11institutetext: Instituto de Ciencias Astronómicas, de la Tierra y del
Espacio (ICATE-CONICET), C.C 467, 5400, San Juan, Argentina. 22institutetext:
Universidad Nacional de San Juan (UNSJ), Facultad de Ciencias Exactas, Físicas
y Naturales (FCEFN), San Juan, Argentina. 33institutetext: Instituto de
Investigación Multidisciplinar en Ciencia y Tecnología, Universidad de La
Serena, Raúl Bitrán 1305, La Serena, Chile 44institutetext: Departamento de
Física y Astronomía, Universidad de La Serena, Av. Cisternas 1200 N, La
Serena, Chile. 55institutetext: Gemini Observatory / NSF’s NOIRLab, Casilla
603, La Serena, Chile 66institutetext: Observatorio Astronómico de Córdoba
(OAC), Laprida 854, X5000BGR, Córdoba, Argentina. 77institutetext: Laboratório
Nacional de Astrofísica (LNA/MCTI), rua Estados Unidos 154, Itajubá, MG,
Brasil 88institutetext: Consejo Nacional de Investigaciones Científicas y
Técnicas (CONICET), Argentina

# Are Am stars and hot-Jupiter planets related?

C. Saffe 112288 J. Alacoria 1188 P. Miquelarena 112288 R. Petrucci 6688 M.
Jaque Arancibia 3344 R. Angeloni 55 E. Martioli 77 M. Flores 112288 E. Jofré
6688 A. Collado and F. Gunella 1122881188

(Received xx, 2022; accepted xx, 2022)

###### Abstract

Context. Metallic-lined Am stars are often components of short-period binary
systems, where tidal interactions would result in low rotational velocities
and help to develop the chemical peculiarities observed. However, the origin
of single Am stars and Am stars that belong to wide binary systems is unclear.

Aims. There is very recent evidence of an Am star hosting a hot-brown dwarf
likely synchronized and other possible Am stars hosting hot-Jupiter planets.
Following literature suggestions, we wonder if these hot-low mass companions
could play a role in the development of an Am star, that is to say, if they
could help to mitigate the ”single Am” problem.

Methods. We carried out a detailed abundance determination via spectral
synthesis of 19 early-type stars hosting hot-brown dwarfs and hot-Jupiter
planets, in order to determine the possible presence of Am stars in this
sample. The abundances were determined iteratively for 25 different species by
fitting synthetic spectra using the SYNTHE program together with local
thermodynamic equilibrium (LTE) ATLAS12 model atmospheres. The abundances of C
I, O I and Mg I were corrected by non-LTE effects. The complete chemical
patterns of the stars were then compared to those of Am stars and other
chemically peculiar stars.

Results. We studied a sample of 19 early-type stars, 7 of them hosting hot-
brown dwarfs and 12 of them hosting hot-Jupiter planets. We detected 4 Am
stars in our sample (KELT-19A, KELT-17, HATS-70 and TOI-503) and 2 possible Am
stars (TOI-681 and HAT-P-69). In particular, we detected the new Am star
HATS-70 which hosts a hot-brown dwarf, and rule out this class for the hot-
Jupiter host WASP-189, both showing different composition than previously
reported. For the first time, we estimated the incidence of Am stars within
stars hosting hot-brown dwarfs (50-75%) and within stars hosting hot-Jupiters
(20-42%). In particular, the incidence of Am stars hosting hot-brown dwarfs
resulted higher than the frequency of Am stars in general. This would imply
that the presence of hot-brown dwarfs could play a role in the development of
Am stars and possibly help to mitigate the ”single Am” problem, different to
the case of hot-Jupiter planets. Notably, these results would also indicate
that the search for hot-brown dwarfs may be benefited by targeting single Am
stars or Am stars in wide binary systems. We encourage the analysis off
additional early-type stars hosting hot-companions in order to improve the
significance of the initial trends found here.

###### Key Words.:

stars: chemically peculiar – stars: early type – stars: abundances – planetary
systems

##  1 Introduction

The metallic-lined A stars (Am stars) were first identified by Titus & Morgan
(1940) as objects with strong metal lines and a weak Ca II K line in their
spectra. On average, Am stars present over-abundances of most heavy elements
in their spectra, particularly Fe and Ni, together with under-abundances of Ca
and Sc (Preston, 1974; Fossati et al., 2007, 2009; Catanzaro et al., 2019).
Starting with the work of Conti (1970), other authors prefer a wider
definition for the group, including stars with underabundances of Ca (and/or
Sc) and/or overabundance of the Fe group elements (e.g. Lane & Lester, 1987;
Adelman, 1994; Adelman et al., 1997; Yüce & Adelman, 2014). Am stars rotate
more slowly than average A-type stars, as first noted by Slettebak (1954) and
then confirmed by a number of works (e.g. Slettebak, 1955; Abt, 2000;
Niemczura et al., 2015). The origin of the peculiar abundances is commonly
attributed to diffusion processes due to gravitational settling and radiative
levitation (Michaud, 1970; Michaud et al., 1976, 1983; Vauclair et al., 1978;
Alecian et al., 1996; Richer et al., 2000; Fossati et al., 2007), where stable
atmospheres of slowly-rotating A-type stars would allow the diffusion
processes to operate111Chemical peculiarities could also reach early F stars,
the Fm stars.. Initial observations showed that Am stars do not pulsate
(Breger, 1970; Breger et al., 1972), while further works showed that indeed
the pulsations are possible in a number of Am stars (e.g. Bessell & Eggen,
1972; Kurtz, 1989; Smalley et al., 2011, 2017). This showed that pulsations do
not necessarily conflict with the peculiar abundances observed in the Am
stars.

A number of works show that Am stars are often components of binary systems
(Abt, 1961; Abt & Levy, 1985; North et al., 1998; Debernardi et al., 2000;
Carquillat & Prieur et al., 2007; Stateva et al., 2012). In fact, the slow
rotation of Am stars is usually attributed to a tidal braking in short-period
(1 d $<$ P $<$ 10 d) binary systems (Abt, 1967; Budaj, 1996, 1997), rather
than to the presence of weak or ultra-weak magnetic fields (e.g., Folsom et
al., 2013; Blazère et al., 2016, 2020). Nonetheless, there are Am stars in
wide binary systems with longer periods and Am stars that appear to be single
(the ”single Am” problem, Murphy, 2013). For instance, Smalley et al. (2014)
suggest that 30%-40% of Am stars are either single or in wide binary systems,
while Murphy et al. (2012) presented the intriguing case of KIC 3429637, a
single pulsating Am star observed by Kepler with no evidence of frequency
modulation by a possible companion. In these cases of single Am stars or in
wide binary systems, the origin of both Am stars and their low rotational
velocities is unclear.

Hot-Jupiter planets present short orbital periods ($<$ 10 d) and large
planetary masses ($>$ 0.1 MJup), that is, they are gas giants orbiting very
close to their stars (e.g., Wang et al., 2015). By analogy, we refer to as
hot-brown dwarfs (hot-BDs) those objects with short orbital periods and having
masses above $\sim$13 MJup, the minimum mass to burn Deuterium (e.g., Grossman
& Graboske, 1973; Saumon et al., 1996). Recently, Ŝubjak et al. (2020)
(hereafter, S20) discovered TOI-503, the first Am star which hosts a hot-BD
companion. They found that the rotation period of the star is similar to the
orbital period of the hot-BD (P${}_{orb}\sim$ 3.67 d), which could be
indicative of synchronism between them. They suggest that the presence of
close low-mass companions to Am stars could be possibly invoked to explain
their small rotational velocities, where the tidal braking of a more distant
stellar companion is not significant. Then, Addison et al. (2021) (hereafter,
A21) found an ultra-hot Jupiter222Ultra-hot Jupiter planets present dayside
temperatures of 2200 K or more (Casasayas-Barris et al., 2019). with a mass of
3.12 MJup orbiting around the star MASCARA-5/TOI-1431, which is found to be
also a peculiar Am star. They wonder if the presence of a giant planet can
contribute to their host being an Am star, without evidence of a stellar
binary. In this way, both works S20 and A21, suggest a possible link between
Am stars and the presence of hot low-mass companions (hot-Jupiters and hot-
BDs), with a lower influence of an additional stellar component. Notably, it
is also worthwhile to mention the recent detection of four additional early-
type stars (KELT-17, KELT-19A, KELT-26 and WASP-189) hosting hot-Jupiter
companions, being confirmed or likely Am stars (Saffe et al., 2020; Siverd et
al., 2018; Rodríguez Martínez et al., 2020; Lendl et al., 2020).

Can hot low-mass companions help to mitigate the ”single Am” problem? Can they
facilitate -in some way- the rise of Am stars, as suggested by S20 and A21?
What is the frequency of Am stars (and chemically peculiar stars in general)
hosting hot-low mass companions? These important questions could be addressed
searching for planets around Am stars, as suggested by A21. Alternatively, it
is also possible to perform a chemical analysis of early-type stars hosting
hot-low mass companions to determine the likely presence of Am stars within
this group. This is one of the motivations of the present work. It would be
easier to confirm systems with low $v\sin i$, which could preferentially be
selecting stars which are susceptible to the Am phenomenon. In the last few
years, the population of hot-low mass companions around early-type stars is
slowly growing compared to late-type stars (e.g. Siverd et al., 2018;
Rodríguez Martínez et al., 2020; Lendl et al., 2020), giving us the
opportunity to perform this analysis. If such possible relation between Am
stars and hot-low mass companions is confirmed, specific surveys targeting
(single) Am stars could be important in helping to detect hot-Jupiters and/or
hot-BDs. In addition, these ideas could provide valuable insights into the
formation and evolution of both Am stars and their companions.

Stars that apparently display a peculiar spectra could be confirmed (or
rejected) through a detailed abundance analysis. For example, Catanzaro et al.
(2019) studied a group of 62 Am stars, previously selected with an abundance-
based criterion (subsolar values of Ca, Sc and overabundances of Ti, Cr, Mn
and Fe). For the case of stars with planets, the exoplanet host star WASP-33
was initially classified as Am (Grenier et al., 1999) and mentioned as such in
different works (S20, A21). However, Collier Cameron et al. (2010) caution
that no obvious Am characteristics are visible in the spectrum of this star,
other than slightly weak Ca II H&K lines. Then, a recent abundance analysis
including several species showed that WASP-33 do not belong to this class
(Saffe et al., 2021). Conversely, the exoplanet host star KELT-17 was
initially considered as a solar-like composition object, while a detailed
chemical analysis showed a clear Am pattern (Saffe et al., 2020, 2021). These
examples highlight the importance of a detailed chemical analysis including
several species, in order to confidently detect a bona fide Am chemical
pattern. This motivates us to perform a detailed chemical analysis of early-
type stars hosting hot low-mass companions in order to compare its chemical
pattern with those of Am stars and other chemically peculiar stars.

Our sample includes a number of stars without previous detailed abundance
analysis (such as TOI-587, TOI-681, CoRoT-3, KELT-1, KELT-25, MASCARA-4,
HAT-P-69 and HAT-P-70), being to our knowledge the largest group of early-type
stars with hot-companions homogeneously studied to date. We report the
detection of the new Am star HATS-70 which hosts a hot-BD, and rule out the Am
class for WASP-189 which hosts a hot-Jupiter planet (in both cases, our
results differ from the composition previously reported). This work is part of
our ongoing program aimed to study early-type stars with planets (Saffe et
al., 2020, 2021), which are poorly studied in general compared to their late-
type FGKM counterparts.

