\sim$ 80 km s-1, which could
be caused by the blend or by nebular contamination at the core of the stellar
lines.

s022 (O8 V): This source was classified as O9.5 III-V by Camacho et al.
(2016), but we reclassify it as an O8 V based on the strength of He ii 4200
and He ii 5411, consistently with the classification criteria of this paper.
We note that the observed magnitude would instead suggest class I. HST
photometry reveals this source is a blend of two stars identified in Bianchi
et al. (2012).

s023 (O8 V): Strong He ii lines are detected. Since He i 4471 is in emission
due to high nebular contamination, we based our classification on the strength
of He ii 4200, the presence of He ii 5411 and the ratio of He ii 4541 with He
i 4388. The strong detection of He ii 4686 and the lack of Si iv suggests a
dwarf luminosity class. The spectrum could be contaminated by the nearby star
s001.

s024 (O8 V): Strong nebular contamination. The spectrum displays strong lines
of He ii. The relative strength of He ii 4541 and He ii 4200 with respect to
He i 4471 and He i 4388 suggests an O8 type. We assign V as its luminosity
class due to the broadening of the Balmer series and the presence of He ii
4686 in absorption.

s025 (O8 V): We detect He ii lines. The relative strength of He ii 4541 with
He i 4471 and He i 4388, and the detection of He ii 5411 points to an O8 type.
Since He ii 4686 is in strong absorption, we classify this star as dwarf.

s026 (O8 V): This source is a blend of three stars identified in Bianchi et
al. (2012)’s catalogue. We assign the photometry and ID of the brightest one.
The spectrum displays strong, broad He ii lines and extremely wide Balmer
lines. The He i 4471 is strongly affected by nebular contamination. Based on
the strength ratio of He ii 4200 with respect to He i+ii 4026, we classify the
star as O8. We assign class V due to the broadening of the Balmer series and
the strong He ii 4686 absorption.

s027 (O8 V): Noisy spectrum. We classify the star as O8 dwarf based on the
strength of He ii 4200 and He ii 4686. However, He ii 4541 is absent. We also
note that all He lines are very broad.

s029 (O8.5 III): This source is s03 in Garcia et al. (2019), but an updated
reduction is presented in this work. The new spectrum shows an emission at
4450 Å due to a slit-ghost. He i 4471 seems double due to nebular
contamination in some observing blocks but was not discarded in the
classification, and He ii 4541 is not seen in all the selected OBs. The target
is marked as a possible SB1 since we detected radial velocity variations of
more than 200 km s-1 in He ii 4541 between different OBs. These variations
were corrected prior to co-adding the spectra of individual OBs. The updated
spectrum exhibits HeII 4200 and HeII 5411, and narrower Balmer lines, hence
the source is reclassified as O8.5 III.

s030 (O8.5 III): The slit was tilted by 7.1°, which results in moderate
degradation of the spectral resolution. This star was observed in the same
slit as s032. Targets are very close to each other, and some contamination is
expected. The stars are located within an H ii bubble and nebular subtraction
was incomplete, as evinced by the core of the Balmer lines.

s031 (O8.5 V): This star was observed within a slit tilted by 8° to include
also star s158. The inclination results in degraded spectral resolution. This
star was close to the slit edge, and sky subtraction using a sky-slit provided
better results than the other stars from the slit. Yet, nebular subtraction is
not fully satisfactory, as demonstrated by the strong emission that overlaps
the wide Balmer lines and the possibly nebular oversubtraction at the core of
He lines. Spectral classification is further hindered by poor S/N as expected
from the faint magnitude. The derived radial velocity is high, suggesting it
may be a binary.

s032 (O8.5 V): This source encompasses two stars unresolved in Massey et al.
(2007)’s catalogue, and it is possibly contaminated by s030, which is also in
the same slit. The slit was tilted by 7.1°, which results in moderate
degradation of the spectral resolution. For the second night, we found that
the sky-slit provided better background subtraction, although that implies
that $\lambda$ < 4100 Å can only be recovered from Night-1, and errors are
larger. Despite this action, nebular subtraction is incomplete in the Balmer
lines, and the He i lines display nebular oversubtraction at their cores.

s033 (O8.5 V): The spectrum shows strong nebular contamination. We assign an
O8.5 type based on the relative strength of He ii 4541 with He i 4471 and the
presence of He ii 5411. The detection of He ii 4686 in absorption and the
broad Balmer lines suggest a dwarf luminosity class.

