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: Consejo Nacional de Investigaciones
Científicas y Técnicas (CONICET), Argentina 77institutetext: Universidad
Nacional de San Juan (UNSJ), Facultad de Filosofía, Humanidades y Artes
(FFHA), San Juan, Argentina

# Testing the accretion scenario of $\lambda$ Boo stars

J. Alacoria 1166 C. Saffe 112266 M. Jaque Arancibia 3344 R. Angeloni 55 P.
Miquelarena 112266 M. Flores 112266 M. E. Veramendi 116677 A. Collado 112266

(Received xxx, xxx ; accepted xxxx, xxxx)

###### Abstract

Context. The group of $\lambda$ Boo stars is known for years, however the
origin of its chemical peculiarity is still strongly debated.

Aims. Our aim is to test the accretion scenario of $\lambda$ Boo stars. This
model predicts that a binary system with two early-type stars passing through
a diffuse cloud should both display the same superficial peculiarity.

Methods. Via spectral synthesis, we carried out a detailed abundance
determination of three multiple systems hosting a candidate $\lambda$ Boo
star: the remarkable triple system HD 15164/65/65C and the two binary systems
HD 193256/281 and HD 198160/161. Stellar parameters were initially estimated
using Strömgren photometry or literature values and then refined by requiring
excitation and ionization balances of Fe lines. The abundances were determined
iteratively for 24 different species by fitting synthetic spectra using the
SYNTHE program together with local thermodynamic equilibrium (LTE) ATLAS12
model atmospheres. Specific opacities were calculated for each star, depending
on its arbitrary composition and microturbulence velocity vmicro through the
opacity sampling (OS) method. The abundances of the light elements C and O
were corrected by Non-LTE effects. The complete chemical pattern of the stars
were then compared to those of $\lambda$ Boo stars.

Results.  The abundance analysis of the triple system HD 15164/65/65C shows a
clear $\lambda$ Boo object (HD 15165) and two objects with near solar
composition (HD 15164 and 15165C). Notably, the presence of a $\lambda$ Boo
star (HD 15165) together with a near solar early-type object (HD 15164) is
difficult to explain under the accretion scenario. Also, the solar-like
composition derived for the late-type star of the system (HD 15165C) could be
used, for the first time, as a proxy for the initial composition of the
$\lambda$ Boo stars. This could help to constrain any model of $\lambda$ Boo
stars formation and not only the accretion scenario. The abundance analysis of
the binary system HD 193256/281 shows no clear $\lambda$ Boo components, while
the analysis of HD 198160/161 shows two mild-$\lambda$ Boo stars. Then, by
carefully reviewing abundance analysis of all known binary systems with
candidate $\lambda$ Boo stars from literature and including the systems
analyzed here, we find no binary/multiple system having two clear ”bonafide”
$\lambda$ Boo stars, as expected from the accretion scenario. The closer
candidates to show two $\lambda$ Boo-like stars are the binary systems HD
84948, HD 171948 and HD 198160; however, in our opinion they show mild rather
than clear $\lambda$ Boo patterns.

Conclusions. We performed for the first time a complete analysis of a triple
system which includes a $\lambda$ Boo candidate. Our results brings little
support to the accretion scenario of $\lambda$ Boo stars. Then, there is an
urgent need of additional binary and multiple systems to be analyzed through a
detailed abundance analysis, in order to test the accretion model of $\lambda$
Boo stars.

###### Key Words.:

Stars: abundances – Stars: binaries – Stars: chemically peculiar – Stars:
individual: HD 15164, HD 15165, HD 15165C, HD 193256, HD 193281, HD 198160, HD
198161

##  1 Introduction

The main feature of $\lambda$ Boo stars is a notable underabundance of most
Fe-peak elements and near solar abundances of lighter elements (C, N, O and
S). They comprise main-sequence late-B to early-F stars, where a maximum of
about 2% of all objects are believed to be $\lambda$ Boo stars (Gray &
Corbally, 1998; Paunzen et al., 2001b). Classification-resolution spectroscopy
shows promising $\lambda$ Boo candidates (e.g., Murphy et al., 2015; Gray et
al., 2017), and a more detailed abundance determination, especially including
the lighter elements, is considered a ultimate test to confirm that a
candidate is indeed a bonafide member of the class (e.g., Andrievsky et al.,
2002; Heiter et al., 2002).