This work is organized as follows. In Sect. 2, we describe the observations
and data reduction. In Sect. 3, we present the stellar parameters and chemical
abundance analysis. In Sect. 4, we show the results and discussion. Finally,
in Sect. 5, we highlight our main conclusions.

##  2 Stellar samples and observations

We start by compiling a list of early-type stars hosting hot-Jupiters and/or
hot-BDs taken from the Extrasolar Planets
Encyclopaedia333http://exoplanet.eu/. We also included a group of early-type
stars with hot-BDs taken from the literature (Deleuil et al., 2008; Collier
Cameron et al., 2010; Siverd et al., 2012, 2018; Bieryla et al., 2014; Lund et
al., 2017; Talens et al., 2017; Gaudi et al., 2017; Temple et al., 2017; Zhou
et al., 2016, 2019a, 2019b; Ŝubjak et al., 2020; Rodríguez Martínez et al.,
2020; Lendl et al., 2020; Dorval et al., 2020; Grieves et al., 2021). The
complete sample studied in this work is shown in Table 1. We present the star
names, visual magnitude V, instrument, the signal-to-noise S/N measured near
5000 Å, together with information about their low-mass companions, that is,
orbital period and mass. The last column presents the reference for the
orbital period and mass of their low-mass companions. In this way, the final
sample comprises 19 objects, including mainly A-type and some early F-type
stars, all having low-mass companions. Within this group, 7 stars host hot-BDs
and 12 stars host hot-Jupiters. Aiming to increase the number of stars, we
included 8 stars with previously obtained stellar parameters (Saffe et al.,
2020, 2021). This allows us to obtain, to our knowledge, the largest possible
sample of early-type stars with hot-low mass companions analyzed to date, with
uniformly derived abundances. We caution that the sample includes also the
star TOI-681 for completeness, with a somewhat larger BD orbital period of
15.78 d (Grieves et al., 2021).

Table 1: Sample of stars studied in this work.

Star name | Alternative | Vmag | Instrument | S/N | Companion | Companion | Companion
---|---|---|---|---|---|---|---
| name |  |  | @5000 Å | period [d] | mass [MJup] | reference
Hot-BD host stars |  |  |  |  |
HATS-70 | TIC 98545929 | 12.57 | FEROS | 125 | 1.89 | 12.9${}^{+1.8}_{-1.6}$ | R1
TOI-503 | BD+13 1880 | 9.43 | GRACES | 230 | 3.68 | 53.7$\pm$1.2 | R2
CoRoT-3 | TIC 392353449 | 13.29 | HARPS | 85 | 4.26 | 21.66$\pm$1.00 | R3
KELT-1 | TOI-1476 | 10.70 | HARPS-N | 105 | 1.22 | 27.38$\pm$0.93 | R4
TOI-681 | TIC 410450228 | 10.89 | REOSC | 180 | 15.78 | 88.7${}^{+2.5}_{-2.3}$ | R5
TOI-587 | HD 74162 | 7.80 | REOSC | 330 | 8.04 | 81.1${}^{+7.1}_{-7.0}$ | R5
KELT-25 | TOI-626 | 9.66 | FEROS | 290 | 4.40 |  $<$ 64 | R6
Hot-Jupiter host stars |  |  |  |  |
WASP-33 | HD 15082 | 8.30 | HIRES | 250 | 1.22 |  $<$ 4.1 | R7
WASP-167 | KELT-13 | 10.50 | HARPS | 205 | 2.02 |  $<$ 8.0 | R8
WASP-189 | HR 5599 | 6.62 | HARPS | 1105 | 2.72 | 1.99${}^{+0.16}_{-0.14}$ | R9
KELT-9 | HD 195689 | 7.59 | HARPS-N | 550 | 1.48 | 2.88$\pm$0.84 | R10
KELT-17 | BD+14 1881 | 9.29 | REOSC | 210 | 3.08 | 1.31${}^{+0.28}_{-0.29}$ | R11
KELT-20 | MASCARA-2 | 7.59 | HARPS-N | 450 | 3.47 |  $<$ 3.382 | R12
MASCARA-1 | HD 201585 | 8.30 | HARPS-N | 560 | 2.15 | 3.7$\pm$0.9 | R13
HAT-P-49 | HD 340099 | 10.33 | SOPHIE | 135 | 2.69 | 1.730$\pm$0.205 | R14
MASCARA-4 | HD 85628 | 8.19 | FEROS | 180 | 2.82 | 3.1$\pm$0.9 | R15
HAT-P-70 | HD 287325 | 9.47 | FEROS | 210 | 2.74 |  $<$ 6.78 | R16
KELT-19A | BD+07 1721 | 9.89 | GRACES | 240 | 4.61 |  $<$ 4.1 | R17
HAT-P-69 | TYC 215-1594-1 | 9.77 | GRACES | 220 | 4.79 | 3.58${}^{+0.58}_{-0.58}$ | R16

Companions data: R1 (Zhou et al., 2019a), R2 (Ŝubjak et al., 2020), R3
(Deleuil et al., 2008), R4 (Siverd et al., 2012), R5 (Grieves et al., 2021),
R6 (Rodríguez Martínez et al., 2020), R7 (Collier Cameron et al., 2010), R8
(Temple et al., 2017), R9 (Lendl et al., 2020), R10 (Gaudi et al., 2017), R11
(Zhou et al., 2016), R12 (Lund et al., 2017), R13 (Talens et al., 2017), R14
(Bieryla et al., 2014), R15 (Dorval et al., 2020), R16 (Zhou et al., 2019b),
R17 (Siverd et al., 2018).

We have an ongoing observational campaign aimed to study early-type stars with
planets. The spectra of the stars TOI-503, KELT-19A and HAT-P-69 were acquired
through the Gemini Remote Access to CFHT ESPaDOnS Spectrograph (GRACES). This
device takes advantage of the high-resolution ESPaDOnS444Echelle
SpectroPolarimetric Device for the Observation of Stars spectrograph, located
at the Canada-France-Hawaii Telescope (CFHT) and fed by an optical fiber
connected to the 8.1-m Gemini North telescope at Maunakea, Hawaii. We used the
1-fiber object-only observing mode, which provides an average resolving power
of $\sim$67500 between 4500 and 8500
Å555http://www.gemini.edu/sciops/instruments/graces/spectroscopy/spectral-
range-and-resolution. The observations were taken during the observing runs
2021B and 2022A (Programs ID: GN-2021B-Q-106 and GN-2022A-Q-305, PI: Carlos
Saffe). The final spectral coverage is 4050-9000 Å, and the signal-to-noise
ratio (S/N) per pixel resulted $\sim$230 measured at $\sim$5000 Å in the
combined spectra. GRACES spectra were reduced using the code OPERA666Open
source Pipeline for ESPaDOnS Reduction and Analysis (Martioli et al., 2012).
More recent documentation on OPERA can be found at the ESPECTRO project
webpage777http://wiki.lna.br/wiki/espectro.

The spectra of the stars HATS-70, KELT-25, MASCARA-4 and HAT-P-70 were taken
with the Fiber-fed Extended Range Optical Spectrograph (FEROS), which provides
a high-resolution (R$\sim$48000) spectra when illuminated via the 2.0 arcsec
aperture on the sky in the unbinned mode. FEROS is installed in the Max Planck
Gesselschaft (MPG) 2.2-m telescope at the European Southern Observatory (ESO)
in La Silla, Chile. The data were taken during the observing runs 2021B and
2022A (Programs ID: 0108.A-9012(A) and 0109.A-9023(A), PI: Marcelo Jaque
Arancibia), including three individual spectra per object, followed by a ThAr
lamp in order to obtain an appropriate wavelength solution. The spectra were
reduced using the FEROS Data Reduction
System888https://www.eso.org/sci/facilities/lasilla/instruments/feros/tools/DRS.html
(DRS). The spectral coverage resulted between 3800-8600 Å, approximately, and
S/N per pixel measured at $\sim$5000 Å resulted in $\sim$210.

In addition, the spectra of the stars TOI-587, TOI-681 and KELT-17 were
obtained at the Complejo Astronómico El Leoncito (CASLEO) during the observing
runs 2019A, 2021B and 2022A (PI: Carlos Saffe). We used the _Jorge Sahade_
2.15-m telescope equipped with a REOSC echelle spectrograph999On loan from the
Institute d’Astrophysique de Liege, Belgium. and a TEK 1024$\times$1024 CCD
detector. The REOSC spectrograph uses gratings as cross-dispersers. We used a
grating with 400 lines mm-1, which provides a resolving power of $\sim$ 13000
covering the spectral range of $\lambda\lambda$3800–6500. Three individual
spectra for each object were obtained and then combined, reaching a final S/N
per pixel of $\sim$240 measured at $\sim$5000 Å. The data were reduced with
Image Reduction and Analysis Facility (IRAF)101010IRAF is distributed by the
National Optical Astronomical Observatories, which is operated by the
Association of Universities for Research in Astronomy, Inc., under a
cooperative agreement with the National Science Foundation. following the
standard recipe for echelle spectra (i.e., bias and flat corrections, order-
by-order normalization, scattered light correction, etc.).

Finally, we used archive spectra for the case of HARPS, HARPS-N, HIRES and
SOPHIE spectrographs. General characteristics of these instruments are shown
in the Table 2, including the resolving power, CCD detector, pixel size,
telescope and approximate wavelength range. The reduction was performed by
using the Data Reduction Software (DRS) pipeline for the case of HARPS and
HARPS-N
spectra111111https://www.eso.org/sci/facilities/lasilla/instruments/harps/doc.html,
using the reduction package MAKEE 3 with HIRES
spectra121212http://www.astro.caltech.edu/ tb/makee/, and the DRS pipeline
with SOPHIE spectra131313http://www.obs-hp.fr/guide/sophie/sophie-
eng.shtml#drs. The continuum normalization and other operations (such as
Doppler correction and combining spectra) were performed using IRAF.

Table 2: General characteristics of the spectrographs used in this work. Instrument | R | CCD | Pixel | Telescope | Approx.
---|---|---|---|---|---
|  | Detector | size |  | wave. range
FEROS | 48000 | 2k x 4k |  15 $\mu$m | MPG 2.2m | 3800 - 8600
REOSC | 13000 | 1k x 1k |  24 $\mu$m | CASLEO 2.15m | 3800 - 6500
GRACES | 67000 | 2k x 4k |  13.5 $\mu$m | Gemini 8.2m | 4050 - 9000
HARPS | 115000 | 4k x 4k |  15 $\mu$m | La Silla 3.6m | 3800 - 6800
HARPS-N | 115000 | 4k x 4k |  15 $\mu$m | TNG 3.6m | 3800 - 6800
HIRES | 67000 | 2k x 4k |  15 $\mu$m | Keck 10m | 3750 - 9000
SOPHIE | 75000 | 4k x 2k |  15 $\mu$m | OHP 1.93m | 3900 - 6800

##  3 Stellar parameters and abundance analysis

Stellar parameters were determined similarly to previous works. The effective
temperature Teff and superficial gravity $\log g$ were first estimated by
using the Strömgren uvby$\beta$ mean photometry of Hauck & Mermilliod (1998)
or by taking previously published results. We applied the calibration of
Napiwotzki et al. (1993) within the program TempLogG (Kaiser, 2006) and
derredenned colors according to Domingo & Figueras (1999), in order to derive
the fundamental parameters. These initial values were then refined (when
necessary and/or possible) by enforcing excitation and ionization balances of
the iron lines. The same strategy was previously applied in the literature
with early-type stars (e.g., Saffe & Levato, 2014; Saffe et al., 2021;
Alacoria et al., 2022). The values derived in this way are listed in the Table
3, including the values taken from our previous work (Saffe et al., 2021). The
resulting average dispersions are $\sim$186 K and $\sim$0.15 dex for Teff and
$\log g$, respectively.