s034 (O8.5 V): The spectrum displays He ii 4541 that, compared with the He i
4471 line suggest an O8.5 type. He ii 5411 is also detected. Based on the
broadening of the Balmer series, we assign a luminosity class V, but we note
that this is in contradiction with the absence of He ii 4686 in absorption.

s035 (O8.5 V): At the assigned $v_{\mathrm{rad}}$ of 170 km s-1, we detect a
strong absorption of He ii 4541, partially affected by nebular
oversubtraction. We classify the source as O8.5 due to the relative strength
of He i 4144 and He ii 4200. The Balmer series shows nebular contamination and
seems inconsistent with the adopted $v_{\mathrm{rad}}$.

s036 (O9 I): This star was observed in three observing runs, MOS2, LSS2 and
LSS3. No variations in spectral features or radial velocities were detected.
In MOS2, it was included in a very short slit that hindered sky subtraction.
In LSS2, the star was off-slit, and exposure times were shorter, yet the sky
subtraction was slightly better. The He ii 4541 transition is detected
completely in emission in MOS2, but the stellar component is partially
recovered in LSS2 and LSS3. After checking against the spectral lines, we co-
added the spectra to increase the spectral S/N. We note that the He ii 4541
line is lost in the co-added spectrum, but we used the LSS2 line strength for
spectral classification. The strong absorption at 4650 Å is due to an artefact
detected in MOS2 run, but not in the LSS2 or LSS3 runs.

s037 (O9 I): This star was observed in runs LSS1, MOS2 and LSS2. Camacho et
al. (2016) classified this star, its OB623, as an O8 Ib. We detect no
significant morphological or radial-velocity variations in the subsequent MOS2
and LSS2 runs. In MOS2, the spectrum was taken in a too-short slit that
compromised the sky subtraction with apall. For the spectrum of the second
night, with worse effective seeing, we instead used the sky spectrum from one
of the sky-slits, although sky subtraction was incomplete. Sky subtraction in
LSS2 is not fully satisfactory either. After checking against spectral
variations, we co-added the LSS2 and MOS2 spectra to increase the S/N and
discarded the LSS1 spectrum that only introduced noise in the co-addition.
The strong emission close to He ii 4541 is remnant from the sky-spectrum and
precludes the detection of He ii 4541. This line is detected in absorption
both in LSS1 and LSS2. Alternatively, He ii 4200 is clearly observed in the
co-added spectrum and He ii 4686 is still neither detected in emission nor
absorption, thus pointing to class I. Considering the strength of the He ii
4200 and Si iv 4089 lines relative to He i transitions, we reclassify the star
as O9 I.

s038 (O9 I((f))): This star was classified as O9.7 I((f)) by Camacho et al.
(2016), who also noted indications of strong winds and increased He abundance.
In this work, we reclassify the star as O9 I((f)) based on the strength of He
ii 5411 and He ii 4200, more consistently with the current criteria.

s039 (O9 III): An unknown emission detected near He ii 4200, at 4183 Å. The
presence of He ii 4541 and its relative strength with respect He i 4471
suggest an O9 type. The lack of He ii 4686 points to a supergiant class, but
the width of the Balmer series and the lack of Si iv suggest luminosity class
III.

s040 (O9 III + mid-B?): The spectrum is noisy and presents significant nebular
contamination. Because of the relative strength of He ii 4200 with He i 4387,
we assign an O9 type. The source also shows weak He ii 4541 and He ii 5411.
The width of the Balmer lines and the relative strength of He ii 4541 with He
ii 4686 points to the adopted class III, although the faint magnitude suggests
class V. The Mg ii 4481 and Si iii 4552 absorptions would be consistent with a
secondary component at $v_{\mathrm{rad}}$ = 300 km s-1 and mid-B spectral
type. Space photometry from Bianchi et al. (2012) cannot resolve this source.

s041 (O9.5 II): This star was observed in two observing runs, MOS1 and MOS2,
where no radial velocity variation was detected. The poorer S/N of the MOS2
observation prevented the co-addition of the two spectra.

s042 (O9.5 III): An emission at 4450 Å is only seen in the observing blocks
taken in the first year (as in s029), and it may be due to a slit-ghost. In
addition, an anomalous sky line emission is detected at 900 km s-1 of $\rm
H\delta$ and $\rm H\gamma$ in some observing blocks, which leads to an
artificial broadening of the lines. He i 4471 was not considered for
classification since it seems to be an artefact. Apart from the MOS1, the
source was also observed in the LSS3 observing run, where He i 4471 has a much
weaker absorption. No radial velocity variations were detected, but since the
LSS3 spectrum has worse S/N and poorer spectral resolution, the spectra were
not co-added. HST photometry reveals the source is a blend of two stars.