The origin of the $\lambda$ Boo peculiarity still remains as a challenge (see,
e.g., the recent discussion of Murphy & Paunzen, 2017, and references
therein). Their rotational velocities do not necessarily point toward lower
values, marking thus a difference with chemically peculiar Am and Ap stars
(Abt & Morrell, 1995; Murphy et al., 2015). A possible explanation consist in
the interaction of the star with a diffuse interstellar cloud (Kamp & Paunzen,
2002; Martinez-Galarza et al., 2009). In this work, we refer to this model as
the ”accretion scenario”, in which the underabundances are produced by
different amounts of volatile accreted material, and the more refractory
species are possibly separated and repelled from the star. More recently, Jura
(2015) proposed that this peculiar pattern possibly originates from the winds
of hot-Jupiter planets111Hot-Jupiter planets present short orbital periods
($<$10 d) and large masses ($>$ 0.1 MJup).. In this case, the planet acts as
the source of gas poor in refractory species. However, Saffe et al. (2021)
have recently shown that eight early-type stars hosting hot-Jupiter planets do
not display the $\lambda$ Boo peculiarity. This would let the interaction of
the star with a diffuse cloud, as the more plausible scenario to explain the
$\lambda$ Boo phenomena in main-sequence stars.

Under the accretion scenario, two early-type stars passing through a diffuse
cloud should display, in principle, the same superficial peculiarity (e.g.,
Paunzen et al., 2012a, b). At the same time, hotter stars (Teff $>$
$\sim$12000 K) with strong winds, and cooler stars (Teff $<$ $\sim$6500 K)
with larger convective zones, should not notably change their composition.
These predictions make the analysis of binary and multiple systems an
important tool to test the accretion scenario. However, the number of known
candidate $\lambda$ Boo stars in binary/multiple systems is limited to a dozen
of objects (e.g., Paunzen et al., 2012a, b), where most of them are
spectroscopic binary (SB) systems. To our knowledge, only five of these
systems present a detailed chemical analysis of the two components (see the
Appendix for a more detailed review). Notably, some stars of these binary
systems were recently identified as non-members or uncertain members of the
$\lambda$ Boo class (see Gray et al., 2017). Based on literature data, we
selected for this study three binary/multiple systems that possibly confront
the accretion scenario. In addition, they are spatially resolved (in contrast
to most candidate $\lambda$ Boo stars that belong to SB systems, Paunzen et
al., 2012a, b), allowing a individual analysis without a strong contribution
from the companion. We also review all known binary or multiple systems with
candidate $\lambda$ Boo stars, with data taken from the literature (see
Appendix).

In this work, we present an analysis of the remarkable triple system HD 15165.
It is composed by HD 15165, HD 15164 and HD 15165C (stars A, B and C) with
spectral types ”F2 V kA2mA2 $\lambda$ Boo?”, ”F1 V kA7mA6 ($\lambda$ Boo)?”
and ”K2 V” (Murphy et al., 2015). Some previous works suggest that the A star
belong to the $\lambda$ Boo class (Andrievsky et al., 1995; Chernyshova et
al., 1998), while the B star seem to display, notably, a solar composition
(Andrievsky et al., 1995). If these abundances are confirmed, this could
seriously defy the accretion scenario. In addition, currently there is no
analysis of the 3rd star, the late-type component of the system. Therefore, we
take the opportunity and perform a detailed abundance analysis including for
the first time the three stars of the system, using a spectra with higher
resolving power than previous works.

We also present an analysis of the binary systems HD 193256/281 and HD
198160/161. Both systems show solar values for C and subsolar Fe, similar to
other candidate $\lambda$ Boo stars (Stürenburg, 1993). However, more recent
classification spectra suggest that only one star of the system belong to the
$\lambda$ Boo class (see Tables 1 and 4 of Murphy et al., 2015; Gray et al.,
2017), which would be difficult to explain under the accretion scenario. This
convert both systems in very interesting targets to study in detail, and are
included in our analysis.