Table 3: Fundamental parameters adopted for the stars in this work.

Star | Teff | $\log g$ | vmicro | $v\sin i$ | Reference
---|---|---|---|---|---
| [K] | [dex] | [km s-1] | [km s-1] |
HATS-70 | 7905 $\pm$ 180 | 4.17 $\pm$ 0.12 | 3.28 $\pm$ 0.82 | 39.7 $\pm$ 1.2 | this work
TOI-503 | 7790 $\pm$ 140 | 4.08 $\pm$ 0.10 | 3.24 $\pm$ 0.81 | 27.2 $\pm$ 0.9 | this work
CoRoT-3 | 6785 $\pm$ 150 | 4.25 $\pm$ 0.15 | 1.92 $\pm$ 0.48 | 17.8 $\pm$ 0.3 | this work
KELT-1 | 6720 $\pm$ 210 | 4.23 $\pm$ 0.22 | 1.81 $\pm$ 0.45 | 50.3 $\pm$ 1.7 | this work
TOI-681 | 7490 $\pm$ 160 | 4.27 $\pm$ 0.14 | 2.99 $\pm$ 0.75 | 43.8 $\pm$ 1.8 | this work
TOI-587 | 10250 $\pm$ 300 | 4.25 $\pm$ 0.13 | 1.18 $\pm$ 0.30 | 31.5 $\pm$ 0.8 | this work
KELT-25 | 8100 $\pm$ 190 | 3.86 $\pm$ 0.18 | 3.31 $\pm$ 0.83 | 112.9 $\pm$ 4.2 | this work
WASP-33 | 7373 $\pm$ 164 | 4.14 $\pm$ 0.20 | 2.86 $\pm$ 0.71 | 82.5 $\pm$ 2.0 | S21 & this work
WASP-167 | 6770 $\pm$ 210 | 4.05 $\pm$ 0.14 | 1.90 $\pm$ 0.47 | 52.0 $\pm$ 2.0 | S21 & this work
WASP-189 | 7946 $\pm$ 136 | 3.85 $\pm$ 0.12 | 3.30 $\pm$ 0.82 | 95.2 $\pm$ 3.3 | S21 & this work
KELT-9 | 9329 $\pm$ 118 | 4.00 $\pm$ 0.14 | 2.27 $\pm$ 0.57 | 113.5 $\pm$ 4.6 | S21 & this work
KELT-17 | 7471 $\pm$ 210 | 4.20 $\pm$ 0.14 | 2.50 $\pm$ 0.63 | 43.0 $\pm$ 0.6 | S21 & this work
KELT-20 | 8652 $\pm$ 160 | 4.11 $\pm$ 0.20 | 3.03 $\pm$ 0.76 | 111.0 $\pm$ 3.1 | S21 & this work
MASCARA-1 | 7687 $\pm$ 238 | 4.09 $\pm$ 0.14 | 3.17 $\pm$ 0.79 | 99.0 $\pm$ 3.9 | S21 & this work
HAT-P-49 | 6730 $\pm$ 234 | 4.02 $\pm$ 0.14 | 1.82 $\pm$ 0.46 | 14.8 $\pm$ 0.2 | S21 & this work
MASCARA-4 | 7810 $\pm$ 165 | 3.98 $\pm$ 0.12 | 3.25 $\pm$ 0.81 | 42.3 $\pm$ 1.3 | this work
HAT-P-70 | 8450 $\pm$ 195 | 4.08 $\pm$ 0.14 | 3.19 $\pm$ 0.80 | 103.5 $\pm$ 4.1 | this work
KELT-19A | 7500 $\pm$ 130 | 4.13 $\pm$ 0.21 | 3.00 $\pm$ 0.75 | 81.5 $\pm$ 2.1 | this work
HAT-P-69 | 7750 $\pm$ 245 | 4.15 $\pm$ 0.09 | 3.21 $\pm$ 0.80 | 76.1 $\pm$ 1.7 | this work

S21 (Saffe et al., 2021).

The projected rotational velocities, $v\sin i,$ were estimated by fitting most
Fe I and Fe II lines in the spectra. Synthetic spectra were calculated using
the program SYNTHE (Kurucz & Avrett, 1981) together with ATLAS12 (Kurucz,
1993) model atmospheres, convolved with a rotational profile (using the
Kurucz’s command rotate) and with an instrumental profile for each
spectrograph (using the command broaden). Microturbulence velocity vmicro was
estimated as a function of Teff following the formula of Gebran et al. (2014),
which is valid for $\sim$6000 K $<$ Teff $<$ $\sim$10000 K. We adopt for
vmicro an uncertainty of $\sim$25 $\%$, as suggested by Gebran et al. (2014),
and then this uncertainty was taken into account in the abundance error
calculation.

We applied an iterative procedure to determine the chemical abundances for the
stars in our sample. In a first step, we computed an ATLAS12 (Kurucz, 1993)
model atmosphere adopting initially solar abundances from Asplund et al.
(2009). The corresponding abundances are then obtained by fitting the observed
spectra with the program SYNTHE (Kurucz & Avrett, 1981). With the new
abundance values, we derived a new model atmosphere and restarted the process
again. In each step, opacities were calculated for an arbitrary composition
and vmicro using the opacity sampling (OS) method, similar to previous works
(Saffe et al., 2020, 2021; Alacoria et al., 2022). If necessary, Teff and
$\log g$ were refined to achieve the balance of Fe I and Fe II lines. In this
way, parameters and abundances are consistently derived using specific
opacities rather than solar-scaled values. Possible differences originating
from the use of solar-scaled opacities instead of an arbitrary composition
were recently estimated for solar-type stars (Saffe et al., 2018, 2019). These
differences could became particularly important when modeling chemically
peculiar stars, where solar-scaled models could result in a very different
atmospheric structure (e.g. Piskunov & Kupka, 2001) and reach abundance
differences up to 0.25 dex (Khan & Shulyak, 2007).

We derived the chemical abundances for 25 different species, including C I, O
I, Na I, Mg I, Mg II, Al I, Al II, Si II, Ca I, Ca II, Sc II, Ti II, Cr II, Mn
I, Fe I, Fe II, Ni II, Co I, Zn I, Sr II, Y II, Zr II, Ba II, Nd II and Eu II.
The laboratory data and atomic line list used in this work are described in
Saffe et al. (2021). Figure 1 presents an example of observed and synthetic
spectra (magenta and blue dotted lines) for some stars in our sample.

Figure 1: Example of observed and synthetic spectra (magenta and blue dotted
lines) for some stars in our sample, sorted by $v\sin i$.

The uncertainty in the abundance values was estimated considering different
sources. We estimated the measurement error e1 from the line-to-line
dispersion as $\sigma/\sqrt{n}$, where $\sigma$ is the standard deviation and
n is the number of lines. For elements with only one line, we adopted for
$\sigma$ the standard deviation of the iron lines. Then, we determined the
contribution to the abundance error due to the uncertainty in stellar
parameters. We modified Teff, $\log g$ and vmicro by their uncertainties and
recalculated the abundances, obtaining the corresponding differences e2, e3
and e4141414We adopt a minimum of 0.01 dex for these errors.. Finally, the
total error etot was estimated as the quadratic sum of e1, e2, e3 and e4. The
abundances with their total error etot, the individual errors e1 to e4 and the
number of lines n, are presented in the Tables 4 to 14 of the Appendix.

In general, departures from local thermodynamic equilibrium (LTE) are
pronounced in stars with high temperature and with low gravity and
metallicity. Then, some particular transitions should be taken with caution
for early-type stars. For the case of O i, Sitnova et al. (2013) estimated
non-LTE (NLTE) corrections for stars with Teff between 5000 K and 10000 K.
They showed that NLTE effects lead to a strengthening of O i lines. Then, we
estimated O i NLTE abundance corrections for the IR triplet at 7771 Å and for
the line 5328.98 Å for the stars in our sample, depending on their fundamental
parameters. If available, we prefer to use the line at 5328.98 Å which
presents considerably lower NLTE effects than the IR triplet (Sitnova et al.,
2013).

NLTE effects could also be present in species such as Mg i and Mg ii (e.g.
Przybilla et al., 2001; Bergemann et al., 2017; Alexeeva et al., 2018).
Przybilla et al. (2001) found that the intense line Mg II 4481 Å
systematically yields notably higher abundances due to NLTE effects (between
0.2 - 0.8 dex for early-type stars), while Alexeeva et al. (2018) estimate
that the average difference between Mg I and Mg II diminish from $\sim$0.23
dex in LTE to $\sim$0.09 dex in NLTE. This could explain, at least in part,
the higher abundances of Mg II compared to Mg I observed in some stars.
However, the authors also caution that the difference Mg II $-$ Mg I even in
NLTE could amount up to $\sim$0.24 dex for metal-poor stars, for a reason that
requires future investigation. We estimated Mg I NLTE abundance corrections
following Bergemann et al. (2017), who calculated Mg abundance departures
between LTE and NLTE by using 1D and spatially averaged $<$3D$>$ model
atmospheres. They showed that Mg I NLTE effects grow with increasing Teff, and
decreasing $\log g$ and [Fe/H], that is to say, they are sensitive to stellar
parameters. The lower NLTE effects of Mg I lines compared to the intense line
Mg II 4481 Å, result in more reliable abundance values for Mg. In addition to
these considerations, we prefer to avoid the line Mg I 5167.32 Å which is
often strongly blended with Fe I 5167.49 Å.

For the case of the calcium lines, we estimated NLTE abundance corrections
following Mashonkina et al. (2007), who determined abundances in LTE and NLTE.
They calculated a detailed Ca model atom and then obtained statistical
equilibrium populations by adopting a plane-parallel homogeneous MAFAGS model
atmosphere. We also note that, beyond NLTE effects, we prefer to avoid the Ca
II line 3968.47 Å which is strongly blended with the Balmer line H$\epsilon$
in many early-type stars.

##  4 Discussion

###  4.1 Defining Am stars using abundances

As explained in the Introduction, a number of authors consider Am stars as
objects showing underabundances of Ca and/or Sc, together with overabundances
of Fe-peak and heavier elements (Preston, 1974; Fossati et al., 2007, 2009;
Catanzaro et al., 2019). In this work we consider ”Am stars” those early-type
stars showing the mentioned characteristics. However, following the work of
Conti (1970), other authors prefer a wider definition for the class, showing
underabundances of Ca and/or Sc, and/or overabundances of the Fe group and
heavier elements (Conti, 1970; Lane & Lester, 1987; Adelman, 1994; Adelman et
al., 1997; Yüce & Adelman, 2014). As a consequence of the wider definition, it
would be possible to increase the number of Am stars found. Those objects
satisfying the wider definition rather than the first definition, are referred
in this work to as ”possible Am” stars. We will estimate the incidence of Am
stars using both definitions, ”Am stars” and ”possible Am stars”, where the
reader could decide the more appropriate for their own preference.

###  4.2 Am stars in our sample

We discuss in this Section those stars in our sample that could be related to
Am stars in general, that is to say, Am stars and possible Am stars. We
compared the abundances of the stars in our sample with an average pattern of
62 Am stars derived by Catanzaro et al. (2019). In particular, we present in
Fig. 2 the abundances of the star KELT-19A (black) compared with those of Am
stars (blue). The square bracket notation [N/H] corresponds to abundance
values relative to the Sun151515[N/H] $=$ log(N/H)star $-$ log(N/H)Sun, with
solar values taken from Asplund et al. (2009). The Figure includes two panels,
corresponding to elements with atomic number z$<$32 and z$>$32\. Ca and Sc
show subsolar values (with Ca slightly higher than Sc), while Fe and other
metals show overabundances that resemble those of Am stars. An apparent
disagreement is Na (z$=$11) showing suprasolar rather than subsolar values,
however some Am stars also present suprasolar values (e.g. HD 267 and HD
176716, Catanzaro et al., 2019). Then, abundance values agree with the
definition of Am stars given in the previous Section. Siverd et al. (2018)
found a transiting hot-Jupiter planet orbiting around this star, and suggest
its possible Am nature by noticing in the spectra an enhancement in the
metallic lines together with a calcium deficiency. Then, we confirm the
previous suggestion of Siverd et al. (2018) and consider that KELT-19A could
be safely identified as an Am star.

Figure 2: Abundances of the star KELT-19A (black), compared to the average
pattern of Am stars (blue). The panels correspond to elements with z$<$32 and
z$>$32.

Zhou et al. (2019a) first identified a transiting companion to the early-type
star HATS-70, with a mass of 12.9 MJup and a period of 1.89 d. They performed
an initial estimation of the stellar parameters by fitting the observed
spectra with a grid of synthetic spectra with steps of 0.5 dex in [Fe/H].
Then, the final stellar parameters adopted were estimated within a global
analysis via a spectral energy distribution (SED) fit. However, the authors
caution that the metallicity is poorly constrained in the SED fitting and
adopt an approximate near-solar value of [Fe/H]$\sim$0.03 dex (see their
Appendix 1). Hosting a hot-companion, we take the opportunity and performed a
detailed abundance analysis for this early-type star. We present in the Fig. 3
the observed spectra of the star HATS-70 (black), compared to synthetic
spectra with our abundances (blue dotted) and adopting a solar composition
(red). The line identifications are shown using three numbers: wavelength,
atomic number (and ionization state as decimals), and intensity (between 1 and
0). For instance, the label ”3933.68 20.01 0.1234” would correspond to the
line Ca ii 3933.68 Å. The two panels show an example of different spectral
regions. The left panel shows that some metals are clearly enhanced, while the
right panel shows that the line Sc ii 5226 Å is notably weaker than expected
for a solar composition. Then, Fig. 3 suggests that it would be difficult to
fit the spectra of HATS-70 adopting a solar composition.

Figure 3: Observed spectra of the star HATS-70 (black), compared to synthetic
spectra with our abundances (blue dotted) and adopting a solar composition
(red). The two panels correspond to two different spectral regions. Line
identifications present the wavelength, atomic number (with ionization state
as decimals), and intensity (between 1 and 0).

We present in Fig. 4 the abundances of the star HATS-70 (black) compared to
the average pattern of Am stars (blue). Most abundance values reasonably agree
with those of Am stars. Ca and Sc present subsolar values (being Ca slightly
higher than Sc), similar to other Am stars. Fe presents slightly lower
abundances than average Am stars, being still suprasolar ([Fe
i/H]$=$0.18$\pm$0.10 dex, [Fe ii/H]$=$0.22$\pm$0.11 dex). For example, the Am
stars HD 134214, HD 139939 and HD 159545 present a very similar Fe enhancement
(Catanzaro et al., 2019). The Ni ii abundance (0.50$\pm$ 0.14 dex) was derived
using the line 4015.47 Å, showing an enhancement in good agreement with other
Am stars. The heavy elements (Sr, Y, Zr) and the most heavy elements (Ba, Nd)
show strong overabundances between $\sim$0.5-2.0 dex, corresponding to a
chemically peculiar (non-solar composition) star. Having subsolar Ca and Sc,
together with suprasolar values of most Fe-peak and heavier elements, HATS-70
fits the previous definition of Am stars.

Figure 4: Abundances of the star HATS-70 (black), compared to the average
pattern of Am stars (blue). The panels correspond to elements with z$<$32 and
z$>$32.

S20 reported the discovery of TOI-503b, the first BD known to orbit a
metallic-lined A star (Am star). They derived abundances for some key elements
of the host star (Fe, Ni, Ca, Sc and Mg) which suggest its Am nature. We
included this star in our sample and performed a detailed abundance analysis
including as many species as possible. We present in Fig. 5 the abundances of
the star TOI-503 (black) compared to the average pattern of Am stars (blue).
The general behavior of the chemical species strongly resembles those of Am
stars, confirming the previous suggestion of S20 about the Am nature of this
object. We note relatively low abundances for C and O compared to average Am
stars. We corrected O abundances (derived from the oxygen triplet near
$\sim$7774 Å) of NLTE effects with an average correction of $\sim$-0.4 dex.
Also, we note that the Li i line 6707.8 Å is relatively weak although still
present in the spectra of TOI-503. We present in the Fig. 6 observed and
synthetic spectra of the star TOI-503 (black and blue dotted lines) in a
spectral region near the Li line. The corresponding abundance derived for
lithium ([Li i/H]$=$1.98$\pm$0.11 dex) resulted similar to other Am stars
(e.g. Iliev et al., 1998; Stateva et al., 2012). Then, showing underabundances
of both Ca and Sc, together with suprasolar values of most Fe-peak and heavier
elements, we consider TOI-503 as Am star.

Figure 5: Abundances of the star TOI-503 (black), compared to the average
pattern of Am stars (blue). The panels correspond to elements with z$<$32 and
z$>$32. Figure 6: Observed and synthetic spectra of the star TOI-503 (black
and blue dotted lines) near the Li i line 6707.8 Å. Line identifications are
similar to those of Fig. 3.

We included in our sample of early-type stars with hot-companions the star
KELT-17. The Am nature of this object was reported in our previous works
(Saffe et al., 2020, 2021) using the same model atmospheres and laboratory
data, showing subsolar Ca and Sc, together with suprasolar Fe-peak elements.
Then, in the complete sample of 19 early-type stars with hot-companions
studied in this work, we detected 4 Am stars: KELT-19A, HATS-70, TOI-503 and
KELT-17.

Next, we discuss some stars with a less clear Am identification. We present in
the Figs. 7, 8 and 9. the chemical pattern of the stars TOI-681, HAT-P-69 and
TOI-587 compared to the average pattern of Am stars. The star TOI-681 presents
subsolar Ca, solar Fe, together with enhanced Ti, Cr and some rare earths.
Then, following the definition given in the previous section, we consider this
object as a possible Am star. As we can see in the Fig. 8, the star HAT-P-69
presents subsolar Sc, solar Fe-peak species and some enhancements in heavier
elements (Sr, Y, Zr, Ba). These characteristics indicate that HAT-P-69 could
be identified as a possible Am star. The star TOI-587 presents subsolar Ca and
Sc, however most elements also present subsolar abundance values (O, Mg, Si,
Ti, Cr and Fe), together with some rare-earths enhancements. Its chemical
pattern shows mostly a metal-poor star, with the exception of few elements. So
we prefer to be cautious and not relate this object to the Am class. Finally,
in the complete sample of 19 stars, we have identified 4 Am stars and,
adopting a wider definition for the class, 2 possible Am stars (TOI-681 and
HAT-P-69).

Figure 7: Abundances of the star TOI-681 (black) compared to the average
pattern of Am stars (blue). Upper and lower panels correspond to elements with
z$<$32 and z$>$32. Figure 8: Abundances of the star HAT-P-69 (black) compared
to the average pattern of Am stars (blue). Upper and lower panels correspond
to elements with z$<$32 and z$>$32. Figure 9: Abundances of the star TOI-587
(black) compared to the average pattern of Am stars (blue). Upper and lower
panels correspond to elements with z$<$32 and z$>$32.

###  4.3 Other chemical patterns in our sample

We discuss in this section the chemical pattern of stars that do not seem to
be Am stars. We present in the Fig. 10 the chemical pattern of the star
KELT-1. This object presents slight overabundances in Si, Cr and Mn; Fe
resulted almost solar within errors; Ca is overabundant and Sc slightly
deficient (although solar within the errors). Then, although KELT-1 presents
slight suprasolar values in some species, the general behavior of the
abundances shown in Fig. 10 does not seem to be related to Am stars.

Figure 10: Abundances of the star KELT-1 (black) compared to the average
pattern of Am stars (blue). Upper and lower panels correspond to elements with
z$<$32 and z$>$32.

Within the group of stars with abundances determined in this work, some of
them resulted with solar or near-solar composition (CoRoT-3 and KELT-25), as
we can see in the Tables 4 to 14 of the Appendix. In the sample we also found
stars with slightly subsolar abundances (HAT-P-70 and MASCARA-4). Then, we
should add to this group 7 stars with chemical patterns derived in our
previous work (Saffe et al., 2021) that do not seem to display an Am spectra
(WASP-33, WASP-167, WASP-189, KELT-9, KELT-20, MASCARA-1 and HAT-P-49). In
this way, in the complete sample of 19 early-type stars hosting hot-
companions, we identified 4 Am stars and 2 possible Am stars, while the 13
remaining objects do not seem to be related to this class.

###  4.4 Tidal effects between Am stars and low-mass companions?

We explore in this section the possible tidal effects between the Am stars
detected and their hot-companions. S20 discovered TOI-503, the first Am star
which hosts a very close (a$\sim$0.06 au) BD companion. They found that the
rotation period of the star is similar to the orbital period of the BD
(P${}_{orb}=$3.67 days). The authors tried to study the tidal evolution and
circularization of the system. However, as discussed by the authors (their
section 4.3), the general uncertainty of tidal evolution models difficult to
estimate the interaction between the star and their hot companion. They showed
that the tidal circularization and tidal synchronization timescales are
notably shorter for lower values of the tidal quality factor Q’, using the
tidal model of Jackson et al. (2008). Alternatively, using the tidal model of
Zahn et al. (1977), they predict timescales one order of magnitude larger than
those using the model of Jackson et al. (2008). Then, they cannot draw firm
conclusions on the tidal evolution of the system (that is to say, if the star
TOI-503 circularized or not the BD’s orbit), given the uncertainty in the
appropriate tidal model and the values of Q’. However, they do not discard the
possibility that the presence of BDs and Am stars could be linked.

Zhou et al. (2019a) discovered a BD with a mass of 12.9 MJup orbiting the star
HATS-70. They estimate a rotational velocity of the star 5$\sigma$ different
than expected from a spin synchronized with the orbit of the BD companion.
Following the procedure of Matsumura et al. (2010), they found that the
angular momentum of HATS-70 does not reach the value required for the Darwin
stability criteria. They also estimated the timescale required by the planet
to modify the stellar spin following Hansen (2012), and found a excessively
large value. Then, although the obliquity of the BD is relatively low
($\lambda=$ 8.9${}^{+5.6}_{-4.5}$ deg), current models do not seem to support
a strong star-planet synchronization.

Siverd et al. (2018) found a hot-Jupiter planet orbiting the Am star KELT-19A.
The star belongs to a binary system together with KELT-19B, a late-type star
separated by $\sim$160 au. Although KELT-19A is a slow rotating star, the
authors suggest that its stellar companion is too distant (and its planetary
companion likely too low in mass) to cause a significant tidal braking
(although without a numerical estimation of these effects). Instead, they
consider that the slow rotation of the star is probably either primordial or
was induced by a more efficient tidal braking mechanism. They also modeled the
orbital evolution of the planet due to tidal friction using the program POET
(Penev et al., 2014), under the assumption of a non-rotating star, ignoring a
larger-distance Type II migration, scattering-induced migration and other
possible hot-Jupiter planets. It only considers a close-in migration due to
tidal friction. However, even under these simplifying assumptions, the authors
caution that the estimated evolution of the semimajor-axis with time strongly
depends of the stellar tidal quality factor Q’, which presents large
uncertainties and is difficult to estimate (e.g. Ogilvie, 2014; Petrucci et
al., 2020).

Stateva et al. (2012) searched for possible tidally driven abundance anomalies
in a sample of Am stars than belong to binary systems. They found that the
”metallicism” defined as [Ca/Fe] $=$ [Ca/H] $-$ [Fe/H] (see e.g. Conti, 1970)
seem to slightly increase with increasing orbital eccentricity, while no trend
was detected with orbital period. They suggest that a mechanism should
stabilize the atmosphere and reduce the mixing, due to the presence of a
stellar companion. In our case, the small number statistics (only 2 Am stars
hosting hot-BDs and 2 Am stars hosting hot-Jupiter planets) preclude us to
reach robust conclusions. These correlations would indicate that tidal effects
could play a role in the development of Am stars.

Current estimations of the tidal evolution (including a number of
uncertainties) do not seem to support a strong star-planet synchronization
between Am stars and their hot-companions. However, as a consequence of the
small separations, a tidal interaction is not ruled out (e.g. Mathis et al.,
2013; Ogilvie, 2014). A close star-planet interaction could also involve
additional mechanisms, such as the recently proposed resonant oscillations in
late-type stars (Lanza, 2022a, b).

###  4.5 Searching stellar companions to Am stars

We explore possible stellar companions to the Am stars found in our sample. We
searched our targets stars in different binary catalogs from the literature
(Pourbaix et al., 2004; Shaya & Olling, 2011; Tokovinin & Lepine, 2012;
Andrews et al., 2017). Shaya & Olling (2011) developed a Bayesian method to
discover non-random pairs using Hipparcos proper motions and parallaxes, while
Tokovinin & Lepine (2012) constructed a catalog of common-proper motion pairs
within 67 pc of the Sun based on the SUPERBLINK proper-motion survey. Andrews
et al. (2017) performed a search of binaries through a Bayesian formulation in
the Tycho-Gaia DR2 catalogs and derived likelihoods of Keplerian orbits. We
find no record for the Am stars KELT-19A, HATS-70, TOI-503 and KELT-17 in the
mentioned binary catalogs. We also searched the Am stars in the SB9 catalog of
spectroscopic binary orbits161616https://sb9.astro.ulb.ac.be/ (Pourbaix et
al., 2004) and found no entry.

In addition, we searched for signs of additional substellar and/or stellar
companions around the four Am stars KELT-19A, KELT-17, HATS-70, and TOI-503,
by using the 2-min cadence photometric data provided by the TESS mission
(Ricker et al., 2015). These publicly available data were analyzed with the
tools in the Lightkurve Python package (Lightkurve Collaboration et al.,
2018). Briefly, we first eliminated the data-points of the transits produced
by the planet around each star, then performed a detrending to remove any
remaining systematics in the light curve and, finally, ran the Transit Least
Squares code (TLS, Hippke & Heller, 2019), designed to detect periodic
transits from time-series photometry171717Although TLS is optimized to find
planetary transits, it can be used to search for signals of eclipses produced
by stellar and substellar companions.. We did not detect any feature that
could be clearly attributed to the transit of a new planet or the eclipse of a
companion around KELT-19A, KELT-17, HATS-70, and TOI-503 .

Siverd et al. (2018) performed adaptive optics observations of the Am star
KELT-19A with the Palomar Observatory Hale Telescope and found a late-type
G9V/K1V companion. The projected separation between both stars is
0”.64$\pm$0”.03, which corresponds to $\sim$160 au at a distance of 255$\pm$15
pc. They suggest that the companion is bound based on radial velocity
observations. However, given the separation between both stars, they consider
unlikely that the companion caused the tidal braking of the Am star. Zhou et
al. (2019a) found no additional stellar companions to the Am star HATS-70
within 2” in the Astralux lucky imaging observations, nor within 5” in the
Gaia DR2 catalog. They also note an agreement between the observed spectral
energy distribution (SED) and single star spectral models, ruling out a
possible nearby star. In addition, searching for fainter stellar blends, they
injected and attempt to recover an additional stellar signal in the out-of-
transit line profiles. No companions were found in this case.

S20 performed Y-band imaging observations of the Am star TOI-503 with the
Sinistro camera on the Las Cumbres Observatory and the Siding Spring
Observatory (SSO). They confirmed that there are no nearby or background
eclipsing binaries within 2’.5 of the star. In addition, no nearby sources
were identified in Gemini/NIRI and Keck/NIRC2 images taken in the Br-$\gamma$
filter within 2” of the target. Zhou et al. (2016) found a transiting planet
orbiting the star KELT-17, detected using radial velocities and Doppler
tomography. They discard possible blend scenarios (such as background
eclipsing binaries) by the detection of the Doppler tomographic signal. The
authors also performed a spectroscopic follow-up taking multi-epoch spectra of
the star using the WiFeS integral field spectrograph, which help to eliminate
the possibility of nearby eclipsing binaries.

###  4.6 Are Am stars and hot-low mass companions related?

As discussed in the previous sections, early-type stars with hot companions
show different chemical patterns, some of them being Am stars. To determine if
the presence of hot-companions could facilitate (in some way) the rise of Am
stars, it would be interesting to compare with the frequency of Am stars in
general. This comparison should be taken with caution, given the (still)
relatively low number of early-type stars with hot-companions detected. Smith
(1971) found an incidence of $\sim$31% of Am stars in a random sample of 70
bright A-type stars. Then, Abt (1981) studied a sample of 865 components of
visual multiples and found a frequency of $\sim$32% in the range A4-F1 (IV and
V), and decreasing to $\sim$5% in the hotter range A0-A3 (IV and V). More
recently, Gray et al. (2016) studied the spectra of 80447 low-resolution
spectra of the LAMOST-Kepler survey and applied the automatic classification
code MKCLASS. In this way, they found an incidence of 36.4% in the spectral
range A4-F1. Then, Qin et al. (2019) searched for metallic-lined Am stars in
more than 200000 low-resolution spectra of the LAMOST DR5 survey. They used a
combination of machine learning algorithms and then manual identifications,
and found a total of 9372 Am stars and 1131 Ap stars181818Ap stars are
chemically peculiar objects showing, for instance, intensified lines of Cr, Sr
and Eu (see e.g. Ghazaryan et al., 2018; Saffe & Levato, 2004).. They found
that most Am stars belong to temperature interval between $\sim$7250-8250 K
($\sim$F0-A4) with an incidence of $\sim$17%, while in hotter and colder
intervals the frequency is lower. The authors suggest that the lower frequency
compared to previous studies (Smith, 1971; Abt, 1981; Gray et al., 2016) is
due to the strict screening applied by Qin et al. (2019).

Chemical peculiarities similar to those of Am stars develop for stars hotter
than the Kraft break (Kraft, 1967) near $\sim$6200 K. Beyond this limit,
hotter stars present gradually thinner convective envelopes that allow the
diffusion processes to operate. The coldest Am star found in the catalog of
Ghazaryan et al. (2018) presents T${}_{\rm eff}\sim$6780 K. These values
indicate that stars hotter than $\sim$6750 K, could develop an Am peculiarity.
Our sample includes 4 stars with temperatures similar or lower than this value
(CoRoT-3, KELT-1, WASP-167 and HAT-P-49, see Table 3). These stars were
analyzed for completeness and will be excluded from the incidence
calculations. On the other hand, we also prefer to exclude the hot star
TOI-587 from the incidence calculations, having an estimated Teff of 10250
$\pm$ 300 K (the hotter Am star in the Ghazaryan’s catalog presents T${}_{\rm
eff}\sim$ 10350 K).

The group of hot-BD hosts studied in this work includes 7 objects, with 4 of
them in the approximate Teff range of Am stars. Within this group, we found
the Am stars HATS-70 and TOI-503, being roughly a fraction of 2/4 ($\sim$50%).
If we also include the star TOI-681, classified as a possible Am star, the
fraction of Am stars between hot-BDs would rise to 3/4 ($\sim$75%). In other
words, in the group of hot-BD hosts, the fraction of Am stars would be roughly
between 50-75%. Similarly, the group of hot-Jupiters host stars studied in
this work includes 12 objects, where 10 of them present Teff hotter than
$\sim$6750 K. Within this group, we found the Am stars KELT-17 and KELT-19A,
being roughly a fraction of 2/10 ($\sim$20%). Recent literature works showed
that the hot-Jupiter host stars MASCARA-5 (Addison et al., 2021) and KELT-26
(Rodríguez Martínez et al., 2020) are also likely Am stars (they are not
included in our sample). If we include these two likely Am stars from the
literature, and also the possible Am star HAT-P-69 identified in this work,
the fraction of Am stars would rise to 5/12 ($\sim$42%). Then, in the group of
hot-Jupiters host stars, the fraction of Am stars would be roughly 20-42%.

If we combine the groups of hot-BD and hot-Jupiters hosts (all the stars in
our sample) we would have 19 stars, with 14 of them in the approximate Teff
range of Am stars. In this group, we detected a total of 4 Am stars, being
roughly a fraction of 4/14 ($\sim$28%). If we also include the likely Am stars
MASCARA-5, KELT-26 and the possible Am star HAT-P-69, the fraction would rise
to 7/16 ($\sim$44%). Then, in the group of hot-low mass companions (including
both hot-BDs and hot-Jupiters), the incidence of Am stars would be roughly
28-44%.

The number of early-type stars having hot-companions detected is still low,
and then, we stress that the comparisons should be taken with caution. On the
one hand, the incidence of Am stars between hot-BDs is roughly 50-75%, which
would be higher than the frequency of Am stars in general (31-36%). On the
other hand, the incidence of Am stars with hot-Jupiters hosts is roughly
between 20-42%, which is not significantly higher than those of Am stars in
general. If confirmed, these incidence estimations would support the idea that
the presence of hot-BDs could help in the development of Am stars, different
to the case of hot-Jupiter planets. Tidal synchronization calculations
indicate an apparent lack of a strong star-planet (or star-BD) effect, however
they still present a number of uncertainties, as discussed in the previous
section. The initial trends found here would imply that the presence of hot-
BDs could possibly help to mitigate the ”single Am” problem, different to the
case of hot-Jupiter planets. Then, for the case of Am stars hosting hot-
Jupiter planets, other mechanisms should be invoked to explain the presence of
these single Am stars or Am stars in wide binary systems.

If confirmed, the higher incidence of Am stars between stars hosting hot-BDs
(50-75%) than in stars hosting hot-Jupiter planets (20-42%), would support the
idea that the presence of hot-BDs could affect their host stars. This could be
related to the higher mass of BDs compared to giant planets, or to a different
formation mechanism. These interesting results should be taken with caution,
and would deserve a further confirmation using a larger sample of stars.

###  4.7 On the origin of single Am stars

A search for eclipsing binaries among Am stars suggests that around 60–70 per
cent of Am stars are spectroscopic binaries (Smalley et al., 2014), which is
consistent with radial velocity studies (Abt & Levy, 1985; Carquillat & Prieur
et al., 2007). However, the presence of single Am stars and the presence of Am
stars in wide binary systems cannot be ruled out. Some works mention the
possible existence of single Am stars (e.g. Conti, 1969; Rao & Abhyankar,
1991), including stars such as HD 29499 (Debernardi et al., 2000), HD 8801
(Henry & Fekel, 2005) and KIC 3429637 (Murphy et al., 2012). Also, different
studies of binarity indicate that there are some apparently single Am stars
that were presumably born with slow rotation (Smalley et al., 2014; Balona et
al., 2015). In Sect. 4.5 we showed that the Am stars detected here belong to
wide binary systems (KELT-19A) or seem to be single (KELT-17, HATS-70 and
TOI-503). In the previous Section, we showed that the presence of hot-BDs
could possibly play a role in the development of single Am stars, however this
is different for the case of hot-Jupiter planets. Then, other mechanisms
should be explored to explain the origin of these Am stars.

The presence of a magnetic field during pre-main-sequence is considered an
efficient braking mechanism. However, Am stars only present weak or ultra-weak
magnetic fields (e.g., Folsom et al., 2013; Blazère et al., 2016, 2020). Abt &
Levy (1985) suggest that evolutionary effects (such as expansion) and
rotational spin-down is the reason for the Am characteristics in long-period
binaries and some single stars. However, Henry & Fekel (2005) noted that the
main-sequence nature of HD 8801 argues against this evolutionary scenario.
Instead, they speculate that some additional factor could be responsible for
the slow rotation of the single Am stars, such as the presence of planetary
systems. Rao & Abhyankar (1991) suggest that rapidly rotating A stars could
develop shell structures at some time, then the star has to slow down in its
shell and exhibit Am characteristics. However, to our knowledge, this
possibility was not further explored. More recently, Murphy (2013) suggested
the possibility that some slow rotators that emerge from pre-main sequence
without the presence of magnetic fields, with the fields perhaps buried,
possibly become single Am stars. Irwin & Bouvier (2009) discuss the importance
of the pre-main sequence disks in regulating the rotational period of the
stars. They suggest that the disk holds the angular velocity of the star by
removing angular momentum during the lifetime of the disk (which, in turn,
depends on stellar mass). For instance, they showed that pre-main sequence
stars in the Orion Nebula Cluster (ONC) with no detectable disk, rotate faster
than stars with disks (Herbst et al., 2001; Irwin & Bouvier, 2009). This shows
that the disk could also play a role in determining the rotation of the star.
We wonder if single Am stars could perhaps present long-lived pre-main
sequence disks, or present an efficient removal of angular momentum through
jets and outflows, resulting in the slow rotation of Am stars. These
speculative scenarios would require further investigation.

We also note that Abt (1979) found likely Am stars in young clusters with
rotational velocities (v $\sin i$) greater than $\sim$120 km/s. This would
imply that Am stars are not necessarily slow rotators. However, diffusion
models of Am stars require slow rotation, and the result of Abt (1979) is
based on few data points (Murphy, 2013; Stepień, 2000). It would be valuable
to check the result of Abt (1979) about the possible high rotation for some Am
stars.

##  5 Conclusions

We performed a detailed abundance analysis of a sample of 19 early-type stars,
7 of them hosting hot-BDs and 12 of them hot-Jupiter planets, in order to
determine if the presence of these hot-low mass companions could have an
effect on the chemical pattern of their host stars. In particular, we wanted
to determine if hot-low mass companions could play a role in the possible
development of Am stars. The main results of this work are as follows:

-The stars studied do not show a single abundance pattern. Instead, the stars analyzed show near-solar abundance values, slightly subsolar and also suprasolar abundances, including the detection of four Am stars (KELT-19A, KELT-17, HATS-70 and TOI-503) and two possible Am stars (TOI-681 and HAT-P-69).

-We detected the new Am star HATS-70 which hosts a hot-BD, and rule out this class for the hot-Jupiter host WASP-189, different than previous reports. This highlights the importance to perform a detailed chemical analysis.

-We obtained a rough estimation for the incidence of Am stars within stars hosting hot-BDs (50-75%) and within stars hosting hot-Jupiters (20-42%). In particular, the incidence of Am stars hosting hot-BDs resulted higher than the frequency of Am stars in general (31-36%). These starting trends would indicate that the presence of hot-BDs could possibly play a role in the development of Am stars, different to the case of hot-Jupiter planets. This is possibly due to the higher mass of BDs compared to the giant planets, or to a different formation mechanism. In this way, hot-BDs could possibly help to mitigate the ”single Am” problem.

-If confirmed, the initial trends found here would indicate that the search of hot-BDs could be benefited by targeting single Am stars or Am stars in wide binary systems.

The detection of additional low-mass companions around early-type stars such
as those obtained from the survey KELT, the satellite TESS and the future
mission PLATO would help to increase the significance of the interesting
results found here.

###### Acknowledgements.

We thank the anonymous referee for constructive comments that greatly improved
the paper. The authors thank Dr. R. Kurucz for making their codes available to
us. CS acknowledge financial support from FONCyT (Argentina) through grant
PICT 2017-2294 and the National University of San Juan (Argentina) through
grant CICITCA E1134. PM and JA acknowledge financial support from CONICET
(Argentina). IRAF is distributed by the National Optical Astronomical
Observatories, which is operated by the Association of Universities for
Research in Astronomy, Inc., under a cooperative agreement with the National
Science Foundation. Based on data acquired at Complejo Astronómico El
Leoncito, operated under agreement between the Consejo Nacional de
Investigaciones Científicas y Técnicas de la República Argentina and the
National Universities of La Plata, Córdoba and San Juan. This paper includes
data collected by the TESS mission that are publicly available from the
Mikulski Archive for Space Telescopes (MAST). We acknowledge the use of public
TESS data from pipelines at the TESS Science Office and at the TESS Science
Processing Operations Center. Resources supporting this work were provided by
the NASA High-End Computing (HEC) Program through the NASA Advanced
Supercomputing (NAS) Division at Ames Research Center for the production of
the SPOC data products. Funding for the TESS mission is provided by NASA’s
Science Mission Directorate. This study is based on observations obtained
through the Gemini Remote Access to CFHT ESPaDOnS Spectrograph (GRACES).
ESPaDOnS is located at the Canada–France–Hawaii Telescope (CFHT), which is
operated by the National Research Council of Canada, the Institut National des
Sciences de l Univers of the Centre National de la Recherche Scientifique of
France, and the University of Hawaii. ESPaDOnS is a collaborative project
funded by France (CNRS, MENESR, OMP, LATT), Canada (NSERC), CFHT, and ESA.
ESPaDOnS was remotely controlled from the international Gemini Observatory, a
program of NSFs NOIRLab, which is managed by the Association of Universities
for Research in Astronomy (AURA) under a cooperative agreement with the
National Science Foundation on behalf of the Gemini partnership: the National
Science Foundation (United States), the National Research Council (Canada),
Agencia Nacional de Investigación y Desarrollo (Chile), Ministerio de Ciencia,
Tecnología e Innovación (Argentina), Ministério da Ciência, Tecnologia e
Inovação (Brazil), and Korea Astronomy and Space Science Institute (Republic
of Korea). Based on observations collected with the SOPHIE spectrograph on the
1.93 m telescope at the Observatoire de Haute-Provence (CNRS), France, by the
SOPHIE Consortium.

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##  Appendix A Chemical abundances

We present in this section the chemical abundances obtained for the stars in
this work. The columns in the Tables 4 to 14 present the chemical specie,
average abundance $\pm$ total error etot, the line-to-line dispersion e1, the
error in the abundances (e2, e3, and e4) when varying Teff, $\log g$, and
vmicro by their corresponding uncertainties, and the number of lines n. For
more details, see Section 3.

Table 4: Chemical abundances for the star CoRoT-3. Specie | [X/H] $\pm$ etot | e1 | e2 | e3 | e4 | n
---|---|---|---|---|---|---
C I | 0.02 $\pm$ 0.09 | 0.03 | 0.07 | 0.05 | 0.01 | 3
Na I | 0.25 $\pm$ 0.15 | 0.11 | 0.06 | 0.02 | 0.08 | 1
Mg I | 0.05 $\pm$ 0.13 | 0.06 | 0.09 | 0.07 | 0.04 | 3
Mg II | 0.16 $\pm$ 0.19 | 0.11 | 0.09 | 0.04 | 0.12 | 1
Al I | -0.48 $\pm$ 0.18 | 0.11 | 0.10 | 0.05 | 0.08 | 1
Si II | 0.01 $\pm$ 0.15 | 0.10 | 0.09 | 0.06 | 0.05 | 2
Ca I | 0.10 $\pm$ 0.21 | 0.11 | 0.16 | 0.07 | 0.03 | 1
Ca II | 0.04 $\pm$ 0.13 | 0.11 | 0.04 | 0.06 | 0.03 | 1
Sc II | -0.09 $\pm$ 0.19 | 0.05 | 0.03 | 0.05 | 0.17 | 3
Ti II | 0.01 $\pm$ 0.15 | 0.03 | 0.06 | 0.04 | 0.13 | 22
Cr II | 0.01 $\pm$ 0.10 | 0.03 | 0.02 | 0.04 | 0.09 | 14
Mn I | 0.11 $\pm$ 0.14 | 0.05 | 0.08 | 0.01 | 0.10 | 10
Fe I | 0.01 $\pm$ 0.17 | 0.01 | 0.10 | 0.03 | 0.13 | 67
Fe II | 0.01 $\pm$ 0.14 | 0.02 | 0.07 | 0.01 | 0.12 | 26
Co I | 0.09 $\pm$ 0.23 | 0.11 | 0.17 | 0.02 | 0.11 | 1
Zn I | -0.18 $\pm$ 0.12 | 0.04 | 0.07 | 0.01 | 0.08 | 2
Sr II | 0.29 $\pm$ 0.14 | 0.01 | 0.12 | 0.01 | 0.08 | 2
Y II | 0.05 $\pm$ 0.09 | 0.01 | 0.03 | 0.05 | 0.07 | 3
Zr II | 0.18 $\pm$ 0.07 | 0.03 | 0.01 | 0.06 | 0.03 | 2
Ba II | 0.37 $\pm$ 0.26 | 0.14 | 0.09 | 0.02 | 0.20 | 3
Table 5: Chemical abundances for the star KELT-1. Specie | [X/H] $\pm$ etot | e1 | e2 | e3 | e4 | n
---|---|---|---|---|---|---
C I | 0.03 $\pm$ 0.12 | 0.01 | 0.03 | 0.09 | 0.08 | 2
Na I | 0.41 $\pm$ 0.12 | 0.04 | 0.08 | 0.05 | 0.06 | 2
Mg I | 0.16 $\pm$ 0.18 | 0.06 | 0.13 | 0.10 | 0.03 | 3
Al I | -0.17 $\pm$ 0.21 | 0.12 | 0.15 | 0.09 | 0.03 | 1
Si II | 0.26 $\pm$ 0.23 | 0.12 | 0.18 | 0.06 | 0.03 | 1
Ca I | 0.20 $\pm$ 0.29 | 0.12 | 0.23 | 0.11 | 0.05 | 1
Ca II | 0.37 $\pm$ 0.22 | 0.12 | 0.14 | 0.11 | 0.01 | 1
Sc II | -0.13 $\pm$ 0.17 | 0.12 | 0.01 | 0.09 | 0.07 | 1
Ti II | 0.09 $\pm$ 0.13 | 0.06 | 0.01 | 0.06 | 0.10 | 11
Cr II | 0.13 $\pm$ 0.08 | 0.05 | 0.01 | 0.04 | 0.05 | 12
Mn I | 0.23 $\pm$ 0.16 | 0.04 | 0.09 | 0.03 | 0.12 | 4
Fe I | 0.16 $\pm$ 0.22 | 0.02 | 0.17 | 0.06 | 0.13 | 51
Fe II | 0.15 $\pm$ 0.18 | 0.02 | 0.12 | 0.02 | 0.12 | 24
Zn I | -0.26 $\pm$ 0.13 | 0.04 | 0.12 | 0.01 | 0.04 | 2
Sr II | 0.29 $\pm$ 0.20 | 0.04 | 0.19 | 0.01 | 0.07 | 2
Y II | 0.06 $\pm$ 0.23 | 0.14 | 0.11 | 0.11 | 0.10 | 2
Ba II | 0.33 $\pm$ 0.16 | 0.04 | 0.07 | 0.04 | 0.14 | 2
Table 6: Chemical abundances for the star HATS-70. Specie | [X/H] $\pm$ etot | e1 | e2 | e3 | e4 | n
---|---|---|---|---|---|---
O I | 0.70 $\pm$ 0.22 | 0.03 | 0.02 | 0.01 | 0.22 | 3
Na I | 0.09 $\pm$ 0.15 | 0.13 | 0.05 | 0.01 | 0.04 | 1
Mg I | -0.09 $\pm$ 0.23 | 0.16 | 0.06 | 0.02 | 0.16 | 3
Mg II | 0.03 $\pm$ 0.19 | 0.13 | 0.09 | 0.04 | 0.10 | 1
Al I | 0.03 $\pm$ 0.23 | 0.13 | 0.12 | 0.01 | 0.14 | 1
Si II | 0.25 $\pm$ 0.14 | 0.08 | 0.09 | 0.04 | 0.06 | 4
Ca I | -0.37 $\pm$ 0.25 | 0.13 | 0.15 | 0.01 | 0.14 | 1
Ca II | -0.29 $\pm$ 0.15 | 0.13 | 0.06 | 0.02 | 0.01 | 1
Sc II | -0.56 $\pm$ 0.18 | 0.13 | 0.04 | 0.05 | 0.11 | 1
Ti II | 0.42 $\pm$ 0.18 | 0.04 | 0.03 | 0.04 | 0.16 | 24
Cr II | 0.51 $\pm$ 0.14 | 0.05 | 0.01 | 0.04 | 0.12 | 18
Mn I | -0.13 $\pm$ 0.15 | 0.10 | 0.07 | 0.01 | 0.08 | 4
Fe I | 0.18 $\pm$ 0.10 | 0.02 | 0.07 | 0.01 | 0.07 | 34
Fe II | 0.22 $\pm$ 0.11 | 0.04 | 0.05 | 0.02 | 0.09 | 29
Ni II | 0.50 $\pm$ 0.14 | 0.13 | 0.03 | 0.02 | 0.02 | 1
Zn I | 0.50 $\pm$ 0.09 | 0.04 | 0.07 | 0.01 | 0.03 | 3
Sr II | 1.08 $\pm$ 0.22 | 0.12 | 0.11 | 0.02 | 0.15 | 2
Y II | 1.02 $\pm$ 0.14 | 0.08 | 0.03 | 0.04 | 0.11 | 5
Zr II | 0.65 $\pm$ 0.18 | 0.17 | 0.03 | 0.04 | 0.01 | 3
Ba II | 1.90 $\pm$ 0.28 | 0.17 | 0.11 | 0.01 | 0.19 | 4
Nd II | 1.17 $\pm$ 0.13 | 0.07 | 0.10 | 0.03 | 0.01 | 3
Table 7: Chemical abundances for the star MASCARA-4. Specie | [X/H] $\pm$ etot | e1 | e2 | e3 | e4 | n
---|---|---|---|---|---|---
C I | -0.33 $\pm$ 0.04 | 0.03 | 0.01 | 0.02 | 0.01 | 4
O I | 0.04 $\pm$ 0.25 | 0.03 | 0.02 | 0.01 | 0.24 | 3
Na I | -0.13 $\pm$ 0.08 | 0.07 | 0.04 | 0.01 | 0.03 | 2
Mg I | -0.30 $\pm$ 0.23 | 0.09 | 0.04 | 0.02 | 0.21 | 3
Mg II | -0.10 $\pm$ 0.18 | 0.10 | 0.09 | 0.04 | 0.11 | 1
Al I | -0.64 $\pm$ 0.29 | 0.25 | 0.06 | 0.01 | 0.14 | 2
Si II | -0.02 $\pm$ 0.13 | 0.09 | 0.08 | 0.04 | 0.04 | 4
Ca I | -0.31 $\pm$ 0.22 | 0.10 | 0.12 | 0.01 | 0.15 | 1
Sc II | -0.26 $\pm$ 0.20 | 0.01 | 0.03 | 0.04 | 0.19 | 2
Ti II | -0.30 $\pm$ 0.14 | 0.03 | 0.03 | 0.04 | 0.13 | 27
Cr II | -0.17 $\pm$ 0.06 | 0.02 | 0.01 | 0.03 | 0.05 | 21
Mn I | -0.26 $\pm$ 0.13 | 0.07 | 0.08 | 0.01 | 0.08 | 8
Fe I | -0.26 $\pm$ 0.16 | 0.01 | 0.07 | 0.01 | 0.14 | 67
Fe II | -0.17 $\pm$ 0.12 | 0.02 | 0.04 | 0.02 | 0.11 | 37
Ni II | 0.11 $\pm$ 0.10 | 0.10 | 0.01 | 0.02 | 0.01 | 1
Zn I | -0.02 $\pm$ 0.09 | 0.07 | 0.06 | 0.01 | 0.02 | 2
Sr II | 0.42 $\pm$ 0.20 | 0.08 | 0.09 | 0.04 | 0.16 | 2
Y II | 0.34 $\pm$ 0.08 | 0.05 | 0.04 | 0.03 | 0.04 | 7
Zr II | 0.39 $\pm$ 0.11 | 0.09 | 0.03 | 0.04 | 0.02 | 5
Ba II | 0.53 $\pm$ 0.27 | 0.10 | 0.10 | 0.02 | 0.23 | 4
Eu II | 0.64 $\pm$ 0.14 | 0.10 | 0.07 | 0.07 | 0.01 | 1
Table 8: Chemical abundances for the star KELT-25. Specie | [X/H] $\pm$ etot | e1 | e2 | e3 | e4 | n
---|---|---|---|---|---|---
C I | -0.11 $\pm$ 0.12 | 0.11 | 0.05 | 0.02 | 0.01 | 3
O I | 0.46 $\pm$ 0.22 | 0.07 | 0.01 | 0.01 | 0.21 | 3
Mg I | -0.09 $\pm$ 0.22 | 0.09 | 0.10 | 0.04 | 0.17 | 3
Mg II | 0.02 $\pm$ 0.14 | 0.13 | 0.06 | 0.04 | 0.01 | 2
Al I | -0.05 $\pm$ 0.21 | 0.18 | 0.05 | 0.04 | 0.09 | 2
Si II | 0.06 $\pm$ 0.20 | 0.17 | 0.08 | 0.05 | 0.01 | 3
Ca I | -0.23 $\pm$ 0.29 | 0.14 | 0.16 | 0.02 | 0.19 | 1
Ca II | -0.05 $\pm$ 0.14 | 0.01 | 0.12 | 0.01 | 0.07 | 2
Sc II | -0.18 $\pm$ 0.20 | 0.02 | 0.08 | 0.06 | 0.17 | 2
Ti II | 0.03 $\pm$ 0.19 | 0.06 | 0.04 | 0.06 | 0.17 | 13
Cr II | 0.03 $\pm$ 0.12 | 0.06 | 0.03 | 0.05 | 0.09 | 10
Mn I | -0.18 $\pm$ 0.21 | 0.17 | 0.12 | 0.01 | 0.04 | 2
Fe I | -0.14 $\pm$ 0.17 | 0.02 | 0.12 | 0.01 | 0.12 | 31
Fe II | -0.05 $\pm$ 0.15 | 0.03 | 0.08 | 0.03 | 0.12 | 23
Sr II | -0.18 $\pm$ 0.24 | 0.09 | 0.12 | 0.09 | 0.16 | 2
Y II | 0.23 $\pm$ 0.11 | 0.06 | 0.04 | 0.07 | 0.04 | 3
Zr II | 0.15 $\pm$ 0.19 | 0.14 | 0.12 | 0.04 | 0.02 | 1
Ba II | 0.37 $\pm$ 0.24 | 0.14 | 0.12 | 0.04 | 0.15 | 2