s043 (O9.5 V): The source is a blend of two unresolved stars identified in
Bianchi et al. (2012). It was observed in the LSS2 and LSS3 observing runs,
and we detected no radial velocity variation. Since the S/N of the LSS2
spectrum was poor, the data were not co-added. The spectrum shows strong lines
of He i, perhaps He ii 4200 and no metallic lines. The absorption found at
$\sim$ 4542 Å is an artefact of reduction. Based on the presence of He ii 4200
and the width of the Balmer lines, we classify it as O9.5 V.

s044 (O9.5 V): He i 4471 is absent due to nebular contamination. We clearly
detect the Balmer series, some He i lines, and He ii 4541 and He ii 4686 in
absorption. We classify the star as an O9.5 dwarf due to the relative strength
of He ii 4541 with respect to He i 4387 and the width of the Balmer series.
$\rm H\gamma$, $\rm H\delta$ and He i 4471 show hints of a second component at
higher radial velocity ($\sim$ 600 km s-1).

s045 (O9.5 V): The detection of He ii 4200 and He ii 4541 points to an O9.5
type. He i 4388 is anomalously broad. The star is classified as a dwarf
because of the broadening of the Balmer wings. However, we note that this is
in contradiction with the absence of He ii 4686.

s046 (O9.5 V): This source is a blend of two unresolved stars identified in
Bianchi et al. (2012). We classify it as O9.5 based on the relative strength
of He ii 4200/He ii 4144 since He ii 4541 seems to be affected by noise.
Luminosity class V is assigned due to the width of the Balmer lines.

s047 (O9.5 V): We reclassify the star from an A5 II in Camacho et al. (2016)
to an O9.5 V as we note the presence of He ii 4686 and He ii 4541.

s048 (O9.5 Vn): We classify the star as O9.5 based on the relative strength of
He i 4471 and He ii 4541. Luminosity class Vn is assigned due to the width of
the Balmer lines.

s049 (O9.7 I): This source was observed in two observing runs, LSS2 and LSS3.
We found a radial velocity variation of $\sim$30 km s-1 between the two runs.
However, we did not co-add the spectra since we did not detect any improvement
in the S/N. The spectrum shows He ii 4541 and similarly strong Si iii 4552, we
assigned a compromise spectral type of O9.7 supergiant. HST photometry reveals
this source is a blend of two unresolved stars identified in Bianchi et al.
(2012).

s050 (O9.7 I): This star is OB321 in Camacho et al. (2016), where they
assigned O9.7 I((f)). Since it is brighter than most of the MOS2 targets ($V$
= 19.609 mag), it was observed in a narrow slit to increase resolution and
constrain better its stellar parameters. The S/N of this spectrum is better
than the LSS1 program, so we replaced OB321 with this one in our database. N
iii 4634 does not exhibit emission in the improved quality spectrum, so we
removed the ((f)) qualifier.

s051 (O9.7 I): This star was observed in LSS2 and MOS2 observing runs. We
detect no radial velocity variations and discard the spectrum from the LSS2
run due to poor S/N.

s052 (O9.7 I): This star was observed together with in a very short slit. For
the first night, we found that the subtraction of one of the sky-slits
provided better results. For the second night, apall sky subtraction provided
better results, but only one box was used (located in the middle of the slit).
He ii lines are clearly seen only in the spectrum of the second night, but we
detected no apparent radial velocity variation between both nights.

s053 (O9.7 I): We detect radial velocity inconsistencies between different
Balmer lines and different He i lines, which may be explained by nebular
contamination. Adopting the systemic velocity of the galaxy, the presence of
He ii 4200 and lack of Si iii 4552, Mg ii 4481 and He ii 4686 suggest an O9.7
I. The spectrum may be contaminated by the near star s070.

s054 (O9.7 III): The star was observed in both LSS2 and LSS3 runs. We did not
detect radial velocity variations between the two observations, and we co-
added the spectra to achieve higher S/N. Only He i lines are clearly detected.
There might be hints of He ii 4541, He ii 4686 and Si iv 4089, hence we
classify the star as O9.7. Luminosity class III is assigned due to the width
of the Balmer series.