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

##  2 Observations

We present in Table 1 the visual magnitude V (from Hipparcos), coordinates,
proper motions and parallax (from Gaia DR2, Gaia Collaboration, 2018) for the
stars studied in this work. Spectral data of the triple system HD 15165 were
obtained with the Max Planck Gesselschaft (MPG) 2.2 meter telescope at the
European Southern Observatory (ESO) in La Silla, Chile, on October 10, 2021
(Program ID: 0108.A-9012, PI: Marcelo Jaque Arancibia). We used the Fiber-fed
Extended Range Optical Spectrograph (FEROS), which provides a high-resolution
(R$\sim$48.000) spectra when illuminated via the 2.0 arcsec aperture on the
sky in the unbinned mode. Three individual spectra for each object were
obtained, followed by a ThAr lamp in order to obtain an appropriate wavelength
solution. The data were reduced using the FEROS Data Reduction
System222https://www.eso.org/sci/facilities/lasilla/instruments/feros/tools/DRS.html
(DRS). The spectral coverage resulted between 3700-9000 Å, approximately, and
the S/N per pixel measured at $\sim$5000 Å resulted in $\sim$300.

Table 1: Magnitudes and astrometric data for the stars studied in this work. Star | V | $\alpha$ | $\delta$ | $\mu_{\alpha}$ | $\mu_{\delta}$ | $\pi$ | Spectra
---|---|---|---|---|---|---|---
|  | J2000 | J2000 | [mas/yr] | [mas/yr] | [mas] |
HD 15164 | 8.27 | 02 26 48.29 | +10 34 57.59 | 36.552 | -13.717 | 7.4185 | MPG+FEROS
HD 15165 | 6.69 | 02 26 45.65 | +10 33 55.07 | 36.680 | -13.086 | 7.4414 | MPG+FEROS
HD 15165C | 11.78 | 02 26 47.40 | +10 32 58.89 | 36.805 | -13.131 | 7.5499 | MPG+FEROS
HD 193256 | 7.53 | 20 20 26.57 | -29 11 28.76 | -1.991 | -1.221 | 5.8675 | CASLEO+REOSC
HD 193281 | 6.64 | 20 20 27.88 | -29 11 49.97 | -0.653 | 0.244 | 6.2644 | CASLEO+REOSC
HD 198160 | 6.21 | 20 51 38.51 | -62 25 45.59 | 82.697 | -46.562 | 13.5137 | MPG+FEROS
HD 198161 | 6.56 | 20 51 38.85 | -62 25 45.26 | 82.077 | -42.340 | 13.5315 | MPG+FEROS

The spectra of the binary system HD 193256/281 were obtained at the Complejo
Astrónomico El Leoncito (CASLEO) between May 9 and 11, 2009 (PI: Maria Eugenia
Veramendi). We used the _Jorge Sahade_ 2.15-m telescope equipped with a REOSC
echelle spectrograph333On 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 $\sim$ 12500 covering the spectral range
$\lambda\lambda$3800–6500. Three individual spectra for each object were
obtained and then combined, reaching a final S/N per pixel of $\sim$300
measured at $\sim$5000 Å. The data were reduced with Image Reduction and
Analysis Facility (IRAF) following the standard recipe for echelle spectra
(i.e. bias and flat corrections, order-by-order normalization, scattered light
correction, etc.).

Finally, the FEROS spectra of the binary system HD 198160/161 were obtained
from the ESO Science Archive Facility444http://archive.eso.org/cms.html. The
stars were observed between April 4 and 7, 2017 (Program ID: 099-A-9029). The
spectra were reduced using the FEROS DRS, obtaining a spectral coverage and
S/N similar to those obtained with HD 15165.

##  3 Stellar parameters and abundance analysis

The stellar parameters Teff and $\log g$ were estimated iteratively, similar
to previous works (Saffe et al., 2021). They were first estimated by using the
Strömgren uvby$\beta$ mean photometry of Hauck & Mermilliod (1998) or
considering previously published results. We used the calibration of
Napiwotzki et al. (1993) and deredenned colors according to Domingo & Figueras
(1999) within the program TempLogG (Kaiser, 2006), in order to derive the
fundamental parameters. These initial values were refined (when necessary
and/or possible) by imposing excitation and ionization balances of the iron
lines. A similar strategy was previously applied in the literature (e.g.,
Saffe & Levato, 2014; Saffe et al., 2021). The values derived in this way are
listed in the Table 2, with an average dispersion of $\sim$115 K and
$\sim$0.13 dex for Teff and $\log g$, respectively.