Table 9: Chemical abundances for the star HAT-P-70. Specie | [X/H] $\pm$ etot | e1 | e2 | e3 | e4 | n
---|---|---|---|---|---|---
C I | -0.34 $\pm$ 0.12 | 0.09 | 0.07 | 0.02 | 0.01 | 2
O I | 0.46 $\pm$ 0.20 | 0.05 | 0.04 | 0.01 | 0.19 | 3
Mg I | -0.33 $\pm$ 0.23 | 0.10 | 0.13 | 0.04 | 0.16 | 3
Al I | -0.56 $\pm$ 0.22 | 0.17 | 0.10 | 0.01 | 0.11 | 2
Si II | -0.26 $\pm$ 0.22 | 0.12 | 0.10 | 0.05 | 0.15 | 3
Ca I | -0.35 $\pm$ 0.26 | 0.14 | 0.19 | 0.01 | 0.12 | 1
Ca II | -0.24 $\pm$ 0.17 | 0.06 | 0.15 | 0.01 | 0.07 | 2
Sc II | -0.44 $\pm$ 0.15 | 0.04 | 0.10 | 0.04 | 0.09 | 2
Ti II | -0.10 $\pm$ 0.17 | 0.04 | 0.08 | 0.05 | 0.13 | 18
Cr II | -0.05 $\pm$ 0.15 | 0.10 | 0.03 | 0.03 | 0.09 | 6
Fe I | -0.39 $\pm$ 0.19 | 0.02 | 0.15 | 0.01 | 0.11 | 38
Fe II | -0.29 $\pm$ 0.15 | 0.03 | 0.11 | 0.02 | 0.10 | 28
Sr II | -0.88 $\pm$ 0.26 | 0.19 | 0.11 | 0.04 | 0.13 | 2
Zr II | -0.03 $\pm$ 0.19 | 0.14 | 0.12 | 0.02 | 0.01 | 1
Ba II | -0.04 $\pm$ 0.24 | 0.14 | 0.19 | 0.01 | 0.04 | 1
Table 10: Chemical abundances for the star TOI-503. Specie | [X/H] $\pm$ etot | e1 | e2 | e3 | e4 | n
---|---|---|---|---|---|---
Li I | 1.89 $\pm$ 0.11 | 0.09 | 0.06 | 0.01 | 0.01 | 1
C I | -0.70 $\pm$ 0.09 | 0.09 | 0.01 | 0.02 | 0.01 | 1
O I | -0.45 $\pm$ 0.19 | 0.01 | 0.03 | 0.02 | 0.19 | 3
Na I | 0.41 $\pm$ 0.08 | 0.04 | 0.04 | 0.01 | 0.06 | 2
Mg I | -0.14 $\pm$ 0.18 | 0.11 | 0.05 | 0.01 | 0.13 | 2
Mg II | -0.28 $\pm$ 0.17 | 0.09 | 0.09 | 0.04 | 0.11 | 1
Si II | 0.16 $\pm$ 0.14 | 0.09 | 0.08 | 0.04 | 0.06 | 1
Ca II | -0.45 $\pm$ 0.21 | 0.07 | 0.09 | 0.01 | 0.19 | 2
Sc II | -0.49 $\pm$ 0.10 | 0.09 | 0.02 | 0.03 | 0.01 | 1
Ti II | 0.05 $\pm$ 0.15 | 0.05 | 0.02 | 0.03 | 0.14 | 14
Cr II | 0.46 $\pm$ 0.11 | 0.02 | 0.01 | 0.03 | 0.10 | 19
Mn I | 0.23 $\pm$ 0.08 | 0.05 | 0.05 | 0.01 | 0.04 | 3
Fe I | 0.21 $\pm$ 0.19 | 0.02 | 0.05 | 0.01 | 0.18 | 34
Fe II | 0.23 $\pm$ 0.13 | 0.02 | 0.03 | 0.02 | 0.13 | 29
Zn I | 0.43 $\pm$ 0.08 | 0.02 | 0.05 | 0.01 | 0.06 | 2
Sr II | 0.20 $\pm$ 0.23 | 0.09 | 0.09 | 0.07 | 0.18 | 1
Y II | 0.57 $\pm$ 0.10 | 0.08 | 0.03 | 0.03 | 0.03 | 4
Ba II | 1.33 $\pm$ 0.16 | 0.08 | 0.08 | 0.01 | 0.11 | 4
Table 11: Chemical abundances for the star TOI-681. Specie | [X/H] $\pm$ etot | e1 | e2 | e3 | e4 | n
---|---|---|---|---|---|---
C I | -0.05 $\pm$ 0.13 | 0.13 | 0.01 | 0.02 | 0.01 | 1
Na I | 0.51 $\pm$ 0.12 | 0.07 | 0.05 | 0.02 | 0.08 | 2
Mg I | -0.29 $\pm$ 0.19 | 0.05 | 0.08 | 0.04 | 0.15 | 3
Mg II | -0.20 $\pm$ 0.19 | 0.13 | 0.07 | 0.02 | 0.12 | 1
Si II | 0.42 $\pm$ 0.17 | 0.13 | 0.09 | 0.04 | 0.05 | 1
Ca I | -0.37 $\pm$ 0.28 | 0.13 | 0.20 | 0.04 | 0.15 | 1
Sc II | 0.13 $\pm$ 0.15 | 0.01 | 0.08 | 0.02 | 0.13 | 3
Ti II | 0.23 $\pm$ 0.18 | 0.05 | 0.05 | 0.04 | 0.16 | 15
Cr II | 0.35 $\pm$ 0.08 | 0.03 | 0.01 | 0.03 | 0.06 | 14
Mn I | 0.02 $\pm$ 0.12 | 0.07 | 0.07 | 0.03 | 0.06 | 4
Fe I | 0.05 $\pm$ 0.20 | 0.02 | 0.11 | 0.02 | 0.17 | 35
Fe II | 0.07 $\pm$ 0.14 | 0.03 | 0.07 | 0.01 | 0.12 | 15
Zn I | 0.20 $\pm$ 0.10 | 0.01 | 0.09 | 0.03 | 0.05 | 2
Sr II | 0.33 $\pm$ 0.27 | 0.19 | 0.15 | 0.03 | 0.13 | 2
Y II | 0.30 $\pm$ 0.07 | 0.04 | 0.01 | 0.05 | 0.03 | 3
Ba II | 1.13 $\pm$ 0.18 | 0.10 | 0.11 | 0.05 | 0.09 | 3
Table 12: Chemical abundances for the star TOI-587. Specie | [X/H] $\pm$ etot | e1 | e2 | e3 | e4 | n
---|---|---|---|---|---|---
O I | -0.38 $\pm$ 0.05 | 0.01 | 0.04 | 0.02 | 0.01 | 2
Mg I | -0.48 $\pm$ 0.24 | 0.09 | 0.21 | 0.07 | 0.02 | 3
Mg II | -0.58 $\pm$ 0.09 | 0.09 | 0.01 | 0.01 | 0.01 | 3
Al II | 0.07 $\pm$ 0.11 | 0.08 | 0.06 | 0.05 | 0.01 | 1
Si II | -0.20 $\pm$ 0.06 | 0.04 | 0.05 | 0.01 | 0.02 | 4
Ca I | -0.72 $\pm$ 0.20 | 0.08 | 0.17 | 0.06 | 0.01 | 1
Sc II | -0.77 $\pm$ 0.23 | 0.21 | 0.10 | 0.01 | 0.01 | 2
Ti II | -0.63 $\pm$ 0.14 | 0.02 | 0.13 | 0.01 | 0.02 | 12
Cr II | -0.22 $\pm$ 0.08 | 0.05 | 0.05 | 0.02 | 0.01 | 11
Fe I | -0.16 $\pm$ 0.17 | 0.02 | 0.17 | 0.02 | 0.03 | 18
Fe II | -0.13 $\pm$ 0.07 | 0.01 | 0.05 | 0.02 | 0.04 | 29
Ni II | 0.21 $\pm$ 0.10 | 0.08 | 0.05 | 0.02 | 0.03 | 1
Sr II | 0.31 $\pm$ 0.24 | 0.14 | 0.13 | 0.02 | 0.15 | 2
Ba II | 0.87 $\pm$ 0.29 | 0.18 | 0.23 | 0.04 | 0.01 | 2
Table 13: Chemical abundances for the star KELT-19A. Specie | [X/H] $\pm$ etot | e1 | e2 | e3 | e4 | n
---|---|---|---|---|---|---
O I | -0.57 $\pm$ 0.25 | 0.20 | 0.04 | 0.05 | 0.14 | 1
Na I | 0.51 $\pm$ 0.22 | 0.20 | 0.01 | 0.01 | 0.10 | 1
Mg I | -0.10 $\pm$ 0.18 | 0.02 | 0.04 | 0.08 | 0.15 | 3
Mg II | -0.18 $\pm$ 0.24 | 0.20 | 0.05 | 0.03 | 0.13 | 1
Si II | 0.37 $\pm$ 0.22 | 0.20 | 0.07 | 0.05 | 0.04 | 1
Ca II | -0.55 $\pm$ 0.18 | 0.07 | 0.06 | 0.01 | 0.15 | 2
Sc II | -0.79 $\pm$ 0.16 | 0.09 | 0.01 | 0.07 | 0.10 | 2
Ti II | 0.39 $\pm$ 0.18 | 0.05 | 0.02 | 0.06 | 0.16 | 9
Cr II | 0.62 $\pm$ 0.11 | 0.03 | 0.01 | 0.03 | 0.10 | 13
Mn I | 0.43 $\pm$ 0.21 | 0.20 | 0.03 | 0.02 | 0.05 | 1
Fe I | 0.39 $\pm$ 0.22 | 0.04 | 0.06 | 0.03 | 0.21 | 20
Fe II | 0.41 $\pm$ 0.17 | 0.04 | 0.02 | 0.01 | 0.16 | 14
Zn I | 0.34 $\pm$ 0.11 | 0.09 | 0.04 | 0.01 | 0.05 | 2
Sr II | 0.58 $\pm$ 0.28 | 0.20 | 0.11 | 0.05 | 0.15 | 1
Y II | 1.23 $\pm$ 0.26 | 0.12 | 0.03 | 0.14 | 0.18 | 3
Ba II | 1.46 $\pm$ 0.19 | 0.11 | 0.06 | 0.03 | 0.14 | 2
Table 14: Chemical abundances for the star HAT-P-69. Specie | [X/H] $\pm$ etot | e1 | e2 | e3 | e4 | n
---|---|---|---|---|---|---
O I | -0.40 $\pm$ 0.18 | 0.04 | 0.04 | 0.04 | 0.17 | 2
Na I | 0.36 $\pm$ 0.09 | 0.03 | 0.06 | 0.02 | 0.05 | 2
Mg I | -0.11 $\pm$ 0.21 | 0.09 | 0.09 | 0.01 | 0.16 | 3
Mg II | -0.15 $\pm$ 0.17 | 0.12 | 0.04 | 0.04 | 0.11 | 1
Si II | 0.22 $\pm$ 0.19 | 0.12 | 0.14 | 0.02 | 0.05 | 1
Ca II | -0.05 $\pm$ 0.22 | 0.12 | 0.15 | 0.01 | 0.12 | 2
Sc II | -0.45 $\pm$ 0.13 | 0.12 | 0.02 | 0.03 | 0.01 | 1
Ti II | 0.06 $\pm$ 0.15 | 0.04 | 0.01 | 0.03 | 0.14 | 11
Cr II | 0.17 $\pm$ 0.11 | 0.02 | 0.04 | 0.02 | 0.10 | 9
Mn I | 0.21 $\pm$ 0.09 | 0.07 | 0.06 | 0.01 | 0.01 | 2
Fe I | 0.06 $\pm$ 0.22 | 0.02 | 0.10 | 0.01 | 0.19 | 27
Fe II | 0.07 $\pm$ 0.15 | 0.03 | 0.05 | 0.03 | 0.14 | 23
Zn I | 0.00 $\pm$ 0.15 | 0.12 | 0.08 | 0.01 | 0.01 | 1
Sr II | 0.71 $\pm$ 0.24 | 0.12 | 0.09 | 0.04 | 0.19 | 1
Y II | 0.91 $\pm$ 0.20 | 0.03 | 0.07 | 0.06 | 0.17 | 3
Zr II | 0.75 $\pm$ 0.22 | 0.19 | 0.11 | 0.01 | 0.05 | 3
Ba II | 0.91 $\pm$ 0.21 | 0.13 | 0.14 | 0.03 | 0.09 | 2