s055 (O9.7 III): The main star (J101053.90-044111.0) is contaminated by a
nearby one with very blue $Q$-colour (J101053.94-044110.1) and only 0.2 mag
fainter, slightly outside the slit. The selection of sky-boxes with apall was
hampered by this source and, as a consequence, the nebular subtraction is
incomplete, as seen clearly in the core of the Balmer lines, but also He ii
4686. The emission at $\sim$ 4540 Å that affects He ii 4541 has an unknown
origin and does not seem to be a cosmic ray.

s056 (O9.7 V): In the MOS2 run, this star was observed in a very short slitlet
that severely handicapped sky subtraction (e.g. see the central nebular
emission lines in Balmer and possibly He i oversubstraction). It was also
observed in the MOS1 run, but nebular subtraction was similarly poor.
Considering that the star is far from the large ionized structures of the
galaxy, the strong nebular contamination is quite puzzling. Alternatively, the
star could be an Oe star. After checking against spectral line or flux
continuum level variations, we co-added the spectra for increased S/N. There
is an absorption feature at $\sim$ 4200 Å of unknown origin that is detected
in the spectrum of both observing runs, making it unlikely for it to be an
artefact. The subtracted sky spectrum was checked, and there are no features
at those wavelengths. He ii 4541 is present but weak, so the star is
classified as O9.7. Luminosity class V is assigned based on the width of the
Balmer lines. We note that He ii 4686 should be strong in absorption in an
O9.7 V but is missing in the spectrum of MOS2.

s057 (O9.7 V): This star was observed in LSS2 and LSS3 observing runs, where
no radial velocity variations were detected. The spectra were co-added to
improve S/N. We note the absence of He ii 4686, in contradiction with the
luminosity class Vn, indicated by the Balmer wings.

s058 (O9.7 V + B?): Very noisy spectrum. The O9.7 type is assigned because of
the presence of He ii lines and Si iii 4552. The strong He i lines at
$v_{\mathrm{rad}}$ = 150 km s-1suggest a secondary component.

s059 (O9.7 Vn): This star is located in the galactic outskirts, and it is the
farthest of the southern sources from the galactic centre. Contrastingly, it
shows nebular lines at the core of the Balmer series and GALEX UV emission. We
assign spectral type O9.7 based on the detection of He ii 4686 in absorption.
The modifier "n" is used because of the abnormally broad Balmer lines.

s060 (O9.7 V): Possible blend, which hampers the determination of its radial
velocity. The target does not show He ii 4541, but it presents He ii 4200 and
He ii 4686 in absorption. The relative strength of He ii 4200 with respect to
He i 4144 indicates an O9.7 type. Due to the broadening of the Balmer series
and He ii 4686 in absorption, we assign class V. The broad He i lines would be
consistent with the source being a blend or an unresolved binary.

s061 (B0 I): The spectrum presents a deep depression at 4320 Å due to a defect
in the CCD. The star seems to have N and O.

s063 (B0 III): Si iv 4089 is affected by an artefact and therefore cannot be
used for classification. Since there are no He ii lines except for He ii 4686,
and Si iii 4552 is very weak, type B0 is assigned. The relative He ii 4686 to
He i 4713 line ratio indicates luminosity class III.

s064 (B0 III): This star was observed in two runs, MOS2 and LSS3, where no
radial velocity variations were detected. We discarded the combination of the
observations and used only the MOS2 observation since we did not find S/N
improvement in the co-addition. In this observing run, the star was observed
in a very short slit, and sky subtraction was compromised. The wavelength of
the core of the Balmer lines and the strongest He i lines is consistent with
the systemic velocity of the galaxy ($v_{\mathrm{rad}}$ = 325 km s-1). An
instrumental defect at $\sim$ 4700 Å prevents any conclusion on He ii 4686 and
He i 4713. There could be a second component at $v_{\mathrm{rad}}$ $\sim$ 625
km s-1 showing Si iii 4552 and He ii 4541.

s065 (B0 V): The wings of the Balmer series are abnormally broad, hinting it
may be a binary star. The component with the quoted $v_{\mathrm{rad}}$
$\simeq$ 100 km s-1 would be consistent with the blue wing of the Balmer lines
and would show He i lines, but no metallic lines. We note the detection of He
i 4471 and Mg ii 4481 is hampered by artefacts. We tentatively classify this
star as a B0 V. A second component with $v_{\mathrm{rad}}$ $\simeq$ 700 km s-1
would be consistent with the red wing of the Balmer lines and Si iii 4552. The
emissions lines at $\sim$ 4230 Å and $\sim$ 4480 Å are reminiscents of the Be
star VFTS 822 from Evans et al. (2015).