Table 2: Fundamental parameters derived for the stars in this work. Star | Teff | $\log g$ | vmicro | $v\sin i$
---|---|---|---|---
| [K] | [dex] | [km s-1] | [km s-1]
HD 15164 | 7150 $\pm$ 70 | 3.74 $\pm$ 0.08 | 2.54 $\pm$ 0.63 | 17.9 $\pm$ 0.7
HD 15165 | 6950 $\pm$ 139 | 3.80 $\pm$ 0.19 | 2.21 $\pm$ 0.55 | 125.7 $\pm$ 5.4
HD 15165C | 4960 $\pm$ 51 | 4.40 $\pm$ 0.03 | 0.46 $\pm$ 0.07 | 2.4 $\pm$ 0.3
HD 193256 | 7780 $\pm$ 146 | 3.97 $\pm$ 0.19 | 3.23 $\pm$ 0.81 | 257.0 $\pm$ 8.2
HD 193281 | 8700 $\pm$ 140 | 3.60 $\pm$ 0.15 | 2.99 $\pm$ 0.75 | 91.5 $\pm$ 3.9
HD 198160 | 8010 $\pm$ 130 | 4.09 $\pm$ 0.15 | 3.31 $\pm$ 0.83 | 190.0 $\pm$ 6.8
HD 198161 | 8010 $\pm$ 130 | 4.09 $\pm$ 0.15 | 3.31 $\pm$ 0.83 | 185.0 $\pm$ 7.2

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. Microturbulence velocity vmicro was estimated as a function
of Teff following the formula of Gebran et al. (2014), (valid for $\sim$6000 K
$<$ Teff $<$ $\sim$10000 K), except for the late-type star HD 15165C for which
we used the formula of Ramirez et al. (2013) for FGK stars. We adopt for
vmicro an uncertainty of $\sim$25 $\%$, as suggested by Gebran et al. (2014).

Chemical abundances were determined iteratively by fitting a synthetic spectra
using the program SYNTHE (Kurucz, 1993). In the first step, we use an ATLAS12
model atmosphere calculated with solar abundances. With the new abundance
values, we derived a new model atmosphere and started 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). Possible differences originated from the use of opacities
with solar-scaled composition instead of an arbitrary composition, were
recently estimated for solar-type stars (Saffe et al., 2018, 2019). If
necessary, Teff and $\log g$ were refined to achieve the balance of Fe I and
Fe II lines. In this way, abundances and parameters are consistently derived
until reach the same input and output abundance values (for more details, see
Saffe et al., 2021).

Chemical abundances were derived for 24 different species. The atomic line
list and laboratory data used in this work are the same described in Saffe et
al. (2021). In Figs. 1 and 2 we present an example of observed and synthetic
spectra (black and blue dotted lines, almost superimposed) together with the
difference spectra (magenta) for the stars in our sample. For clarity, Fig. 1
corresponds to stars with the higher $v\sin i$ values ($>$ 91 km s-1), while
Fig. 2 corresponds to stars with the lower $v\sin i$ values ($<$ 17.9 km s-1).
The stars are sorted in these plots by increasing $v\sin i$. There is a good
agreement between modeling and observations for the lines of different
chemical species. To determine the uncertainty in the abundance values, we
considered different sources. The total error etot was derived as the
quadratic sum of the line-to-line dispersion e1 (estimated as
$\sigma/\sqrt{n}$ , where $\sigma$ is the standard deviation), and the error
in the abundances (e2, e3 and e4) when varying Teff, $\log g$ and vmicro by
their corresponding uncertainties555We adopt a minimum of 0.01 dex for the
errors e2, e3 and e4.. For chemical species with only one line, we adopt as
$\sigma$ the standard deviation of iron lines. The abundances, the total error
etot and the individual errors e1 to e4 are presented in Tables B.1 to B.7 of
the Appendix.

Figure 1: Observed, synthetic, and difference spectra (black, blue dotted, and
magenta lines) for the stars in our sample, sorted by $v\sin i$. Figure 2:
Observed, synthetic, and difference spectra (black, blue dotted, and magenta
lines) for the stars in our sample, sorted by $v\sin i$.

###  3.1 NLTE effects

Light element Non-Local Thermodynamic Equilibrium (NLTE) abundances are
particularly important for the case of $\lambda$ Boo stars. For instance,
Paunzen et al. (1999) derived for a sample of $\lambda$ Boo stars an average O
I correction of -0.5 dex, while for C I they estimated an average correction
of -0.1 dex. Rentzsch-Holm (1996) calculated carbon NLTE abundance corrections
by using a multilevel model atom for stars with Teff between 7000 K and 12000
K, log g between 3.5 and 4.5 dex, and metallicities from -0.5 dex to +1.0 dex.
She showed that C I NLTE effects are negative (calculated as NLTE-LTE) and
depend basically on equivalent width Weq. Near $\sim$7000 K, the three lower
levels of C I are always in LTE; however, by increasing the Teff increase the
underpopulation of these levels respect to LTE by UV photoionization. Then, we
estimated NLTE abundance corrections of C I for the early-type stars in our
sample by interpolating in their Figs. 7 and 8 as a function of Teff, Weq and
metallicity.

Sitnova et al. (2013) performed NLTE abundance corrections for O I for stars
with spectral types from A to K (Teff between 10000 and 5000 K). They showed
that NLTE effects lead to an strengthening of O I lines, producing a negative
NLTE correction. We estimated NLTE abundance corrections of O I (IR triplet
7771 Å and 6158 Å) for the stars in this work, by interpolating in the Table
11 of Sitnova et al. (2013) as a function of Teff. Other O I lines present
corrections lower than $\sim$-0.02 dex (see, e.g., Table 5 of Sitnova et al.,
2013).

###  3.2 Comparison with literature

We present in Fig. 3 a comparison of [Fe/H] values derived in this work, with
those taken from literature for the stars HD 15164 (Andrievsky et al., 1995),
HD 15164 (Paunzen et al., 2002), HD 193256, HD 193281, HD 198160 and HD 198161
(Stürenburg, 1993). In general, there is a reasonable agreement with
literature, where the star HD 193281 present the larger difference (marked in
the plot).

Stürenburg (1993) estimated for HD 193281 an iron abundance of
[Fe/H]=-1.0$\pm$0.2. However, we estimated for this star a somewhat higher
value of [FeI/H]=-0.36$\pm$0.13 ([FeII/H]=-0.48$\pm$0.13). We explored the
possible sources for this difference. They estimated a Teff of 8080 K (without
quoting uncertainties) by using the Strömgren photometry, while we estimated
for this object a Teff of 8700$\pm$140K, having a difference of 620 K. This
could be one of the reasons for the different [Fe/H] that we obtained.
Different works estimated for this star temperatures of 8700 K (Gray et al.,
2017), 8623 K (Koleva & Vazdekis, 2012), and recently 8695 K (Arentsen et al.,
2019). Then, our estimated Teff is more in agreement with these works. We also
note that this star presents different metallicities in literature:
-1.0$\pm$0.2 dex (Stürenburg, 1993), -0.68 dex (Koleva & Vazdekis, 2012) and
more recently -0.37 dex (Arentsen et al., 2019). Our estimated metallicity of
[FeI/H]=-0.36$\pm$0.13 is closer to the work of Arentsen et al. (2019).

In addition, there is evidence that HD 193281 could be contaminated by a
nearby star. Simbad database reports that the star ADS 13702 B (= TYC
6918-1823-2) is located at $\sim$3.5 arcsec from HD 193281, having spectral
type ”F5:V”. Ivanov et al. (2019) present a library of stellar spectra taken
with the integral field spectrograph
MUSE666https://www.eso.org/sci/facilities/develop/instruments/muse.html in low
spectral resolution (R$\sim$2000) although with high spatial resolution
(0.3-0.4 arcsec). They report that HD 193281 is a binary with $\sim$3.8 arcsec
separation and the components cross-contaminate each other. They identified
the components as HD 193281 A and B, and estimated spectral types A2 III and
K2 III, respectively (updating the spectral type F5:V reported by Simbad for
the star HD 193281 B). This possible contamination could explain, at least in
part, the different parameters and metallicities obtained from different works
for this object. In this study, we estimated parameters and abundances of HD
193281 taken as single, for which the resulting values should then be
considered with caution.

Figure 3: Comparison of [Fe/H] values derived in this work with those from
literature. