11institutetext: Institute of Astronomy, KU Leuven, Celestijnenlaan 200D,
B-3001 Leuven, Belgium
11email<EMAIL_ADDRESS>22institutetext: Department of Physics,
Faculty of Science, University of Zagreb, 10 000 Zagreb, Croatia
33institutetext: School of Mathematics, Statistics and Physics, Newcastle
University, Newcastle upon Tyne, NE1 7RU, UK 44institutetext: Max Planck
Institute for Astronomy, Königstuhl 17, 69117 Heidelberg, Germany
55institutetext: Astrophysics Group, Keele University, Staffordshire ST5 5BG,
UK

# Observational mapping of the mass discrepancy in eclipsing binaries:
Selection of the sample and its photometric and spectroscopic
properties.††thanks: Based on observations collected at the European Southern
Observatory, La Silla, Chile under program 0106.A-0906(A), and the HERMES
spectrograph mounted on the KU Leuven Mercator telescope, La Palma, Spain.

Andrew Tkachenko Observational mapping of the mass discrepancy in eclipsing
binaries: Selection of the sample and its photometric and spectroscopic
properties.††thanks: Based on observations collected at the European Southern
Observatory, La Silla, Chile under program 0106.A-0906(A), and the HERMES
spectrograph mounted on the KU Leuven Mercator telescope, La Palma,
Spain.Observational mapping of the mass discrepancy in eclipsing binaries:
Selection of the sample and its photometric and spectroscopic
properties.††thanks: Based on observations collected at the European Southern
Observatory, La Silla, Chile under program 0106.A-0906(A), and the HERMES
spectrograph mounted on the KU Leuven Mercator telescope, La Palma, Spain.
Krešimir Pavlovski Observational mapping of the mass discrepancy in eclipsing
binaries: Selection of the sample and its photometric and spectroscopic
properties.††thanks: Based on observations collected at the European Southern
Observatory, La Silla, Chile under program 0106.A-0906(A), and the HERMES
spectrograph mounted on the KU Leuven Mercator telescope, La Palma,
Spain.Observational mapping of the mass discrepancy in eclipsing binaries:
Selection of the sample and its photometric and spectroscopic
properties.††thanks: Based on observations collected at the European Southern
Observatory, La Silla, Chile under program 0106.A-0906(A), and the HERMES
spectrograph mounted on the KU Leuven Mercator telescope, La Palma,
Spain.Observational mapping of the mass discrepancy in eclipsing binaries:
Selection of the sample and its photometric and spectroscopic
properties.††thanks: Based on observations collected at the European Southern
Observatory, La Silla, Chile under program 0106.A-0906(A), and the HERMES
spectrograph mounted on the KU Leuven Mercator telescope, La Palma,
Spain.Observational mapping of the mass discrepancy in eclipsing binaries:
Selection of the sample and its photometric and spectroscopic
properties.††thanks: Based on observations collected at the European Southern
Observatory, La Silla, Chile under program 0106.A-0906(A), and the HERMES
spectrograph mounted on the KU Leuven Mercator telescope, La Palma, Spain.
Nadezhda Serebriakova Observational mapping of the mass discrepancy in
eclipsing binaries: Selection of the sample and its photometric and
spectroscopic properties.††thanks: Based on observations collected at the
European Southern Observatory, La Silla, Chile under program 0106.A-0906(A),
and the HERMES spectrograph mounted on the KU Leuven Mercator telescope, La
Palma, Spain.Observational mapping of the mass discrepancy in eclipsing
binaries: Selection of the sample and its photometric and spectroscopic
properties.††thanks: Based on observations collected at the European Southern
Observatory, La Silla, Chile under program 0106.A-0906(A), and the HERMES
spectrograph mounted on the KU Leuven Mercator telescope, La Palma, Spain.
Dominic M. Bowman Observational mapping of the mass discrepancy in eclipsing
binaries: Selection of the sample and its photometric and spectroscopic
properties.††thanks: Based on observations collected at the European Southern
Observatory, La Silla, Chile under program 0106.A-0906(A), and the HERMES
spectrograph mounted on the KU Leuven Mercator telescope, La Palma,
Spain.Observational mapping of the mass discrepancy in eclipsing binaries:
Selection of the sample and its photometric and spectroscopic
properties.††thanks: Based on observations collected at the European Southern
Observatory, La Silla, Chile under program 0106.A-0906(A), and the HERMES
spectrograph mounted on the KU Leuven Mercator telescope, La Palma,
Spain.Observational mapping of the mass discrepancy in eclipsing binaries:
Selection of the sample and its photometric and spectroscopic
properties.††thanks: Based on observations collected at the European Southern
Observatory, La Silla, Chile under program 0106.A-0906(A), and the HERMES
spectrograph mounted on the KU Leuven Mercator telescope, La Palma,
Spain.Observational mapping of the mass discrepancy in eclipsing binaries:
Selection of the sample and its photometric and spectroscopic
properties.††thanks: Based on observations collected at the European Southern
Observatory, La Silla, Chile under program 0106.A-0906(A), and the HERMES
spectrograph mounted on the KU Leuven Mercator telescope, La Palma, Spain. Luc
IJspeert Observational mapping of the mass discrepancy in eclipsing binaries:
Selection of the sample and its photometric and spectroscopic
properties.††thanks: Based on observations collected at the European Southern
Observatory, La Silla, Chile under program 0106.A-0906(A), and the HERMES
spectrograph mounted on the KU Leuven Mercator telescope, La Palma,
Spain.Observational mapping of the mass discrepancy in eclipsing binaries:
Selection of the sample and its photometric and spectroscopic
properties.††thanks: Based on observations collected at the European Southern
Observatory, La Silla, Chile under program 0106.A-0906(A), and the HERMES
spectrograph mounted on the KU Leuven Mercator telescope, La Palma, Spain.
Sarah Gebruers Observational mapping of the mass discrepancy in eclipsing
binaries: Selection of the sample and its photometric and spectroscopic
properties.††thanks: Based on observations collected at the European Southern
Observatory, La Silla, Chile under program 0106.A-0906(A), and the HERMES
spectrograph mounted on the KU Leuven Mercator telescope, La Palma,
Spain.Observational mapping of the mass discrepancy in eclipsing binaries:
Selection of the sample and its photometric and spectroscopic
properties.††thanks: Based on observations collected at the European Southern
Observatory, La Silla, Chile under program 0106.A-0906(A), and the HERMES
spectrograph mounted on the KU Leuven Mercator telescope, La Palma,
Spain.Observational mapping of the mass discrepancy in eclipsing binaries:
Selection of the sample and its photometric and spectroscopic
properties.††thanks: Based on observations collected at the European Southern
Observatory, La Silla, Chile under program 0106.A-0906(A), and the HERMES
spectrograph mounted on the KU Leuven Mercator telescope, La Palma,
Spain.Observational mapping of the mass discrepancy in eclipsing binaries:
Selection of the sample and its photometric and spectroscopic
properties.††thanks: Based on observations collected at the European Southern
Observatory, La Silla, Chile under program 0106.A-0906(A), and the HERMES
spectrograph mounted on the KU Leuven Mercator telescope, La Palma, Spain.
John Southworth Observational mapping of the mass discrepancy in eclipsing
binaries: Selection of the sample and its photometric and spectroscopic
properties.††thanks: Based on observations collected at the European Southern
Observatory, La Silla, Chile under program 0106.A-0906(A), and the HERMES
spectrograph mounted on the KU Leuven Mercator telescope, La Palma,
Spain.Observational mapping of the mass discrepancy in eclipsing binaries:
Selection of the sample and its photometric and spectroscopic
properties.††thanks: Based on observations collected at the European Southern
Observatory, La Silla, Chile under program 0106.A-0906(A), and the HERMES
spectrograph mounted on the KU Leuven Mercator telescope, La Palma, Spain.

(Received XXX / Accepted XXX)

###### Abstract

Context. Eclipsing spectroscopic double-lined binaries are the prime source of
precise and accurate measurements of masses and radii of stars. These
measurements provide a stringent test of models of stellar evolution that are
persistently reported to contain major shortcomings.

Aims. The mass discrepancy observed for the eclipsing spectroscopic double-
lined binaries is one of the manifestations of shortcomings in stellar
evolution models. The problem reflects the inability of the models to
accurately predict effective temperature and surface gravity or luminosity of
the star for a given mass. Our ultimate goal is to provide an observational
mapping of the mass discrepancy and propose a recipe for its solution.

Methods. We initiate a spectroscopic monitoring campaign of 573 candidate
eclipsing binaries classified as such based on their TESS light curves. In
this work, we present a sub-sample of 83 systems for which orbital-phase
resolved spectroscopy has been obtained and subsequently analysed with the
methods of least-squares deconvolution and spectral disentangling. In
addition, we employ TESS space-based light curves to provide photometric
classification of the systems according to the type of their intrinsic
variability.

Results. We confirm 69 systems as either spectroscopic binaries or higher-
order multiple systems. Twelve stars are classified as single and two more
objects are found at the interface of their line profile variability being
interpreted as due to binarity and intrinsic variability of the star.
Moreover, 20 eclipsing binaries are found to contain at least one component
that exhibits stellar oscillations. Spectroscopic orbital elements are
obtained with the method of spectral disentangling and reported for all
systems classified as either SB1 or SB2. The sample presented in this work
contains both detached and semi-detached systems and covers a range in the
effective temperature and mass of the star of $T_{\rm eff}$ $\in[7000,30000]$
K and $M\in[1.5,15]$ M⊙, respectively.

Conclusions. From the comparison of our own results with those published in
the literature for well-studied systems, we conclude an appreciable capability
of the spectral disentangling method to deliver precise and accurate
spectroscopic orbital elements from as few as 6-8 orbital phase-resolved
spectroscopic observations. Orbital solutions obtained this way are accurate
enough to deliver age estimates with accuracy of 10% or better for
intermediate-mass F-type stars, an important resource for calibration of
stellar evolution models for future space-based missions like PLATO. Finally,
despite its rather small size relative to 573 system that we will ultimately
monitor spectroscopically, the sample presented in this work is already
suitable to kick off observational mapping of the mass discrepancy in
eclipsing binaries.

###### Key Words.:

methods: data analysis – methods: observational – techniques: spectroscopic –
(stars:) binaries: eclipsing – (stars:) binaries: spectroscopic – stars:
oscillations

##  1 Introduction

Eclipsing spectroscopic double-lined (hereafter eSB2) binaries are the prime
source of precise and accurate mass measurements of stars, the parameter that,
along with the initial composition of the star determines its evolutionary
path and fate. Owing to the fact that their mass determination is largely
model-independent, eSB2s represent an ideal test bed for the theory of stellar
structure and evolution (SSE).

One of the striking outcomes of the SSE model test is the mass discrepancy.
The term was introduced originally by Herrero et al. (2000) in the context of
single stars, where the authors report a discrepancy between stellar masses
inferred from the spectroscopic properties of stars and those found from the
comparison of the position of stars in the Hertzsprung–Russell diagram (HRD)
with evolutionary tracks. Although the authors conclude that the mass
discrepancy is an observational manifestation of shortcomings in SSE models,
it is worth noting that the spectroscopic masses are also model-dependent as
they are inferred from the spectroscopic estimates of $\log\,g$ and radii of
the stars, both largely dependent on physics employed in the models of stellar
atmospheres and winds.

The conclusions regarding possible imperfections in SSE models are more sound
when based on a comparison between model-independent masses inferred from the
binary dynamics (dynamical masses hereafter) and those dictated by SSE models
(evolutionary masses hereafter). Although earlier works suggest a good
agreement between the dynamical and evolutionary masses for stars up to 25 M⊙
(e.g. Burkholder et al. 1997), more recent studies report a systematic
discrepancy between the two types of measurement. Notably, two independent
studies of the V380 Cyg system (Guinan et al. 2000; Pavlovski et al. 2009)
found a large mass discrepancy for the more evolved primary component. While
Guinan et al. (2000) claim the discrepancy can be resolved with an increased
efficiency of the near-core mixing in the form of core overshooting, Pavlovski
et al. (2009) focus on the effect of rotation in SSE models and conclude that
it cannot explain the observed mass discrepancy. Moreover, Tkachenko et al.
(2014b) re-investigate the V380 Cyg system based on space-based photometric
and high-resolution spectroscopic data and find the mass discrepancy for both
binary components. The authors conclude that the discrepancy can only be
resolved when the combined effect of rotation and core overshooting is
included in SSE models.

The need for a more efficient near-core mixing becomes one of the streamlined
theoretical hypotheses to explain the origin of the mass discrepancy in
eclipsing binaries. Massey et al. (2012) and Morrell et al. (2014) argue the
need for an increased efficiency of interior (rotational or core overshoot)
mixing in SSE models to accommodate a small (some 10%) but systematic
discrepancy between dynamical and evolutionary masses in five high-mass
eclipsing binaries in the LMC. The same conclusion is presented by Pavlovski
et al. (2018) based on the study of four high-mass eclipsing binaries in the
Galaxy. Furthermore, Rosu et al. (2020, 2022a, 2022b) invoke a simultaneous
measurement of the stellar parameters, rate of the apsidal motion and $k_{2}$
internal structure constant in the HD 152248, HD 152219, CPD-41° 7742, and HD
152218 binary systems to demonstrate the lack of efficiency of interior mixing
in standard SSE models, in line with the conclusions drawn by Guinan et al.
(2000) for the V380 Cyg system.

Claret & Torres (2016, 2017, 2018, 2019) present a series of papers focusing
on the inference of the efficiency of near-core mixing in the form of the
convective core overshooting in a sample of some 50 eSB2 systems from Torres
et al. (2010) and the catalogue of the Optical Gravitational Microlensing
Experiment (OGLE)111http://ogle.astrouw.edu.pl/. The authors consider two
implementations of the core overshooting in Granada (Claret 2004, 2012) and
mesa (Paxton et al. 2011, 2013, 2015, 2018, 2019) models, i.e. the convective
penetration (Zahn 1991) and exponential diffusive approximation (Freytag et
al. 1996; Herwig 2000). Irrespective of the prescription used, the authors
report a strong dependency of the inferred overshooting parameter on the
stellar mass such that there is an almost linear transition from no
overshooting to approximately $\alpha_{\rm ov}$ ($f_{\rm ov}$) = 0.2 (0.02)
$H_{\rm p}$ in the mass range from some 1.2 to 2.0 M⊙. The authors find that
the distribution flattens for stars more massive than 2.0 M⊙. Viani & Basu
(2020) investigate an asteroseismic sample of stars observed with the Kepler
mission in the mass range between 1.1 and 1.5 M⊙. The authors report a strong
positive correlation between the inferred overshooting parameter with the
stellar mass.

Costa et al. (2019) revisit the sample of Claret & Torres (2016, 2017, 2018,
2019) using a combination of the Bayesian statistical framework and parsec
stellar evolution models. The authors confirm the need for a mild convective
core overshooting for stars with masses $M\geq 1.9$ M⊙ and, additionally, find
a large spread in the inferred $\alpha_{\rm ov}$ values ranging from some 0.3
to 0.8 $H_{\rm p}$. Costa et al. (2019) arrive at the conclusion that while
the low boundary of the inferred $\alpha_{\rm ov}$ values is likely to be
explained by an insufficient amount of core-boundary mixing in standard SSE
models, the large spread in the inferred values is likely a manifestation of
the natural distribution of initial rotation rates, and hence active
rotational mixing in these intermediate-mass stars.

Finally, Tkachenko et al. (2020) present a study of eleven high-mass eSB2
systems for which high-precision stellar quantities were derived in a
homogeneous way (Pavlovski et al. 2018, 2023; Tkachenko et al. 2014b; Johnston
et al. 2019a). No dependence of the observed mass discrepancy on the stellar
mass is reported but a strong anti-correlation with the surface gravity of the
star is found. The same correlation persists when the stellar mass is replaced
with the mass of the convective core of the star, the finding that the authors
interpret as the need for larger convective core masses in standard models of
SSE. Remarkably, Johnston (2021) arrived at the same conclusion from the study
of a combined large sample of binary components and asteroseismically active
single stars. Furthermore, Martinet et al. (2021) investigate a large sample
of stars from literature spanning a mass range from 7 to 25 M⊙ and using a
large grid of SSE models computed for different values of the overshooting
parameter and initial rotation, and under the assumptions of a moderate and
strong angular momentum transport in the models. The authors confirm the
findings by Tkachenko et al. (2020) and Johnston (2021) and report the need
for larger convective cores at higher stellar masses.

An alternative hypothetical explanation of the mass discrepancy comes from the
nature of binarity itself and the fact that the majority of high-mass stars
are expected to be found in binary systems and/or have experienced binary
interactions in the past (e.g., Sana et al. 2012, 2014). Mahy et al. (2020)
investigate 26 eclipsing binary systems in the LMC and report a good agreement
between their dynamical and spectroscopic masses whereas the evolutionary
masses are found to be systematically overestimated. Upon a closer inspection
of the obtained results, the authors find the mass discrepancy to be more
pronounced in semi-detached binary configurations, the finding that suggests
that binary interactions might be (at least partially) responsible for the
observed mass discrepancy.

Despite being known for some three decades and many attempts made to
observationally quantify and pinpoint the most likely theoretical explanation
for it, the mass discrepancy persists as a phenomenon and the only definite
statement the stellar astrophysics community can make about it is that the
problem exists. The majority of previous studies focus on a particular mass
regime and there is a large diversity in the employed data analysis methods
and input physics used in SSE models, which makes it non-trivial to account
for systematic uncertainties in an attempt to quantify and ultimately resolve
the mass discrepancy problem. Motivated by the current state-of-the-art, we
have initiated a systematic search for intermediate- to high-mass (spectral
types OBAF and masses larger than some 1.2 M⊙) eclipsing binaries in TESS
(Ricker et al. 2015) space-based photometric data and have organised a
complementary ground-based spectroscopic follow-up campaign of the detected
candidate eclipsing binaries. Our goal is to build a stellar sample that
covers the entire mass range of stars born with the convective core, all the
way up to some 30 M⊙, and includes binaries with and without pulsating
component and in different orbital configurations.

Our newly compiled sample is a generalisation of the binary sample presented
in Tkachenko et al. (2020) to cover a much larger range of the stellar mass,
account for intrinsic variability of stars in binaries, and for different
orbital configurations such as detached versus semi-detached systems. We
ultimately aim at a self-contained and homogeneous detailed analysis of all
suitable eSB2, such that the analysis methods remain largely the same, thus
minimising systematic uncertainties on the inferred parameters and
observational mapping of the mass discrepancy across the HRD. In this paper,
we present the sample selection of candidate eclipsing binary systems and
provide an overview of the ground-based spectroscopic observations we have
obtained so far (Section 2). We provide a more detailed look into the
spectroscopic properties of the observed systems with the methods of least-
squares deconvolution and spectral disentangling in Section 3. Analysis of the
TESS photometric data with the aim to provide a photometric classification
according to the type of intrinsic variability of the sample stars is
presented in Section 4. We proceed with a discussion of the obtained results
in the context of the previous studies (Section 5) and close the paper with
some concluding remarks and a list of future prospects (Section 6).

##  2 Sample selection

Our sample selection is based on the work of IJspeert et al. (2021) that
presents a magnitude limited catalogue of eclipsing binaries observed by the
TESS mission. While a detailed description of all the analysis steps is
provided in IJspeert et al. (2021), here we provide a brief summary of
relevance for the rest of the analysis in this work.

IJspeert et al. (2021) start with a global all-sky selection of OB(A)-type
stars from the TESS Input Catalogue (version TICv8 Stassun et al. 2019). The
authors limit the selection to TESS magnitudes below 15 and exclude all
targets that are resolved in the Gaia DR2 catalogue (Gaia Collaboration et al.
2018) but are recognised as a single object in the 2MASS catalogue. The 2MASS
(J-K)–(J-H) colour-colour plane is employed to select some 205 000 OB(A)-type
candidate stars with the exact colour cuts informed by the OB-type star
variability samples of Bowman et al. (2019a) and Pedersen et al. (2019) (cf.
Figure 1 in IJspeert et al. 2021). In the next step, the authors exclude white
dwarfs and giants from the sample using dedicated Gaia flags and limit
themselves to stars for which either TESS-SPOC222DOI: 10.17909/t9-wpz1-8s54 or
MIT QLP333DOI: 10.17909/t9-r086-e880 light curves could be retrieved from the
MAST database444http://archive.stsci.edu/tess/all_products.html. This step
reduces the stellar sample of unique targets to some 91 000. Ultimately, in a
search for eclipsing binaries in the selected sample of OB(A)-type stars with
TESS light curves, IJspeert et al. (2021) develop an algorithm for automated
detection of eclipses and, if detected, determination of the binary orbital
periods. The authors report 5502 candidate eclipsing binaries of which 2077
and 3425 are found to be false positives and true eclipsing binaries,
respectively.

###  2.1 Spectroscopic observations

From the above-mentioned sample of 3425 eclipsing binaries, we select all
targets with V magnitude below 11 and accessible from the Roque de los
Muchachos Observatory (La Palma, Spain) in the course of the year. A total of
545 systems are proposed for observations with the High-Efficiency and high-
Resolution Mercator Echelle Spectrograph (hermes; Raskin et al. 2011)
instrument attached to the 1.2-m Mercator telescope, as part of a dedicated
large programme. For each of those systems, we request eight epochs of
spectroscopic observations with uniform sampling of the orbital phase. The
latter is computed from the orbital period value derived by IJspeert et al.
(2021) and a reference date $T_{0}$ taken to be the date of the first
spectroscopic observation. A few systems have been followed up more
extensively to resolve the effect of intrinsic stellar oscillations on
spectral line profiles, with the time-series ranging from a couple of dozen
spectra to almost 200 measurements (e.g. the case of 16 Lac in Table
LABEL:Tab:targets; see Southworth & Bowman 2022).

At the time of writing, we have completed spectroscopic observations with the
hermes instrument for 58 systems from our sample. The criteria of completion
are: (i) individual epoch observations meet the minimum signal-to-noise-ratio
(S/N) requirement of 60 per resolution element, and (ii) a quasi-uniform
orbital phase coverage to maximise the chances for detecting lines of a
companion star in the observed composite spectra. The former requirement
guarantees S/N in excess of 100 in the disentangled spectra of SB2 systems or
the combined spectrum for SB1 binaries. Table LABEL:Tab:targets gives an
overview of the sample of 58 systems with columns 1 through 7 indicating a
star name and a catalogue identifier (col. 1-3), orbital period (col. 4),
number of spectra acquired (col. 5), and spectroscopic (col. 6) and
photometric (col. 7) classifications obtained in this work.

A complementary, albeit smaller, sample of 28 southern hemisphere eclipsing
binaries is proposed by us for spectroscopic observations with the Fiber-fed
Extended Range Optical Spectrograph (feros; Kaufer et al. 1999) instrument
attached to the MPG/ESO 2.2-m telescope (under program 0106.A-0906(A)).
Similar to the hermes-based programme, we request eight epochs of
spectroscopic observations and use exactly the same completion criteria for
all the observed systems. The criteria are met for the total of 25 systems
while the remaining 3 stars have received only 2 epochs of spectroscopic
observations. Detailed information is provided in Table LABEL:Tab:targets for
the 25 systems that we consider completed, with the meaning of the columns
being the same as for the hermes sample.

The hermes spectra are reduced with version 7.0 of the dedicated hermes
pipeline while the feros spectra are processed with a modification of the
ceres pipeline (Brahm et al. 2017) presented in Gebruers et al. (2022). The
data reduction steps are the standard ones and include bias subtraction, flat
fielding, cosmic ray removal, wavelength calibration, barycentric correction,
and order merging. Normalisation to the pseudo-continuum is done by selecting
knot points and fitting a low-degree polynomial through them. Special care is
taken in the regions of broad Balmer lines to ensure the outer line wings are
not altered by the normalisation process.

Figure 1: Examples of the LSD profiles computed with the $T_{\rm eff}$=11 000
K mask. From top left to bottom right: HD 34382 - a single star showing line
profile variability, HD 138305 - an SB2 binary system, HD 234650 - an SB3
triple system, and HD 57158 - an SB4 quadruple system. The individual LSD
profiles in each panel have been shifted vertically by a constant factor for
clarity.

Figure 2: Example orbital solutions for SB2 systems with differing numbers of
observed spectra and fairly good phase coverage. Black lines depict the
orbital solutions while blue and red dots represent orbital phases of the
observed spectra for the primary and secondary components, respectively. V436
Per, $\alpha$ CrB and NY Cep are highly eccentric binary systems with
eccentricities of about 0.37, 0.38 and 0.44, respectively. RX Dra represents
the case of a circular-orbit system with the previously undetected secondary
component that is revealed in this work.

Figure 3: Examples of the disentangled spectra for the binary systems shown in
Fig 2. In the upper row, V436 Per and $\alpha$ CrB represent systems with
extreme light ratio. In the lower row, binary systems with hot (NY Cep) and
cool (RX Dra) components in our sample are shown for comparison.

##  3 Spectroscopic analysis

In this section, we present the results of our preliminary spectroscopic
analysis of 83 systems, that is 58 in the hermes sample and 25 in the feros
sample. We start with spectroscopic classification based on the least-squares
deconvolution profiles (Section 3.1). The classification is validated and, if
necessary, refined with the method of spectral disentangling which is
subsequently used to compute the spectroscopic orbital elements of confirmed
binary star systems (Section 3.2).

###  3.1 Least-squares deconvolution

We employ the method of least-squares deconvolution (LSD; Donati et al. 1997)
as implemented by Tkachenko et al. (2013). The LSD method offers an efficient
way to compute a mean profile from an ensemble of spectral lines present in
the spectrum. Two fundamental assumptions of the method need to be kept in
mind when computing an LSD profile: (i) all spectral lines should have similar
shape, meaning that hydrogen, helium, and metal lines with damping wings need
to be excluded from the calculations, and (ii) all spectral lines add up
linearly implying that the depths of spectral blends formed of lines whose
absorption coefficients overlap are not accurately reproduced by the method.
Within these fundamental assumptions, stellar spectrum $I$ is represented as a
convolution of a line mask $M$ with a priori unknown average profile $Z(v)$,
such that

$I=M\ast Z(v).$ (1)

The line mask consists of a set of delta functions that specify wavelength
positions and theoretical strengths of spectral lines. This implies that the
exact content of a line mask depends on the assumed atmospheric parameters
($T_{\rm eff}$, $\log\,g$, $\xi$, [M/H]) of the star, where the effective
temperature has by far the most dominant effect. For the purpose of the
analysis in this work, and informed by the spectral classification as
OB(A)-type stars in IJspeert et al. (2021), we compute line masks for two
values of $T_{\rm eff}$, 8 000 K and 11 000 K, and use both of them for the
calculation of the LSD profiles from every single spectrum of every object in
the sample. The two masks are required and sufficient at the same time in
order to account for the variable complexity of the stellar spectrum between
these two temperature regimes.

We perform a visual classification of all systems based on the time series of
their LSD profiles. Figure 1 shows time-series of the LSD profiles for four
objects, ranging from a single star to a quadruple system. HD 34382 (top left)
is classified by us as a single star exhibiting line profile variations
(LPVs). The latter are caused by the intrinsic variability of the star either
due to rotational modulation of inhomogeneities on its surface or stellar
pulsations. HD 138305 (top right) is found by us to be an SB2 system composed
of two similar stars. HD 234650 (bottom left) and HD 57158 (bottom right) are
higher-order multiple systems with the former being a spectroscopic triple-
lined (SB3) system while the latter being a spectroscopic quadruple-lined
(SB4) system.

###  3.2 Spectral disentangling

The method of spectral disentangling (SPD) was originally introduced and
formulated by Simon & Sturm (1994). The method allows for a simultaneous,
self-consistent inference of the orbital elements for spectroscopic binary and
higher-order multiple systems and the disentangled spectra of the components
these systems are composed of. SPD relies on the principle that the complex
spectrum of a binary or multiple-star system is a linear combination of the
individual components’ spectra, each shifted due to the Doppler effect in the
course of the orbital cycle and diluted by the component’s fractional light
contribution.

In the original formulation of Simon & Sturm (1994), the disentangling problem
is described by the following matrix equation:

$\mathbf{M}\cdot\mathbf{x}=\mathbf{c},$ (2)

where the single-column matrices $\mathbf{x}$ and $\mathbf{c}$ represent a
priori unknown disentangled spectra of the components and the time-series of
the observed spectra, respectively. The rectangular matrix $\mathbf{M}$
represents a linear transformation of $\mathbf{x}$ to $\mathbf{c}$, has a
block structure, and is constructed using individual Doppler shifts for each
of the components at each epoch of the observations. Matrix $\mathbf{M}$ also
contains information on the light dilution factors for each component of the
system and each observed spectrum in the time-series. In the case of high-
resolution spectra and when the number of observed spectra exceeds the number
of individual components, the system of linear equations is overdetermined
(i.e. number of equations is greater than the number of unknowns).
Mathematically, this is an ill-posed problem and regularisation conditions are
required to solve it. Simon & Sturm (1994) employ an algebraic technique known
as a singular value decomposition (SVD). Solving the system is typically a
computationally demanding problem owing to the need to invert matrix
$\mathbf{M}$ in Eq. (2).

Independently, a Fourier-based formulation of spectral disentangling is
developed by Hadrava (1995). With the help of the the Fourier transform, a
large set of equations can be uncoupled into many small sets of equations for
each Fourier mode. Such an approach is significantly less computationally
demanding than the original formulation. In either formulation, the SPD method
can be exploited in two basic modes: spectrum disentangling, where a priori
unknown components’ spectra are optimised along with the orbital elements of
the system, and spectrum separation, where RVs or orbital elements are assumed
to be known and individual component’s spectra are inferred from the
respective fixed orbital solution (e.g. Pavlovski & Hensberge 2010; Serenelli
et al. 2021).

The SPD method builds on the following fundamental assumptions: (i) the
dominant variability present in the observed composite spectra of a system
should be due to the motion of the components around their common centre of
mass; and (ii) observationally the orbital phase coverage should be complete
and close to uniform to exploit the full power of the method. In this work, as
well as in the future exploitation of our sample (cf. Section 6), we are
interested in the detection of high-contrast binary SB2 systems (i.e. where
one of the binary components is significantly fainter than the other),
including those where one of the components is a pulsating star. While the
former represents a challenge for the method, the latter is, in principle, a
direct violation of one of the method’s fundamental assumption. Hence, we
comment on both of these challenges in more detail to justify the choice of
the SPD method for our purpose.

In the method of SPD, the gain in S/N in the final disentangled spectra is
proportional to $\sqrt{N}$, where $N$ is the number of the observed spectra in
a time-series. The gain is distributed according to the light ratio between
the binary components. Previous studies demonstrate that SPD is a powerful
tool to reveal the presence of a faint component (1–2% contribution to the
total light of the system) in high-contrast systems (e.g. Kolbas et al. 2015;
Themeßl et al. 2018; Pavlovski et al. 2022; Johnston et al. 2023).

While time-dependent LPVs due to the effect of non-radial pulsations represent
an extra source of variability that is not recognised by the SPD method,
previous studies of binary systems with pulsating component(s) demonstrate
that the method delivers reliable results provided that: (i) the LPV amplitude
due to the intrinsic variability of the star is much smaller than due to its
binary motion, and (ii) the orbital period of the system differs from that of
the dominant intrinsic variability of the star (e.g. Uytterhoeven et al.
2005a, b; Ausseloos et al. 2006; Tkachenko et al. 2012). For the case of a
radially pulsating star where the RV amplitude of the dominant (radial)
pulsation mode is comparable to or exceeds that of the orbital motion, a large
set of observed spectra with uniform coverage of the pulsation and orbital
cycles is required. Indeed, in such a scenario, a time series of the observed
spectra can be binned with respect to the orbital phase of the binary system
to average out the effect of oscillations on the line profile of a pulsating
component. As a result, spectral line profiles in the orbital phase-binned
observed spectra contain a minimum amount of distortion due to stellar
pulsations, bringing us to the former case of orbital motion being the
dominant source of the apparent line profile variability (e.g., Tkachenko et
al. 2014a, 2016).

In this work, we employ the Fourier-based method of SPD as implemented in the
FDBinary software package555http://sail.zpf.fer.hr/fdbinary/ (Ilijic et al.
2004). The disentangling method is used by us to verify and if necessary
refine the LSD-based spectroscopic classification and, more importantly,
derive a preliminary set of orbital elements for all systems in our sample.
The most common, yet not frequent, case of a refined classification is the one
of an SB1 (LSD-based) to an SB2 (SPD-based) system. This is thanks to the
power of the SPD method to detect spectral contributions as faint as a few
percent (in the continuum flux units) which may stay unnoticed in the LSD
profiles should not the most optimal line mask be chosen for the respective
calculations.

The results of our combined LSD- and SPD-based spectroscopic classification
are provided in Table LABEL:Tab:targets (sixth column). For all SB1 and SB2
systems, we employ the FDBinary code to compute their spectroscopic orbital
elements, that is orbital eccentricity $e$, argument of periastron $\omega$,
and semi-amplitudes of the individual components $K_{\rm i}$. Owing to the
fact that the orbital period value can be constrained with typically a much
higher precision from the eclipses in space-based photometric data than from a
limited set of ground-based spectroscopic observations, we do not optimise the
orbital period value in the SPD method by default. Instead, we improve the
period separately using an algorithm described in detail in IJspeert et al.
(in prep.) that searches globally for the best orbital period and then refines
it with a granularity of one part in one hundred thousand. The global merit
function is a combination of the phase dispersion measure (Stellingwerf 1978)
and Lomb-Scargle amplitude (Lomb 1976; Scargle 1982), as well as the number
and completeness factor of orbital harmonic frequencies found in the light
curve. This period finding algorithm is part of a larger methodological
framework for analysing eclipsing binary light curves developed by IJspeert et
al. (in prep.) and implemented in STAR_SHADOW666STAR_SHADOW will be part of
the mentioned publication and published on GitHub. Furthermore, eccentricity
is optimised by us in those case only where evidence for a value different
from zero is found in the TESS light curve. Ultimately, the SPD solution is
computed in three different wavelength regions free of Balmer lines and rich
in metal lines, 4375-4490 Å, 4510-4725 Å, and 4900-5230 Å. The orbital
elements and their uncertainties are reported in Table LABEL:Tab:SPDOrbit and
are the mean and standard deviation of the mean values, respectively. We note
that a typical S/N of about 80 or higher has been achieved in all
spectroscopic epochs for all systems in the studied sample. This is
sufficiently high to have a negligible effect on the estimated uncertainties
of the obtained orbital elements compared to other contributors to the total
error budget. These dominant contributors to the uncertainties are: (i) flux
ratio of the binary components where lower flux contribution implies larger
uncertainties on the K-value of the star and eccentricity of the system; (ii)
spectral types of the binary components where a high density of lines in the
spectra of cooler stars leads to more appreciable constraints on the K-values
of stars and eccentricity of the orbit; and (iii) spectral line broadening
with narrow spectral lines being more constraining on individual spectral
shifts in the Fourier space and, as a result, on the eccentricity of the
system and K-value of the respective binary component. While we opt for a
simple method of uncertainty estimate in this work, our future studies will
include a more sophisticated way of estimating uncertainties (e.g., the
bootstrap method or jackknife estimator) that will allows us to account for
the above-mentioned contributors to the total error budget.

Figure 2 shows example orbital solutions for a selection of four systems from
the studied sample. Black lines represent orbital solutions for both binary
components as determined with the SPD method. Red and blue filled circles
indicate, respectively, the predicted primary and secondary RVs from the
obtained orbital solution and at the epochs of the acquired spectroscopic
observations. We note that the method of SPD bypasses the step of RV
determination from the observed spectra, hence the RVs shown in Figure 2 are
inferred quantities and used to illustrate typical phase coverage we achieve
for our systems irrespective of their orbital eccentricity value. Examples of
the disentangled spectra for the same four systems are shown in Figure 3. In
the top row, we show a system with two similar components (V436 Per, light
ratio $l_{\rm B}/l_{\rm A}=0.798\pm 0.021$; Southworth & Bowman 2022) and a
system with an extreme light ratio between the two stars ($\alpha$ CrB, light
ratio $l_{\rm B}/l_{\rm A}=0.0180\pm 0.0002$, Pavlovski et al., in
preparation). In the latter case, spectral lines of the secondary component
have a depth of at most 0.5% in the continuum units owing to their rotational
broadening. In the bottom row, the disentangled spectra are shown for the NY
Cep (left) and RX Dra (right) systems. The primary component in the detached
eclipsing binary NY Cep is one of the hottest stars in our sample which is
evidenced by the presence of He ii lines at 4541 Å and 4687 Å. SPD also
reveals the spectrum of the secondary component in the RX Dra system (right)
which makes it an SB2 system with a $\gamma$ Dor pulsating component.

##  4 Photometric analysis

Previous efforts have demonstrated the power of combining high-quality TESS
light curves with spectroscopic radial velocities for relatively small numbers
of high-mass pulsating EBs (see e.g. Bowman et al. 2019b; Southworth et al.
2020, 2021; Southworth & Bowman 2022). As discussed in Section 2, our large
sample of EBs was originally identified as eclipsing systems by IJspeert et
al. (2021) using the products of automated light curve extraction tools, such
as the MAST SPOC pipeline (Jenkins et al. 2016) and/or the TESS quick-look
pipeline (QLP). However, in our current work we extract a light curve
optimised for each pulsating EB using the TESS full-frame images (FFIs)
following a similar methodology described in Bowman et al. (2022), which is
briefly summarised here for completeness.

TESS pixel cutouts sized $25\times 25$ are extracted from the FFIs using
astrocut (Brasseur et al. 2019). We use the lightkurve software package
(Lightkurve Collaboration et al. 2018) to select aperture masks optimised to
maximise the signal-to-noise of the eclipses (and pulsations if present)
whilst minimising the contribution from nearby contaminating sources based on
cross-checking with the location and fluxes of sources in the Gaia catalogue.
The background flux is estimated from the median observed flux per frame after
excluding pixels that contain flux from the target or other sources. The
background flux is subtracted, and the extracted light curves are normalised
by dividing through the median flux. Finally, we perform a principal component
analysis (PCA) to remove remaining instrumental systematics in the light
curves (Bowman et al. 2022).

Due to the difference in integration times between different TESS data cycles
and the large (i.e. $>1$ yr) gaps between cycles, we choose to analyse data
from each cycle, if available, separately for each star. To classify the
dominant variability beyond the eclipses caused by binarity, and ascertain if
any additional signal is caused by pulsations or rotational modulation
(RotMod) in our sample, we create a multi-frequency cosinusoid model to remove
the orbital harmonics from the extracted TESS light. Our approach is similar
to the method previously used by Bowman et al. (2019b) for the pulsating
eclipsing binary system U Gru. After optimising the orbital period and the
number of significant orbital harmonic frequencies using a least-squares fit
to the light curve, we subtract the resultant optimised analytical model from
the light curve and calculate the residual amplitude spectrum (Kurtz 1985).
The dominant variability is identified as pulsations or RotMod based on visual
inspection of the residual amplitude spectrum. We apply this method to all the
combined sectors within each TESS cycle to compare our classifications. For
all stars, we find that our classification results are consistent across the
multiple cycles, if available.

An example figure summarising this methodology is provided in Fig. 4, which
contains our FFI-extracted light curve for all available cycle 1 data, the
corresponding binary orbit phase-folded light curve, and the original and
residual amplitude spectra for TIC 150442264 (HD 46792). In this example,
several significant g-mode frequencies can be seen around 1 d-1, and also a
series of harmonics typical of rotational modulation at 0.1 d-1. There are no
significant pulsation frequencies above about 2 d-1. We leave the combined
binary modelling and pulsational frequency analysis for a future paper, but
this approach allows us to classify the dominant pulsational variability in
our targets and identify promising systems for future asteroseismic modelling,
which are listed in Table LABEL:Tab:targets. In so doing, we dramatically
increase the number of pulsating EBs which span spectral types from O through
to F, and include $\beta$ Cep through to $\gamma$ Dor pulsators.

Figure 4: Example summary figure for TIC 150442264. Top panels: FFI-extracted
light curve versus time (left) and phase-folded on the orbital period for all
available cycle 1 sectors (right). Bottom panels: Original and residual
amplitude spectra in which vertical red lines denote the location of
significant orbital harmonics that comprise the analytic binary model, which
is also shown as a red line in the phase-folded light curve panel.

##  5 Discussion

In this section, we present a summary of the obtained results and discuss them
in the context of the previous studies. We recall that our sample selection
was done in a blind way based on colour information and available TESS light
curves. This approach to the target selection guarantees the presence in our
sample of some of the well-studied systems in the literature allowing us to
use those as the methodology validators.

We report the detection of 12 single stars that are found by us to show either
LPVs or no variations at all. We note that two more systems are found at the
interface of their (low-amplitude) LPVs being interpreted as due to binarity
and intrinsic variability of the star. Fifteen stars are specrtoscopically
classified by us as apparent SB1 systems. We note that the SB1 classification
is not set in stone and may be revised into an SB2 when/if more spectra become
available in the future. The method of spectral disentangling delivers average
disentangled spectra with a gain in S/N of roughly a factor $\sqrt{N}$ with
respect to the individual observed composite spectra, where $N$ stands for the
number of spectra in a time-series. This implies that the detection limit of a
hypothetical faint companion star with the SPD method increases with the
number of observed spectra in a time-series which may ultimately lead to
reclassification of a star from SB1 into SB2. Furthermore, 50 targets are
classified by us as SB2 systems. These are the primary candidates for future
studies in the context of the observational mapping of the mass discrepancy.
Finally, we report the detection of four high-order multiple systems, of which
two objects are classified by us as triple (SB3) and two as quadruple (SB4)
systems.

Twenty eclipsing systems are classified by us as presumably containing at
least one oscillating component. Of these twenty systems, one (u Tau) is a
tentative detection where we cannot exclude rotational modulation as a cause
of the intrinsic variability of the star, and two systems (V436 Per and HD
84493) are likely to contain a component showing stochastic low-frequency
variability (e.g., Bowman et al. 2019a, 2020). Sixteen systems are found to
contain either a g- (eight binaries) or p-mode (eight binaries) pulsator, and
one system (V350 And) contains a g-/p-mode hybrid pulsator. Finally, the CD-45
4393 system is found by us to show a steadily decreasing eclipse depth across
different TESS sectors such that the signal practically disappears in sectors
61 and 62. Such a phenomenon can be a manifestation of the precession of the
binary orbit (e.g. due to the presence of a third body in the system) that
causes the orbital inclination angle with respect to the observer’s line-of-
sight to change with time.

Below, we discuss systems that have been studied in the literature and, where
applicable, compare our photometric and spectroscopic results with those
reported by other research groups. In doing that, we focus primarily on
dedicated studies rather than on large catalogues, hence the overview provided
below is probably not inclusive.

BD+13 1880: Šubjak et al. (2020) report the system to comprise a metallic-line
Am star and a brown dwarf companion in a 3.6772$\pm$0.0001 d orbit. From the
analysis of a set of some 50 spectra obtained with multiple high-resolution
spectrographs, the authors report the RV semi-amplitude $K$ of the host star
of 4.64$\pm$0.03 km s-1, in a good agreement with the value of 5.0$\pm$0.1 km
s-1 derived in this work from a time-series of 10 hermes spectra.

BD+36 3317: Özdarcan et al. (2012) classify the system as an Algol-type binary
based on the analysis of ground-based multi-colour photometry, and provide
evidence of the system being a member of the $\delta$ Lyrae cluster. The
system is revisited by Kıran et al. (2016) based on photometric observations
from the literature and a newly obtained time-series of 20 medium-resolution
($R=11\,700$) spectra. The authors reinforce the conclusion on the cluster
membership of the system, and report the RV semi-amplitudes of the primary and
secondary to be $K_{\rm 1}$=80.8$\pm$1.1 and $K_{\rm 2}$=124.1$\pm$1.3,
respectively. Both values are in good agreement with those of $K_{\rm
1}$=82.93$\pm$0.03 and $K_{\rm 2}$=127.3$\pm$0.1 derived in this work from 8
hermes spectra.

65 UMa B: the star is discovered as magnetic by Bychkov et al. (2003) while
Aurière et al. (2007) measure its longitudinal magnetic field and rotation
period to be $B_{l}=-166\pm 20$ G and 15.830 d, respectively. Joshi et al.
(2010) present a detailed spectroscopic study of the star based on a time-
series of high-resolution spectra obtained with the 2.56-m Nordic Optical
Telescope (NOT). The authors report on the absence of RV variations in the
spectra and on the abundance pattern characteristic of evolved Ap stars.
Zasche et al. (2012, see their fig. 7) report 65 UMa to be a sextuple system
consisting of: (i) a close eclipsing pair of nearly identical stars (Aa1+Aa2)
orbiting each other with a period of $\sim$1.73 d; (ii) a distant third
component (Ab) orbiting the close pair with a period of about 640 d; (iii) a
fourth component (65 UMa B) resolved interferometrically whose period is about
118 yr; and (iv) fifth (65 UMa C) and sixth (65 UMa D) components with periods
of $\sim$14 kyr and $\sim$591 kyr, respectively. Our spectroscopic
classification of 65 UMa B as a single star not showing notable RV variability
is in full agreement with the literature. The TESS light curve reveals the
presence of shallow eclipses indicative of an eclipsing binary with a period
of about 1.73 d. We conclude that the eclipse signal in the TESS light curve
of 65 UMa B is due to a contaminating light from the component A which has a
similar surface brightness.

EL CVn: Maxted et al. (2014) study the system based on a combined WASP
photometry and ses high-resolution (R=60 000) spectroscopy. The authors find
the system to be composed of an A-type primary and a pre-He white dwarf
companion, residing in a circular orbit with the RV semi-amplitude of the
primary $K_{\rm 1}=29.0\pm 0.4$ km s-1. The latter is in agreement with our
determination of $K_{\rm 1}=30.0\pm 1.6$ km s-1. Wang et al. (2020) measure
the spectral lines of both binary components with the far-UV HST/COS
spectroscopy and report the detection of a spectral signature of the pre-He-WD
companion in the Mg ii 4481 Å optical line in the medium-resolution (R=31 500)
arces spectra. The authors report $K_{\rm 1}=29.0\pm 0.4$ km s-1 and $K_{\rm
2}=236.2\pm 1.1$ km s-1, in good agreement with the determination of the
respective parameters in this work.

$\alpha$ CrB is a well-studied object whose orbit is reported for the first
time by Jordan (1910). The system is solved spectroscopically by Ebbighausen
(1976) who reports $e=0.404\pm 0.004$ and $K_{\rm 1}=35.8\pm 0.2$ km s-1.
Tomkin & Popper (1986) report the detection of a secondary component based on
newly obtained spectroscopic data with the coudé spectrograph attached to the
2.7-m McDonald Observatory telescope. The authors report $e=0.371\pm 0.005$,
$K_{\rm 1}=35.4\pm 0.5$ km s-1, and $K_{\rm 2}=99.0\pm 0.5$ km s-1, in good
agreement with the respective parameters derived in this work from 14 hermes
spectra. Of the more recent studies, Schmitt et al. (2016) provide RV
measurements of the fainter secondary component using the method of cross-
correlation and combine those with their own and historical RV data of the
primary component to update the system’s orbital elements and to search for
evidence of an apsidal motion. The authors report $e=0.379\pm 0.002$, $K_{\rm
1}=36.2\pm 0.1$ km s-1, and $K_{\rm 2}=98.0\pm 0.3$ km s-1, in excellent
agreement with the respective orbital elements derived in this work. In
addition, Schmitt et al. (2016) find the apsidal motion period to be between
6600 and 10 600 years in agreement with the previous findings, and report on
the alignment of the orbit and rotation axes.

V994 Her: the system is discovered to have a quadruple-lined double-eclipsing
nature from ground-based multi-colour photometric and high-resolution
spectroscopic data by Lee et al. (2008). The authors report two eclipsing
binaries to have orbital periods of $\sim$2.08 d and $\sim$1.42 d, of which
the former is in agreement with the period used in this work and reported by
IJspeert et al. (2021). Both binaries are found to have slightly elliptic
($e$¡0.1) orbits with the reported masses of 2.83$\pm$0.20 M⊙, 2.30$\pm$0.16
M⊙, 1.87$\pm$0.12 M⊙, and 1.86$\pm$0.12 M⊙ for the Aa, Ab, Ba, and Bb
components, respectively. Zasche & Uhlař (2013) report V994 Her to be a
quintuple system with the two eclipsing binaries orbiting each other with the
period of about 6.3 d. These findings are refined in Zasche & Uhlař (2016)
where the authors report orbital period of $\sim$2.9 d for the two eclipsing
pairs and find evidence for apsidal motion with periods of about 116 and 111
years. The authors also report updated masses of 3.01$\pm$0.06 M⊙,
2.58$\pm$0.05 M⊙, 1.84$\pm$0.03 M⊙, and 1.93$\pm$0.04 M⊙ for the Aa, Ab, Ba,
and Bb components, respectively. Ultimately, Zasche et al. (2023) report the
discovery of a third set of eclipses in the TESS space-based and archival
ground-based photometric data which makes the system a triply-eclipsing
sextuple star system. The authors refine the orbital period of the
A(Aa+Ab)–B(Ba+Bb) core system to 1062$\pm$2 d and report about its close to
3:2 mean motion resonance. The eclipsing pair C(Ca+Cb) is reported to have an
elliptic orbit with the period of 1.96 d and component masses of
1.81${}^{+0.17}_{-0.07}$ M⊙ and 1.08${}^{+0.16}_{-0.11}$ M⊙, respectively. Our
classification of V994 Her as an SB4 system in this work is in agreement with
the previous findings.

V1898 Cyg: the system is classified as a spectroscopic double-lined binary by
Abt et al. (1972) and as an eclipsing system by Halbedel (1985). A detailed
analysis of the system based on the newly obtained spectroscopic and archival
photometric data is presented in Dervisoglu et al. (2011). The authors report
the $K_{1}$ and $K_{2}$ RV semi-amplitudes of 55.2$\pm$0.8 km s-1 and
287.6$\pm$2.1 km s-1, respectively, in good agreement with the findings in
this work.

GK Cep: the system is classified as an eclipsing spectroscopic double-lined
binary with the period of 0.936171 d and spectroscopic mass ratio of 0.92 by
Bartolini et al. (1965). Furthermore, Pribulla et al. (2009) classify the
system as a spectroscopic triple, in agreement with the findings in this work.
A detailed study of GK Cep is presented in Zhao et al. (2021) based on newly
obtained photometric data with the lunar-based ultraviolet telescope and TESS
space-based mission. The authors confirm the presence of a third body in the
system and measure the masses of the close binary components to be 1.93 M⊙ and
2.11 M⊙.

AH Cep: the system is analysed based on a combined photometric and
spectroscopic data set by Bell et al. (1986). The authors report $K_{1}$ =
249$\pm$8 km s-1 and $K_{2}$ = 283$\pm$8 km s-1 under the assumption of a
circular orbit. Holmgren et al. (1990a) and Burkholder et al. (1997) report
$K_{1}$ = 237$\pm$2 km s-1 and $K_{2}$ = 269$\pm$2 km s-1 and $K_{1}$ =
230.0$\pm$3.2 km s-1 and $K_{2}$ = 277.6$\pm$4.4 km s-1, respectively, under
the same assumption of $e=0$. The latter solution is in agreement with the
findings in this work. A recent study by Pavlovski et al. (2018) reports a
slightly higher value for $K_{1}$ = 234.9$\pm$1.1 km s-1 while $K_{2}$ =
276.9$\pm$1.4 km s-1 remains to be consistent with our determination of the
respective parameter.

WW Aur: a detailed analysis of the system is presented by Southworth et al.
(2005). The authors report $K_{1}$ = 116.81$\pm$0.23 km s-1 and $K_{2}$ =
126.49$\pm$0.28 km s-1 and $M_{1}$ = 1.964$\pm$0.007 M⊙ and $M_{2}$ =
1.814$\pm$0.007 M⊙ for the primary and secondary component, respectively. Our
determinations of the RV semi-amplitudes are in agreement for the primary
component and slightly lower for the secondary star.

AW Cam: a single-lined binary solution is presented by Mammano et al. (1967)
with $K_{1}$ = 112$\pm$5 km s-1 under the assumption of a circular orbit. Frey
et al. (2010) provide a simultaneous light- and RV-curve solution estimating
$K_{1}$ to 110 km s-1. Our estimate of $K_{1}$ is in agreement with the
previous findings and we establish the double-lined nature of the system in
this work.

V453 Cyg: Popper & Hill (1991), Simon & Sturm (1994), and Burkholder et al.
(1997) report spectroscopic RV semi-amplitudes of $K_{1}$ = 171$\pm$1.5 km
s-1, $K_{2}$ = 222$\pm$2.5 km s-1; $K_{1}$ = 171.7$\pm$2.9 km s-1, $K_{2}$ =
223.1$\pm$2.9 km s-1; $K_{1}$ = 173.2$\pm$1.3 km s-1, $K_{2}$ = 213.6$\pm$3.0
km s-1, respectively. Of the more recent studies, Southworth et al. (2004)
obtain eccentricity $e=0.022\pm 0.002$ from the apsidal motion solution and
$K_{1}$ = 173.7$\pm$0.8 km s-1, $K_{2}$ = 224.6$\pm$2.0 km s-1 from the RV
fitting. Consistent with these findings are the solutions obtained by
Pavlovski & Southworth (2009); Pavlovski et al. (2018), where the latter study
reports $e=0.022\pm 0.002$, $K_{1}$ = 175.2$\pm$0.7 km s-1, and $K_{2}$ =
220.2$\pm$1.6 km s-1. Orbital elements determined by us in this work are in
agreement with the previous studies, including small eccentricity of the
system. The star is classified as including a $\beta$ Cep variable by
Southworth et al. (2020), in line with our own classification in this work.

NY Cep: Holmgren et al. (1990b) present a detailed spectroscopic study of the
system based on 26 newly obtained high-resolution spectra. Among other things,
the authors report eccentricity, argument of periastron, and components’ RV
semi-amplitudes of $e=0.48\pm 0.02$, $\omega=58^{\circ}\pm 2$, $K_{1}$ =
112$\pm$3 km s-1, and $K_{2}$ = 158$\pm$8 km s-1, respectively. Albrecht et
al. (2011) revisit the system based on 46 newly obtained high-resolution
spectra with the sophie spectrograph. The authors report $e=0.443\pm 0.005$,
$\omega=56.3^{\circ}\pm 1$, $K_{1}$ = 113.8$\pm$1.2 km s-1, and $K_{2}$ =
139$\pm$4 km s-1 assuming $P=15.27566$ d as derived in Ahn (1992). Our
determinations of the eccentricity and argument of periastron of the system,
as well as the RV semi-amplitudes of both binary components are in good
agreement with the findings by Albrecht et al. (2011).

16 Lac ($=$ EN Lac): a detailed spectroscopic analysis of the star as an SB1
system based on some 1200 newly obtained and archival spectra is presented by
Lehmann et al. (2001). The authors report $P=12.096844$ d, $e=0.0392\pm
0.0017$, $\omega=63.7^{\circ}\pm 2.1$, and $K_{1}$ = 23.818$\pm$0.033 km s-1.
While the system has also been known as a single eclipsing since the discovery
by Jerzykiewicz (1980), a grazing secondary eclipse is for the first time
detected in the TESS data by Southworth & Bowman (2022). The spectroscopic
orbital elements derived by us in this work are in good agreement with those
presented by Lehmann et al. (2001). The star is reported as a $\beta$ Cep
variable by Jerzykiewicz (1980); Dziembowski & Jerzykiewicz (1996); Lehmann et
al. (2001); Aerts et al. (2003); Jerzykiewicz et al. (2015); Southworth &
Bowman (2022), in line with our own findings in this work.

V436 Per: Harmanec et al. (1997) present the first detailed spectroscopic
study of the system based on a collection of newly obtained and archival data.
Using the method of spectral disentangling, the authors report the detection
of LPVs and determine orbital eccentricity, argument of periastron, and RV
semi-amplitudes of the components to be $e=0.3882\pm 0.0043$,
$\omega=108.98^{\circ}\pm 0.27$, $K_{1}$ = 98.0$\pm$1.0 km s-1, and $K_{2}$ =
102.5$\pm$1.2 km s-1, respectively. The system is revisited by Janík et al.
(2003) based on a new set of high-resolution spectroscopic observations. The
authors do not confirm the previously reported LPVs and present an updated set
of spectroscopic orbital elements, $e=0.3768\pm 0.0014$,
$\omega=109.83^{\circ}\pm 0.10$, $K_{1}$ = 97.4$\pm$0.1 km s-1, and $K_{2}$ =
91.2$\pm$0.1 km s-1. A notable finding in the latter study is that the
secondary component has a lower RV semi-amplitude suggesting it is a more
massive star in the system. Southworth & Bowman (2022) analyse TESS
photometric data of the V346 Per system, report several local minima in the
obtained solution, and emphasise the importance of obtaining an independent
(spectroscopic) estimate of the components’ light ratio to resolve the faced
degeneracy. Spectroscopic orbital elements of the system derived by us in this
work are in agreement with those reported by Harmanec et al. (1997) except
that both $K$-values found by us are some 0.3% lower. We also note that unlike
Janík et al. (2003) we find the primary component to be the more massive one
in the system. In addition, we find the star to exhibit stochastic low-
frequency photometric variability, in a agreement with the classification by
Southworth & Bowman (2022).

RX Dra: we find the system to contain a $\gamma$ Dor-type pulsator, in
agreement with the classification by Southworth & Van Reeth (2022).

V1425 Cyg: a detailed study of the system based on a combined set of
spectroscopic and photometric observations is presented by Hill & Khalesseh
(1993). The authors report $K_{1}$ = 142.3$\pm$1.8 km s-1 and $K_{2}$ =
221.7$\pm$2.2 km s-1 under the assumption of a circular orbit. The respective
RV semi-amplitudes derived by us in this work are in agreement with the
findings by Hill & Khalesseh (1993).

V446 Cep: our photometric classification of the system as a $\beta$ Cep
pulsator is in agreement with the conclusions by Southworth & Bowman (2022)
who also report the star to presumably exhibit tidally-induced oscillations.

CM Lac: a study of the system based on a combined set of spectroscopic and
photometric observations is presented by Liakos & Niarchos (2012). The authors
determine $K_{1}$ = 119$\pm$2 km s-1, and $K_{2}$ = 156$\pm$2 km s-1 under the
assumption of a circular orbit. Southworth & Van Reeth (2022) revisit the
system based on TESS photometric data and spectroscopic RV measurements from
Liakos & Niarchos (2012). The authors report $K_{1}$ = 120.0$\pm$3.4 km s-1,
and $K_{2}$ = 157.0$\pm$3.3 km s-1, in agreement with the previous findings.
Our determinations of the RV semi-amplitudes are in good agreement with both
of the above-mentioned studies. Southworth & Van Reeth (2022) report the
system to contain a $\gamma$ Dor pulsator, we confirm their findings in this
work.

V398 Lac: the system is investigated spectroscopically by Çakırlı et al.
(2007). The authors report $e=0.230$, $K_{1}$ = 110.3$\pm$3.7 km s-1, and
$K_{2}$ = 128.5$\pm$3.8 km s-1 for the eccentricity and RV semi-amplitudes of
the components, respectively. A moderate orbital eccentricity of 0.273 and
$0.2284\pm 0.0007$ is confirmed by Bulut & Demircan (2008) and Wolf et al.
(2013), respectively, based on the studies of Hipparcos photometry and apsidal
motion of the system. While the orbital eccentricity obtained by us is in
agreement with the previous studies the RV semi-amplitudes of both components
are lower by some 15%-20% than those reported by Çakırlı et al. (2007).

V402 Lac: a detailed analysis of the system based on the newly obtained
spectroscopic data and TESS space-based photometry is presented by Baroch et
al. (2022). The authors determine eccentricity and RV semi-amplitudes to be
$e=0.376\pm 0.003$, $K_{1}$ = 128.5$\pm$0.8 km s-1, and $K_{2}$ =
129.2$\pm$0.8 km s-1 from the combined RV and eclipse timing analysis,
suggesting the system is composed of two stars of a similar mass. Our findings
are similar to but not exactly the same as the determinations by Baroch et al.
(2022). In particular, we find a slightly lower eccentricity and a $K$-value
for the primary component, the latter being suggestive of a slightly smaller
mass ratio of the system.

AE Pic: spectroscopic orbit of the system is presented by Sahade & Landi Dessy
(1950). The authors report $e=0.1$, $\omega=39^{\circ}$, and $K_{1}$ = 119 km
s-1. In this work, we classify the star as an SB2 system with a slightly lower
eccentricity and RV semi-amplitude of the primary than reported previously.
The system is classified as an eclipsing binary containing a pulsator that
additionally exhibits signatures of rotational modulation by Barraza et al.
(2022). We confirm their findings in this work and classify the star as to
having a $\beta$ Cep pulsator with extra signatures of rotational modulation
in the light curve.

##  6 Conclusions

Figure 5: Kiel diagram of stars listed in Table LABEL:Tab:SPDOrbit. In case of
SB2 systems, parameters of the hotter component are displayed. Solid lines
represent mesa evolutionary tracks for the stellar mass values $M$ of, from
right to left, 1.5, 3.0, 4.5, 7,0, and 13 M⊙, and a minimum amount of
overshoot $f_{\rm OV}=0.005$ $H_{\rm p}$. The tracks are from Johnston et al.
(2019b).

Of the total of 573 systems proposed for spectroscopic monitoring with the
hermes and feros spectrographs, in this work we report phase-resolved
spectroscopic observations for 83 of them (cf. Table LABEL:Tab:targets).
Sixty-five systems are classified by us as either SB1 or SB2 and their
respective orbital elements obtained with the method of spectral disentangling
are provided in Table LABEL:Tab:SPDOrbit. For two more systems, HD 246047 and
CD-47 4364, neither the method of spectral disentangling nor the LSD-based
classification could give a definitive answer as to whether these stars are
single and show low-amplitude LPVs or rather single-lined binaries with low
$K$-values. Furthermore, we identify in total four high-order SB3 or SB4
multiple systems, while the remaining 12 targets are classified by us as
single stars.

The $\log(T_{\rm eff})-\log(g)$ Kiel diagram of the above-mentioned sixty five
SB1 and SB2 systems is shown in Figure 5. Where possible, the $T_{\rm eff}$
and $\log\,g$ values have been taken from literature (cf. Section 5),
otherwise the respective values have been estimated from the spectral type and
luminosity class of the star by interpolating in the tables of Schmidt-Kaler
(1982). We also note that $T_{\rm eff}$ and $\log\,g$ of only the hotter and
more massive primary components are displayed in Fig. 5. Therefore, the
positions of the majority of stars in the Kiel diagram are not precise but
rather indicative and are used for the purpose of demonstration of the
parameter space covered by the sample presented in this work. With a slightly
over 10% sub-sample of the total 573 systems that we are currently monitoring
spectroscopically, we achieve the coverage of a large range in the effective
temperature and mass of the star, i.e. $T_{\rm eff}$ $\in[7000,30000]$ K and
$M\in[1.5,15]$ M⊙. Currently, the lower-mass regime with $M\lesssim 3.5$ M⊙ is
populated more densely and the sample consists almost exclusively of stars in
the core hydrogen burning phase.

Nevertheless, the present sample allows us to start looking into the problem
of the mass discrepancy as a function of the stellar mass, surface properties
of the star (i.e. $T_{\rm eff}$, $\log\,g$, and $v\,\sin\,i$), binary
configuration (i.e. detached versus semi-detached systems), and intrinsic
variability of the star (i.e. pulsating versus non-pulsating components).
Indeed, several systems in the sample with AF-type primary components are
known Algols and more of those have light curve characteristics of Algol-type
binaries. A mix of well-detached and semi-detached systems in the sample
allows us to assess quantitatively the role of binary interactions in the mass
discrepancy problem, as has been suggested by Mahy et al. (2020). Furthermore,
diversity of the sample in terms of the effective temperature and mass of the
primary component is an asset in the context of the proposed in the literature
possible connection between the mass discrepancy and amount of the near-core
mixing in the form of the convective core overshoot (e.g., Guinan et al. 2000;
Massey et al. 2012; Morrell et al. 2014; Claret & Torres 2016, 2017, 2018,
2019). The fact that the present sample contains binaries with and without
pulsating components will allow us to study the role of pulsations in the mass
discrepancy problem, methodologically (i.e. to what extent the presence of
intrinsic variability alters inference of the absolute stellar dimensions from
a light curve) and physics wise (e.g. to what extent a wave-induced mixing
alters interior and atmospheric properties of stars in models that are used to
quantify the mass discrepancy). Moreover, detection of stellar oscillations in
at least one of the binary components will enable an independent measurement
of the amount of the near-core mixing with asteroseismic methods. A further
extension of the sample towards inclusion of more evolved stars will allow us
to investigate a possible connection between the mass discrepancy and
evolutionary stage of the star, in particular the role of turbulent and
radiative pressure terms in the inference of the effective temperature of the
star when its surface gravity is accurately known from a photodynamical model
of the system (e.g., Tkachenko et al. 2020).

Summarising, the 15% sub-sample of the total of 573 eclipsing binary systems
that we are currently monitoring spectroscopically, represents one of the
largest (if not the largest) eclipsing binary star ensembles that are being
observed in a largely consistent way (i.e. with the same ground-based
instruments for spectroscopic data and TESS space-based mission for
photometric data sets) and that will be analysed by us with a well-established
modelling framework described in detail in Tkachenko et al. (2020). The high
level of consistency in the observational strategy and modelling approach is
in particular important to minimise otherwise numerous and hardly traceable
systematic uncertainties in the investigation of the mass discrepancy problem.
Coupled with a homogeneous and in the near future complete coverage of the
parameter space, we can conclude that our sample of eclipsing binaries has all
the potential to become the foundation to quantify and reesolve the mass
discrepancy problem in eclipsing binaries.

Finally, we stress an appreciable capability of the spectral disentangling
method to infer precise and accurate spectroscopic orbital elements and
individual components’ spectra from as little as 6-8 orbital phase-resolved
spectroscopic observations. Indeed, using an estimate of the orbital period of
the system as the only input to the planning of our observational campaign, we
demonstrate that a stable SPD solution can be obtained for the majority of
binaries in the sample. This is opposed to the established thinking in the
community that an extensive time-series of typically a couple of dozen
spectroscopic observations is required for the SPD method to deliver
meaningful results. In particular, by comparing our orbital solutions reported
in Table LABEL:Tab:SPDOrbit with those obtained by other research groups based
on typically more extended datasets, we find a good agreement for all but a
few systems (cf. Section 5) in the sample. Moreover, systems like AE Pic or AW
Cam that are known as SB1 in the literature are discovered by us to be double-
lined binaries and thus become ideal candidates for a detailed modelling with
the goal to infer absolute dimensions of both binary components. Furthermore,
the precision with which we infer the RV semi-amplitudes of both binary
components is often comparable to that reported in dedicated studies in the
literature (e.g., Pavlovski et al. 2018; Schmitt et al. 2016, for the AH Cep
and $\alpha$ CrB systems, respectively). This in turn translates into our
ability to infer absolute dimensions of stars with the precision and accuracy
better than 3% when a complementary high-quality ground- or space-based
photometric dataset is available (which is the case for all systems in our
sample).

The above-mentioned accuracy of 3% in stellar mass is sufficiently high to
provide a pertinent calibration of SSE models for scientific exploitation of
future space-based missions like PLATO (Rauer et al. 2014). Indeed, Chaplin et
al. (2014) and Silva Aguirre et al. (2017) demonstrate asteroseismically for
solar-like pulsators (i.e. main-sequence Sun-like stars and evolved
intermediate-mass stars on the red giant brunch) that $\sim$10% precision in
age can be achieved when the mass and radius of the star are measured with the
precision of 4% and 2%, respectively. The 10% age accuracy is one of the most
fundamental science requirements of the PLATO mission to be able to
characterise an Earth-like planet in the habitable zone of a Sun-like star
(Rauer et al. 2014). In this work, we demonstrate that our approach to the
planning of spectroscopic observations and subsequent analysis of the obtained
data precision and accuracy wise is compliant even with the most stringent
requirements of space-based missions such as PLATO.

In the forthcoming papers, we will present detailed analyses of individual
systems based on their combined spectroscopic and TESS photometric data. All
systems will be studied in the context of the mass discrepancy problem
presented in detail in Section 1. Also, the forthcoming papers will include
updates on our ongoing spectroscopic monitoring of the entire sample of 573
candidate eclipsing binaries, in a format similar to that used in the present
study.

###### Acknowledgements.

The research leading to these results has received funding from the KU Leuven
Research Council (grant C16/18/005: PARADISE), from the Research Foundation
Flanders (FWO) under grant agreements G089422N (NS, AT), V506223N (KP,
Scientific stay in Flanders), 11E5620N (SG, PhD Aspirant mandate), 1124321N
(LIJ, PhD Aspirant mandate), and 1286521N (DMB, Senior Postdoctoral
Fellowship), as well as from the BELgian federal Science Policy Office
(BELSPO) through PRODEX grant PLATO. DMB gratefully acknowledges the
Engineering and Physical Sciences Research Council (EPSRC) of UK Research and
Innovation (UKRI) in the form of a Frontier Research grant under the UK
government’s ERC Horizon Europe funding guarantee (grant number
[EP/Y031059/1]), and a University Research Fellowship from the Royal Society
(grant number: [URF\R1\231631]). Based on observations made with the Mercator
Telescope, operated on the island of La Palma by the Flemish Community, at the
Spanish Observatorio del Roque de los Muchachos of the Instituto de
Astrofísica de Canarias. Based on observations obtained with the HERMES
spectrograph, which is supported by the Research Foundation - Flanders (FWO),
Belgium, the Research Council of KU Leuven, Belgium, the Fonds National de la
Recherche Scientifique (F.R.S.-FNRS), Belgium, the Royal Observatory of
Belgium, the Observatoire de Genève, Switzerland and the Thüringer
Landessternwarte Tautenburg, Germany. Some of the observations used in this
work were obtained with the FEROS spectrograph mounted on the 2.2-m MPG/ESO
telescope at the ESO La Silla observatory under program 0106.A-0906(A). The
TESS data in this paper were obtained from the Mikulski Archive for Space
Telescopes (MAST) at the Space Telescope Science Institute (STScI), which is
operated by the Association of Universities for Research in Astronomy, Inc.,
under NASA contract NAS5-26555. Support to MAST for these data is provided by
the NASA Office of Space Science via grant NAG5-7584 and by other grants and
contracts. Funding for the TESS mission is provided by the NASA Explorer
Program. This research has made use of the SIMBAD database, operated at CDS,
Strasbourg, France; the SAO/NASA Astrophysics Data System; and the VizieR
catalog access tool, CDS, Strasbourg, France.

## References

  * Abt et al. (1972) Abt, H. A., Levy, S. G., & Gandet, T. L. 1972, AJ, 77, 138
  * Aerts et al. (2003) Aerts, C., Lehmann, H., Briquet, M., et al. 2003, A&A, 399, 639
  * Ahn (1992) Ahn, Y. S. 1992, Ap&SS, 198, 137
  * Albrecht et al. (2011) Albrecht, S., Winn, J. N., Carter, J. A., Snellen, I. A. G., & de Mooij, E. J. W. 2011, ApJ, 726, 68
  * Aurière et al. (2007) Aurière, M., Wade, G. A., Silvester, J., et al. 2007, A&A, 475, 1053
  * Ausseloos et al. (2006) Ausseloos, M., Aerts, C., Lefever, K., Davis, J., & Harmanec, P. 2006, A&A, 455, 259
  * Baroch et al. (2022) Baroch, D., Giménez, A., Morales, J. C., et al. 2022, A&A, 665, A13
  * Barraza et al. (2022) Barraza, L. F., Gomes, R. L., Messias, Y. S., et al. 2022, ApJ, 924, 117
  * Bartolini et al. (1965) Bartolini, C., Mammano, A., Mannino, G., & Margoni, R. 1965, Information Bulletin on Variable Stars, 85, 1
  * Bell et al. (1986) Bell, S. A., Hilditch, R. W., & Adamson, A. J. 1986, MNRAS, 223, 513
  * Bowman et al. (2019a) Bowman, D. M., Burssens, S., Pedersen, M. G., et al. 2019a, Nature Astronomy, 3, 760
  * Bowman et al. (2020) Bowman, D. M., Burssens, S., Simón-Díaz, S., et al. 2020, A&A, 640, A36
  * Bowman et al. (2019b) Bowman, D. M., Johnston, C., Tkachenko, A., et al. 2019b, ApJ, 883, L26
  * Bowman et al. (2022) Bowman, D. M., Vandenbussche, B., Sana, H., et al. 2022, A&A, 658, A96
  * Brahm et al. (2017) Brahm, R., Jordán, A., & Espinoza, N. 2017, PASP, 129, 034002
  * Brasseur et al. (2019) Brasseur, C. E., Phillip, C., Hargis, J., et al. 2019, in Astronomical Society of the Pacific Conference Series, Vol. 523, Astronomical Data Analysis Software and Systems XXVII, ed. P. J. Teuben, M. W. Pound, B. A. Thomas, & E. M. Warner, 397
  * Bulut & Demircan (2008) Bulut, İ. & Demircan, O. 2008, New A, 13, 252
  * Burkholder et al. (1997) Burkholder, V., Massey, P., & Morrell, N. 1997, ApJ, 490, 328
  * Bychkov et al. (2003) Bychkov, V. D., Bychkova, L. V., & Madej, J. 2003, A&A, 407, 631
  * Çakırlı et al. (2007) Çakırlı, Ö., Frasca, A., İbanoǧlu, C., Soydugan, F., & Deǧirmenci, O. 2007, Astronomische Nachrichten, 328, 536
  * Chaplin et al. (2014) Chaplin, W. J., Basu, S., Huber, D., et al. 2014, ApJS, 210, 1
  * Claret (2004) Claret, A. 2004, A&A, 424, 919
  * Claret (2012) Claret, A. 2012, A&A, 541, A113
  * Claret & Torres (2016) Claret, A. & Torres, G. 2016, A&A, 592, A15
  * Claret & Torres (2017) Claret, A. & Torres, G. 2017, ApJ, 849, 18
  * Claret & Torres (2018) Claret, A. & Torres, G. 2018, ApJ, 859, 100
  * Claret & Torres (2019) Claret, A. & Torres, G. 2019, ApJ, 876, 134
  * Costa et al. (2019) Costa, G., Girardi, L., Bressan, A., et al. 2019, MNRAS, 485, 4641
  * Dervisoglu et al. (2011) Dervisoglu, A., Çakirli, Ö., Ibanoglu, C., & Sipahi, E. 2011, Rev. Mexicana Astron. Astrofis., 47, 297
  * Donati et al. (1997) Donati, J. F., Semel, M., Carter, B. D., Rees, D. E., & Collier Cameron, A. 1997, MNRAS, 291, 658
  * Dziembowski & Jerzykiewicz (1996) Dziembowski, W. A. & Jerzykiewicz, M. 1996, A&A, 306, 436
  * Ebbighausen (1976) Ebbighausen, E. G. 1976, Publications of the Dominion Astrophysical Observatory Victoria, 14, 411
  * Frey et al. (2010) Frey, J. R., Angione, R. J., & Sievers, J. R. 2010, Information Bulletin on Variable Stars, 5946, 1
  * Freytag et al. (1996) Freytag, B., Ludwig, H. G., & Steffen, M. 1996, A&A, 313, 497
  * Gaia Collaboration et al. (2018) Gaia Collaboration, Brown, A. G. A., Vallenari, A., et al. 2018, A&A, 616, A1
  * Gebruers et al. (2022) Gebruers, S., Tkachenko, A., Bowman, D. M., et al. 2022, A&A, 665, A36
  * Guinan et al. (2000) Guinan, E. F., Ribas, I., Fitzpatrick, E. L., et al. 2000, ApJ, 544, 409
  * Hadrava (1995) Hadrava, P. 1995, A&AS, 114, 393
  * Halbedel (1985) Halbedel, E. M. 1985, Information Bulletin on Variable Stars, 2663, 1
  * Harmanec et al. (1997) Harmanec, P., Hadrava, P., Yang, S., et al. 1997, A&A, 319, 867
  * Herrero et al. (2000) Herrero, A., Puls, J., & Villamariz, M. R. 2000, A&A, 354, 193
  * Herwig (2000) Herwig, F. 2000, A&A, 360, 952
  * Hill & Khalesseh (1993) Hill, G. & Khalesseh, B. 1993, A&A, 276, 57
  * Holmgren et al. (1990a) Holmgren, D. E., Hill, G., & Fisher, W. 1990a, A&A, 236, 409
  * Holmgren et al. (1990b) Holmgren, D. E., Hill, G., Fisher, W., & Scarfe, C. D. 1990b, A&A, 231, 89
  * IJspeert et al. (2021) IJspeert, L. W., Tkachenko, A., Johnston, C., et al. 2021, A&A, 652, A120
  * Ilijic et al. (2004) Ilijic, S., Hensberge, H., Pavlovski, K., & Freyhammer, L. M. 2004, in Astronomical Society of the Pacific Conference Series, Vol. 318, Spectroscopically and Spatially Resolving the Components of the Close Binary Stars, ed. R. W. Hilditch, H. Hensberge, & K. Pavlovski, 111–113
  * Janík et al. (2003) Janík, J., Harmanec, P., Lehmann, H., et al. 2003, A&A, 408, 611
  * Jenkins et al. (2016) Jenkins, J. M., Twicken, J. D., McCauliff, S., et al. 2016, in Proc. SPIE, Vol. 9913, Software and Cyberinfrastructure for Astronomy IV, 99133E
  * Jerzykiewicz (1980) Jerzykiewicz, M. 1980, in Nonradial and Nonlinear Stellar Pulsation, ed. H. A. Hill & W. A. Dziembowski, Vol. 125, 125–134
  * Jerzykiewicz et al. (2015) Jerzykiewicz, M., Handler, G., Daszyńska-Daszkiewicz, J., et al. 2015, MNRAS, 454, 724
  * Johnston (2021) Johnston, C. 2021, A&A, 655, A29
  * Johnston et al. (2019a) Johnston, C., Pavlovski, K., & Tkachenko, A. 2019a, A&A, 628, A25
  * Johnston et al. (2019b) Johnston, C., Tkachenko, A., Aerts, C., et al. 2019b, MNRAS, 482, 1231
  * Johnston et al. (2023) Johnston, C., Tkachenko, A., Van Reeth, T., et al. 2023, A&A, 670, A167
  * Jordan (1910) Jordan, F. C. 1910, Publications of the Allegheny Observatory of the University of Pittsburgh, 1, 85
  * Joshi et al. (2010) Joshi, S., Ryabchikova, T., Kochukhov, O., et al. 2010, MNRAS, 401, 1299
  * Kaufer et al. (1999) Kaufer, A., Stahl, O., Tubbesing, S., et al. 1999, The Messenger, 95, 8
  * Kıran et al. (2016) Kıran, E., Harmanec, P., Değirmenci, Ö. L., et al. 2016, A&A, 587, A127
  * Kolbas et al. (2015) Kolbas, V., Pavlovski, K., Southworth, J., et al. 2015, MNRAS, 451, 4150
  * Kurtz (1985) Kurtz, D. W. 1985, MNRAS, 213, 773
  * Lee et al. (2008) Lee, C. U., Kim, S. L., Lee, J. W., et al. 2008, MNRAS, 389, 1630
  * Lehmann et al. (2001) Lehmann, H., Harmanec, P., Aerts, C., et al. 2001, A&A, 367, 236
  * Liakos & Niarchos (2012) Liakos, A. & Niarchos, P. 2012, Ap&SS, 340, 281
  * Lightkurve Collaboration et al. (2018) Lightkurve Collaboration, Cardoso, J. V. d. M., Hedges, C., et al. 2018, Lightkurve: Kepler and TESS time series analysis in Python
  * Lomb (1976) Lomb, N. R. 1976, Ap&SS, 39, 447
  * Mahy et al. (2020) Mahy, L., Almeida, L. A., Sana, H., et al. 2020, A&A, 634, A119
  * Mammano et al. (1967) Mammano, A., Margoni, R., & Stagni, R. 1967, Mem. Soc. Astron. Italiana, 38, 509
  * Martinet et al. (2021) Martinet, S., Meynet, G., Ekström, S., et al. 2021, A&A, 648, A126
  * Massey et al. (2012) Massey, P., Morrell, N. I., Neugent, K. F., et al. 2012, ApJ, 748, 96
  * Maxted et al. (2014) Maxted, P. F. L., Bloemen, S., Heber, U., et al. 2014, MNRAS, 437, 1681
  * Morrell et al. (2014) Morrell, N. I., Massey, P., Neugent, K. F., Penny, L. R., & Gies, D. R. 2014, ApJ, 789, 139
  * Özdarcan et al. (2012) Özdarcan, O., Sipahi, E., & Dal, H. A. 2012, New A, 17, 483
  * Pavlovski & Hensberge (2010) Pavlovski, K. & Hensberge, H. 2010, in Astronomical Society of the Pacific Conference Series, Vol. 435, Binaries - Key to Comprehension of the Universe, ed. A. Prša & M. Zejda, 207
  * Pavlovski et al. (2022) Pavlovski, K., Hummel, C. A., Tkachenko, A., et al. 2022, A&A, 658, A92
  * Pavlovski & Southworth (2009) Pavlovski, K. & Southworth, J. 2009, MNRAS, 394, 1519
  * Pavlovski et al. (2018) Pavlovski, K., Southworth, J., & Tamajo, E. 2018, MNRAS, 481, 3129
  * Pavlovski et al. (2023) Pavlovski, K., Southworth, J., Tkachenko, A., Van Reeth, T., & Tamajo, E. 2023, A&A, 671, A139
  * Pavlovski et al. (2009) Pavlovski, K., Tamajo, E., Koubský, P., et al. 2009, MNRAS, 400, 791
  * Paxton et al. (2011) Paxton, B., Bildsten, L., Dotter, A., et al. 2011, ApJS, 192, 3
  * Paxton et al. (2013) Paxton, B., Cantiello, M., Arras, P., et al. 2013, ApJS, 208, 4
  * Paxton et al. (2015) Paxton, B., Marchant, P., Schwab, J., et al. 2015, ApJS, 220, 15
  * Paxton et al. (2018) Paxton, B., Schwab, J., Bauer, E. B., et al. 2018, ApJS, 234, 34
  * Paxton et al. (2019) Paxton, B., Smolec, R., Schwab, J., et al. 2019, ApJS, 243, 10
  * Pedersen et al. (2019) Pedersen, M. G., Chowdhury, S., Johnston, C., et al. 2019, ApJ, 872, L9
  * Popper & Hill (1991) Popper, D. M. & Hill, G. 1991, AJ, 101, 600
  * Pribulla et al. (2009) Pribulla, T., Rucinski, S. M., DeBond, H., et al. 2009, AJ, 137, 3646
  * Raskin et al. (2011) Raskin, G., van Winckel, H., Hensberge, H., et al. 2011, A&A, 526, A69
  * Rauer et al. (2014) Rauer, H., Catala, C., Aerts, C., et al. 2014, Experimental Astronomy, 38, 249
  * Ricker et al. (2015) Ricker, G. R., Winn, J. N., Vanderspek, R., et al. 2015, Journal of Astronomical Telescopes, Instruments, and Systems, 1, 014003
  * Rosu et al. (2020) Rosu, S., Noels, A., Dupret, M. A., et al. 2020, A&A, 642, A221
  * Rosu et al. (2022a) Rosu, S., Rauw, G., Farnir, M., Dupret, M. A., & Noels, A. 2022a, A&A, 660, A120
  * Rosu et al. (2022b) Rosu, S., Rauw, G., Nazé, Y., Gosset, E., & Sterken, C. 2022b, A&A, 664, A98
  * Sahade & Landi Dessy (1950) Sahade, J. & Landi Dessy, J. 1950, ApJ, 111, 191
  * Sana et al. (2012) Sana, H., de Mink, S. E., de Koter, A., et al. 2012, Science, 337, 444
  * Sana et al. (2014) Sana, H., Le Bouquin, J. B., Lacour, S., et al. 2014, ApJS, 215, 15
  * Scargle (1982) Scargle, J. D. 1982, ApJ, 263, 835
  * Schmidt-Kaler (1982) Schmidt-Kaler, T. 1982, Bulletin d’Information du Centre de Donnees Stellaires, 23, 2
  * Schmitt et al. (2016) Schmitt, J. H. M. M., Schröder, K. P., Rauw, G., et al. 2016, A&A, 586, A104
  * Serenelli et al. (2021) Serenelli, A., Weiss, A., Aerts, C., et al. 2021, A&A Rev., 29, 4
  * Silva Aguirre et al. (2017) Silva Aguirre, V., Lund, M. N., Antia, H. M., et al. 2017, ApJ, 835, 173
  * Simon & Sturm (1994) Simon, K. P. & Sturm, E. 1994, A&A, 281, 286
  * Southworth & Bowman (2022) Southworth, J. & Bowman, D. M. 2022, MNRAS, 513, 3191
  * Southworth et al. (2021) Southworth, J., Bowman, D. M., & Pavlovski, K. 2021, MNRAS, 501, L65
  * Southworth et al. (2020) Southworth, J., Bowman, D. M., Tkachenko, A., & Pavlovski, K. 2020, MNRAS, 497, L19
  * Southworth et al. (2004) Southworth, J., Maxted, P. F. L., & Smalley, B. 2004, MNRAS, 351, 1277
  * Southworth et al. (2005) Southworth, J., Smalley, B., Maxted, P. F. L., Claret, A., & Etzel, P. B. 2005, MNRAS, 363, 529
  * Southworth & Van Reeth (2022) Southworth, J. & Van Reeth, T. 2022, MNRAS, 515, 2755
  * Stassun et al. (2019) Stassun, K. G., Oelkers, R. J., Paegert, M., et al. 2019, AJ, 158, 138
  * Stellingwerf (1978) Stellingwerf, R. F. 1978, ApJ, 224, 953
  * Themeßl et al. (2018) Themeßl, N., Hekker, S., Southworth, J., et al. 2018, MNRAS, 478, 4669
  * Tkachenko et al. (2014a) Tkachenko, A., Aerts, C., Pavlovski, K., et al. 2014a, MNRAS, 442, 616
  * Tkachenko et al. (2012) Tkachenko, A., Aerts, C., Pavlovski, K., et al. 2012, MNRAS, 424, L21
  * Tkachenko et al. (2014b) Tkachenko, A., Degroote, P., Aerts, C., et al. 2014b, MNRAS, 438, 3093
  * Tkachenko et al. (2016) Tkachenko, A., Matthews, J. M., Aerts, C., et al. 2016, MNRAS, 458, 1964
  * Tkachenko et al. (2020) Tkachenko, A., Pavlovski, K., Johnston, C., et al. 2020, A&A, 637, A60
  * Tkachenko et al. (2013) Tkachenko, A., Van Reeth, T., Tsymbal, V., et al. 2013, A&A, 560, A37
  * Tomkin & Popper (1986) Tomkin, J. & Popper, D. M. 1986, AJ, 91, 1428
  * Torres et al. (2010) Torres, G., Andersen, J., & Giménez, A. 2010, A&A Rev., 18, 67
  * Uytterhoeven et al. (2005a) Uytterhoeven, K., Briquet, M., Aerts, C., et al. 2005a, A&A, 432, 955
  * Uytterhoeven et al. (2005b) Uytterhoeven, K., Harmanec, P., Telting, J. H., & Aerts, C. 2005b, A&A, 440, 249
  * Viani & Basu (2020) Viani, L. S. & Basu, S. 2020, ApJ, 904, 22
  * Šubjak et al. (2020) Šubjak, J., Sharma, R., Carmichael, T. W., et al. 2020, AJ, 159, 151
  * Wang et al. (2020) Wang, L., Gies, D. R., Lester, K. V., et al. 2020, AJ, 159, 4
  * Wolf et al. (2013) Wolf, M., Zasche, P., Kučáková, H., et al. 2013, A&A, 549, A108
  * Zahn (1991) Zahn, J. P. 1991, A&A, 252, 179
  * Zasche et al. (2023) Zasche, P., Borkovits, T., Jayaraman, R., et al. 2023, MNRAS, 520, 3127
  * Zasche & Uhlař (2013) Zasche, P. & Uhlař, R. 2013, MNRAS, 429, 3472
  * Zasche & Uhlař (2016) Zasche, P. & Uhlař, R. 2016, A&A, 588, A121
  * Zasche et al. (2012) Zasche, P., Uhlář, R., Šlechta, M., et al. 2012, A&A, 542, A78
  * Zhao et al. (2021) Zhao, E. G., Qian, S. B., Zhou, X., et al. 2021, MNRAS, 504, 5155

##  Appendix A Classification of the sample stars

In this section, we present our spectroscopic and photometric classifications
for all stars included in the present sample.

“N” refers to the number of spectra we have at our disposal. Period
uncertainty is provided in parentheses in terms of the last digit.

Eclipse signal is detected in the TESS data but contamination as the source
cannot be excluded.

Variable eclipse depth such that eclipses practically disappear in the TESS
Sectors 61 & 62.

Table 1: Spectroscopic and photometric classification for stars studied in this work. Star ID | Period | N | Classification
---|---|---|---
Main | Alternative | TIC | (d) | Spec. | Phot.
hermes sample
BD+13 1880 | BD+13 1880 | TIC 186812530 | 3.677(2) | 10 | SB1 | EB
BD+30 3184 | BD+30 3184 | TIC 23470753 | 4.237216(5) | 8 | SB2 | EB, SPB or $\delta$ Sct
BD+36 3317 | BD+36 3317 | TIC 237195907 | 4.302147(3) | 8 | SB2 | EB, $\delta$ Sct
BD+47 1906 | BD+47 1906 | TIC 253049152 | 1.3993(4) | 12 | SB2 | EB
BD+67 1049 | BD+67 1049 | TIC 219110814 | 3.5888926(6) | 8 | SB1 | EB
65 UMa B1 | HD 103498 | TIC 141148944 | 0.8652(2) | 11 | single |  $\delta$ Sct
HD 107379 | HD 107379 | TIC 148895442 | 3.595266(3) | 9 | SB2 | EB
EL CVn | HD 116608 | TIC 165371937 | 0.795624(8) | 11 | SB2 | EB
NO Dra | HD 135437 | TIC 202442974 | 2.738920(2) | 9 | SB2 | EB
HD 1354661 | HD 135466 | TIC 202442982 | 29.26(2) | 8 | LPV |
HD 138305 | HD 138305 | TIC 368291074 | 3.498(3) | 8 | SB2 | EB, RotMod or ellipsoidal
$\alpha$ CrB | HD 139006 | TIC 274945059 | 17.360(5) | 14 | SB2 | EB
HD 13970 | HD 13970 | TIC 264614791 | 3.508808(2) | 9 | SB2 | EB, $\beta$ Cep
HD 150781A | HD 150781a | TIC 349444267 | 7.457(5) | 8 | SB2 | EB, $\gamma$ Dor
V920 Her | HD 151972 | TIC 143009538 | 6.926(3) | 8 | SB2 | EB
HD 158148 | HD 158148 | TIC 351701483 | 17.91032(2) | 9 | SB2 | EB, SPB
HD 160363 | HD 160363 | TIC 311433319 | 1.837710(2) | 11 | SB1 | EB
V994 Her | HD 170314 | TIC 424508303 | 2.083296(1) | 9 | SB4 | EB, high-order multiple
HD 172133 | HD 172133 | TIC 8705972 | 3.598199(4) | 8 | SB1 | EB, RotMod or ellipsoidal
BH Dra | HD 178001 | TIC 377192659 | 1.8172380(2) | 9 | SB2 | EB
V2108 Cyg | HD 191530 | TIC 378395625 | 2.560339(1) | 7 | SB2 | EB
V1898 Cyg | HD 200776 | TIC 273173532 | 1.5131193(2) | 42 | SB2 | EB
GK Cep | HD 205372 | TIC 256352113 | 0.9361715(2) | 7 | SB3 | EB
V383 Cep | HD 208106 | TIC 410522328 | 1.4957437(3) | 11 | SB2 | EB
HD 208510 | HD 208510 | TIC 299494754 | 1.5970321(2) | 8 | SB2 | EB, RotMod or ellipsoidal
AH Cep | HD 216014 | TIC 377506471 | 1.7747483(4) | 11 | SB2 | EB
V350 And | HD 2189 | TIC 58107375 | 1.7111454(2) | 8 | SB2 | EB, $\gamma$ Dor/$\delta$ Sct hybrid
HD 234650 | HD 234650 | TIC 21189379 | 7.590363(6) | 8 | SB3 | EB
u Tau | HD 23466 | TIC 426588729 | 2.4245(9) | 14 | SB2 | EB, RotMod or SPB
HD 234713 | HD 234713 | TIC 48087401 | 3.067219(3) | 8 | SB2 | EB
HD 237866 | HD 237866 | TIC 137905382 | 1.5238(4) | 9 | SB1 | EB
HD 2460471 | HD 246047 | TIC 116331699 | 9.91493(2) | 11 | SB1 or LPV |
IM Aur | HD 33853 | TIC 368180294 | 1.2472686(2) | 11 | SB2 | EB
HD 343821 | HD 34382 | TIC 2234723 | 2.4619009(4) | 13 | LPV | RotMod
HD 348568 | HD 348568 | TIC 342480862 | 11.727766(4) | 8 | SB2 | EB
HD 348725 | HD 348725 | TIC 342794723 | 1.531346(1) | 7 | SB2 | EB
HD 350685 | HD 350685 | TIC 392053854 | 2.873796(2) | 8 | SB1 | EB, $\beta$ Cep
HD 37646 | HD 37646 | TIC 75507062 | 0.6726158(4) | 8 | SB1 | EB
WW Aur | HD 46052 | TIC 172171873 | 2.525018(1) | 9 | SB2 | EB
V459 Aur | HD 46552 | TIC 172421004 | 1.0626466(7) | 14 | SB2 | EB
HD 47046 | HD 47046 | TIC 353759550 | 1.7166(6) | 9 | SB2 | EB, RotMod and ellipsoidal
AW Cam | HD 48049 | TIC 456263109 | 0.77134620(7) | 8 | SB2 | EB
HD 54159 | HD 54159 | TIC 367567347 | 0.80858(3) | 9 | single | RotMod
HD 57158 | HD 57158 | TIC 302907601 | 1.6982479(7) | 8 | SB4 | EB, high-order multiple?
HD 63887 | HD 63887 | TIC 457101125 | 3.791285(3) | 8 | SB2 | EB
V766 Cas | HD 8027 | TIC 241017747 | 2.3296557(4) | 9 | SB2 | EB
HD 89601 | HD 89601 | TIC 150251466 | 12.388(7) | 9 | SB1 | EB
TYC 3529-2494-11 | TYC 3529-2494-1 | TIC 7694914 | 0.4450381(2) | 9 | LPV | RotMod
V453 Cyg | HD 227696 | TIC 90349611 | 3.890(1) | 28 | SB2 | EB, $\beta$ Cep
NY Cep | HD 217312 | TIC 13389059 | 15.2884(2) | 8 | SB2 | EB
16 Lac | HD 216916 | TIC 129538133 | 12.097(1) | 180 | SB1 | EB, $\beta$ Cep
V436 Per | HD 11241 | TIC 403625251 | 25.936(1) | 41 | SB2 | EB, SLF
RX Dra | RX Dra | TIC 377190161 | 3.876(1) | 26 | SB2 | EB, $\gamma$ Dor
V1425 Cyg | HD 202000 | TIC 344456744 | 1.2523886(1) | 44 | SB2 | EB
V446 Cep | HD 210478 | TIC 335265326 | 3.808(1) | 73 | SB2 | EB, $\beta$ Cep
CM Lac | HD 209147 | TIC 331221558 | 1.6047(1) | 27 | SB2 | EB, $\gamma$ Dor
V398 Lac | HD 210180 | TIC 326374705 | 5.406081(1) | 26 | SB2 | EB
V402 Lac | HD 210405 | TIC 468792279 | 3.7818043(8) | 30 | SB2 | EB, SPB
feros sample
HD 1007371 | HD100737 | TIC 290391601 | 2.552548(2) | 9 | SB1 |  $\beta$ Cep?
HD 104233 | HD104233 | TIC 307687961 | 1.823975(1) | 10 | SB2 | EB
HD 121776 | HD121776 | TIC 448375181 | 1.737896(1) | 9 | SB2 | EB
HD 28913 | HD28913 | TIC 170729895 | 1.4906780(1) | 7 | LPV | RotMod
HD 300344 | HD300344 | TIC 274687574 | 2.7902916(3) | 7 | SB2 | EB, RotMod
HD 304241 | HD304241 | TIC 451259413 | 2.730422(3) | 9 | SB1 | EB
HD 309317 | HD309317 | TIC 306139502 | 2.257301(2) | 9 | SB2 | EB, $\beta$ Cep
AE Pic | HD46792 | TIC 150442264 | 2.9816923(9) | 8 | SB2 | EB, RotMod and $\beta$ Cep
HD 519811 | HD51981 | TIC 147314529 | 0.4627711(2) | 7 | single |
HD 52349 | HD52349 | TIC 80041531 | 2.775849(3) | 8 | SB2 | EB, SPB
V386 Pup | HD62738 | TIC 175254818 | 1.649314(1) | 6 | SB2 | EB
HD 66235 | HD66235 | TIC 285413162 | 1.5941375(1) | 7 | LPV | RotMod
HD 67025 | HD67025 | TIC 79935432 | 1.2823665(1) | 7 | SB2 | EB
HD 683401 | HD68340 | TIC 145405941 | 1.3069152(1) | 7 | single |
HD 75872 | HD75872 | TIC 29216374 | 0.94508671(6) | 6 | SB1 | EB
HD 793651 | HD79365 | TIC 74715631 | 0.90672432(6) | 8 | LPV |
HD 82110 | HD82110 | TIC 438089724 | 1.8810195(2) | 8 | SB2 | EB, SPB
HD 84493 | HD84493 | TIC 363146191 | 6.8760(2) | 8 | SB1 | EB, SLF
HD 91141 | HD91141 | TIC 457540424 | 2.3821886(4) | 10 | SB2 | EB
HD 91154 | HD91154 | TIC 457545293 | 3.6631544(5) | 10 | SB1 | EB
HD 927411 | HD92741 | TIC 458561474 | 5.372886(4) | 10 | LPV |  $\beta$ Cep?
HD 979661 | HD97966 | TIC 450276745 | 1.2666420(8) | 9 | single |
CD-45 43932 | TYC8151-937-1 | TIC 141858108 | 2.0494343(2) | 6 | SB1 | EB
CD-47 43641 | TYC8155-1212-1 | TIC 270844716 | 3.6580419(5) | 8 | SB1 or LPV | SPB
CD-56 1160 | TYC8514-106-1 | TIC 382044531 | 0.7903919(1) | 6 | SB2 | EB
|  |  |  |  |  |
Table 1: continued.

##  Appendix B SPD-based orbital solutions

In this section we present SPD-based orbital solutions for all stars that are
classified as either SB1 or SB2 systems in Table LABEL:Tab:targets.

$e$, $\omega$, and $K_{\rm i}$ stand for the eccentricity, time of periastron
passage, and RV semi-amplitudes, respectively. Cases where eccentricity and
argument of periastron were fixed to 0.0 and 90∘, respectively, are indicated
with the superscript “f”. Orbital period values and their uncertainties are
adopted from Table LABEL:Tab:targets.

The star is classified as an SB2 system based on the visual inspection of the
LSD profiles and/or original observed spectra, however the SPD solution could
be obtained for the brighter primary component only.

Same as above but no stable SPD solution could be obtained for either of the
components.

Either an SB1 system with a low $K$ semi-amplitude of about 1.5 km s-1 or a
single star showing LPVs: none of the methods employed in this work favours
any of those hypothesis.

Table 2: Spectroscopic orbital elements for all SB1 and SB2 systems studied in this work. Star ID | Period | $e$ | $\omega$ | $K_{\rm 1}$ | $K_{\rm 2}$
---|---|---|---|---|---
|  |  |  | value | error | value | error | value | error | value | error
Main | Alt. | TIC | (d) |  |  | (degrees) | (km s-1)
hermes sample
BD+13 1880 | BD+13 1880 | TIC 186812530 | 3.677(2) | 0.0f | 0.0f | 90f | 0.0f | 5.0 | 0.1 | — | —
BD+30 3184 | BD+30 3184 | TIC 23470753 | 4.237216(5) | 0.20 | 0.01 | 121.9 | 5.1 | 82.7 | 1.2 | 116.8 | 3.2
BD+36 3317 | BD+36 3317 | TIC 237195907 | 4.302147(3) | 0.0f | 0.0f | 90f | 0.0f | 82.93 | 0.03 | 127.3 | 0.1
BD+47 1906 | BD+47 1906 | TIC 253049152 | 1.3993(4) | 0.0f | 0.0f | 90f | 0.0f | 20.1 | 0.5 | 213.6 | 3.2
BD+67 1049 | BD+67 1049 | TIC 219110814 | 3.5888926(6) | 0.0f | 0.0f | 90f | 0.0f | 18.64 | 0.06 | — | —
HD 107379 | HD 107379 | TIC 148895442 | 3.595266(3) | 0.0f | 0.0f | 90f | 0.0f | 14.9 | 0.8 | 76.4 | 0.6
EL CVn | HD 116608 | TIC 165371937 | 0.795624(8) | 0.0f | 0.0f | 90f | 0.0f | 30.0 | 1.6 | 242.4 | 2.7
NO Dra | HD 135437 | TIC 202442974 | 2.738920(2) | 0.02 | 0.04 | 217 | 11 | 77.8 | 0.5 | 154.9 | 3.0
HD 138305 | HD 138305 | TIC 368291074 | 3.498(3) | 0.0f | 0.0f | 90f | 0.0f | 106.4 | 0.2 | 109.39 | 0.08
$\alpha$ CrB | HD 139006 | TIC 274945059 | 17.360(5) | 0.3795 | 0.0002 | 314.7 | 0.8 | 37.1 | 1.3 | 98.5 | 0.8
HD 13970 | HD 13970 | TIC 264614791 | 3.508808(2) | 0.0f | 0.0f | 90f | 0.0f | 117.1 | 1.4 | 188.2 | 4.3
HD 150781A | HD 150781A | TIC 349444267 | 7.457(5) | 0.139 | 0.001 | 189.0 | 0.6 | 84.8 | 0.6 | 85.7 | 0.1
V920 Her | HD 151972 | TIC 143009538 | 6.926(3) | 0.0f | 0.0f | 90f | 0.0f | 85.4 | 0.4 | 90.98 | 0.09
HD 158148 | HD 158148 | TIC 351701483 | 17.91032(2) | 0.0f | 0.0f | 90f | 0.0f | 18.0 | 1.1 | 129.0 | 0.4
HD 160363 | HD 160363 | TIC 311433319 | 1.837710(2) | 0.0f | 0.0f | 90f | 0.0f | 28.2 | 0.8 | — | —
HD 172133 | HD 172133 | TIC 8705972 | 3.598199(4) | 0.0f | 0.0f | 90f | 0.0f | 43.97 | 0.04 | — | —
BH Dra | HD 178001 | TIC 377192659 | 1.8172380(2) | 0.0f | 0.0f | 90f | 0.0f | 105.8 | 0.6 | 162.8 | 2.0
V2108 Cyg | HD 191530 | TIC 378395625 | 2.560339(1) | 0.0f | 0.0f | 90f | 0.0f | 44.7 | 3.2 | 72.6 | 1.4
V1898 Cyg | HD 200776 | TIC 273173532 | 1.5131193(2) | 0.0f | 0.0f | 90f | 0.0f | 54.6 | 1.5 | 293.0 | 1.4
V383 Cep | HD 208106 | TIC 410522328 | 1.4957437(3) | 0.0f | 0.0f | 90f | 0.0f | 122.9 | 0.7 | 151.8 | 4.2
HD 208510 | HD 208510 | TIC 299494754 | 1.5970321(2) | 0.0f | 0.0f | 90f | 0.0f | 68.0 | 6.1 | 243.7 | 5.6
AH Cep | HD 216014 | TIC 377506471 | 1.7747483(4) | 0.0f | 0.0f | 90f | 0.0f | 230.1 | 2.0 | 277.1 | 3.9
V350 And | HD 2189 | TIC 58107375 | 1.7111454(2) | 0.0f | 0.0f | 90f | 0.0f | 125.5 | 1.0 | 133.7 | 1.3
u Tau | HD 23466 | TIC 426588729 | 2.4245(9) | 0.22 | 0.01 | 230 | 8 | 22.5 | 2.7 | 75.2 | 2.8
HD 234713 | HD 234713 | TIC 48087401 | 3.067219(3) | 0.0f | 0.0f | 90f | 0.0f | 101.87 | 0.03 | 102.9 | 0.6
HD 237866 | HD 237866 | TIC 137905382 | 1.5238(4) | 0.0f | 0.0f | 90f | 0.0f | 19.5 | 1.8 | — | —
HD 2460473 | HD 246047 | TIC 116331699 | 9.91493(2) | 0.0f | 0.0f | 90f | 0.0f | 1.3 | 0.6 | — | —
IM Aur | HD 33853 | TIC 368180294 | 1.2472686(2) | 0.0f | 0.0f | 90f | 0.0f | 67.8 | 0.6 | 285.8 | 1.0
HD 348568 | HD 348568 | TIC 342480862 | 11.727766(4) | 0.0f | 0.0f | 90f | 0.0f | 23.8 | 2.0 | 131.3 | 3.8
HD 348725 | HD 348725 | TIC 342794723 | 1.531346(1) | 0.0f | 0.0f | 90f | 0.0f | 115.4 | 2.5 | 143.9 | 3.1
HD 350685 | HD 350685 | TIC 392053854 | 2.873796(2) | 0.0f | 0.0f | 90f | 0.0f | 34.1 | 2.9 | — | —
HD 37646 | HD 37646 | TIC 75507062 | 0.6726158(4) | 0.0f | 0.0f | 90f | 0.0f | 28.7 | 0.1 | — | —
WW Aur | HD 46052 | TIC 172171873 | 2.525018(1) | 0.0f | 0.0f | 90f | 0.0f | 117.53 | 0.02 | 124.59 | 0.04
V459 Aur | HD 46552 | TIC 172421004 | 1.0626466(7) | 0.0f | 0.0f | 90f | 0.0f | 65.2 | 1.6 | 149.8 | 2.5
HD 47046 | HD 47046 | TIC 353759550 | 1.7166(6) | 0.0f | 0.0f | 90f | 0.0f | 14.0 | 0.7 | 134.5 | 1.3
AW Cam | HD 48049 | TIC 456263109 | 0.77134620(7) | 0.0f | 0.0f | 90f | 0.0f | 113.5 | 2.7 | 252.9 | 3.3
HD 63887 | HD 63887 | TIC 457101125 | 3.791285(3) | 0.0f | 0.0f | 90f | 0.0f | 104.1 | 0.4 | 119.2 | 0.9
V766 Cas | HD 8027 | TIC 241017747 | 2.3296557(4) | 0.09 | 0.02 | 232 | 17 | 129.9 | 1.0 | 154.4 | 8.5
HD 89601 | HD 89601 | TIC 150251466 | 12.388(7) | 0.4207 | 0.0003 | 191.28 | 0.08 | 30.60 | 0.01 | — | —
V453 Cyg | HD 227696 | TIC 90349611 | 3.890(1) | 0.027 | 0.002 | 151.3 | 11.3 | 173.7 | 1.0 | 221.4 | 1.0
NY Cep | HD 217312 | TIC 13389059 | 15.2884(2) | 0.444 | 0.001 | 54.2 | 1.7 | 110.6 | 1.6 | 141.7 | 7.0
16 Lac | HD 216916 | TIC 129538133 | 12.097(1) | 0.047 | 0.004 | 40.5 | 1.6 | 23.75 | 0.06 | — | —
V436 Per | HD 11241 | TIC 403625251 | 25.936(1) | 0.373 | 0.009 | 109.5 | 1.0 | 94.1 | 1.8 | 99.9 | 1.0
RX Dra | RX Dra | TIC 377190161 | 3.876(1) | 0.0f | 0.0f | 90f | 0.0f | 85.8 | 0.2 | 104.6 | 0.3
V1425 Cyg | HD 202000 | TIC 344456744 | 1.2523886(1) | 0.0f | 0.0f | 90f | 0.0f | 143.8 | 1.3 | 218.4 | 2.2
V446 Cep | HD 210478 | TIC 335265326 | 3.808(1) | 0.0150 | 0.0005 | 69.7 | 17.4 | 42.3 | 1.9 | 309.7 | 2.4
CM Lac | HD 209147 | TIC 331221558 | 1.6047(1) | 0.0f | 0.0f | 90f | 0.0f | 121.0 | 0.2 | 155.3 | 0.8
V398 Lac | HD 210180 | TIC 326374705 | 5.406081(1) | 0.23 | 0.01 | 233.6 | 7.5 | 90.7 | 1.5 | 139.3 | 2.6
V402 Lac | HD 210405 | TIC 468792279 | 3.7818043(8) | 0.358 | 0.005 | 56.7 | 2.9 | 125.5 | 0.3 | 128.5 | 0.5
feros sample
HD 100737 | HD100737 | TIC 290391601 | 2.552548(2) | 0.0f | 0.0f | 90f | 0.0f | 56.9 | 0.5 | — | —
HD 104233 | HD104233 | TIC 307687961 | 1.823975(1) | 0.0f | 0.0f | 90f | 0.0f | 42.0 | 0.9 | 92.3 | 4.2
HD 1217761 | HD121776 | TIC 448375181 | 1.737896(1) | 0.0f | 0.0f | 90f | 0.0f | 81.8 | 1.1 | — | —
HD 3003442 | HD300344 | TIC 274687574 | 2.7902916(3) | — | — | — | — | — | — | — | —
HD 304241 | HD304241 | TIC 451259413 | 2.730422(3) | 0.0f | 0.0f | 90f | 0.0f | 84.6 | 1.3 | — | —
HD 309317 | HD309317 | TIC 306139502 | 2.257301(2) | 0.0f | 0.0f | 90f | 0.0f | 34.5 | 0.3 | 78.7 | 1.4
AE Pic | HD46792 | TIC 150442264 | 2.9816923(9) | 0.075 | 0.005 | 58.4 | 26.4 | 114.1 | 1.3 | 158.9 | 3.6
HD 523491 | HD52349 | TIC 80041531 | 2.775849(3) | 0.37 | 0.01 | 312.6 | 7.3 | 54.8 | 0.6 | — | —
V386 Pup2 | HD62738 | TIC 175254818 | 1.649314(1) | — | — | — | — | — | — | — | —
HD 67025 | HD67025 | TIC 79935432 | 1.2823665(1) | 0.0f | 0.0f | 90f | 0.0f | 147.4 | 1.9 | 173.7 | 2.7
HD 75872 | HD75872 | TIC 29216374 | 0.94508671(6) | 0.0f | 0.0f | 90f | 0.0f | 45.9 | 1.0 | — | —
HD 82110 | HD82110 | TIC 438089724 | 1.8810195(2) | 0.0f | 0.0f | 90f | 0.0f | 78.2 | 1.0 | 178.8 | 1.5
HD 84493 | HD84493 | TIC 363146191 | 6.8760(2) | 0.0f | 0.0f | 90f | 0.0f | 29.20 | 0.06 | — | —
HD 91141 | HD91141 | TIC 457540424 | 2.3821886(4) | 0.0f | 0.0f | 90f | 0.0f | 15.58 | 0.08 | 72.7 | 1.8
HD 91154 | HD91154 | TIC 457545293 | 3.6631544(5) | 0.0f | 0.0f | 90f | 0.0f | 36.5 | 1.5 | — | —
CD-45 4393 | TYC8151-937-1 | TIC 141858108 | 2.0494343(2) | 0.0f | 0.0f | 90f | 0.0f | 16.9 | 0.7 | — | —
CD-47 43643 | TYC8155-1212-1 | TIC 270844716 | 3.6580419(5) | 0.0f | 0.0f | 90f | 0.0f | 1.58 | 0.09 | — | —
CD-56 1160 | TYC8514-106-1 | TIC 382044531 | 0.7903919(1) | 0.0f | 0.0f | 90f | 0.0f | 65.7 | 1.5 | 155.6 | 2.2
|  |  |  |  |  |  |  |  |  |  |
Table 2: continued.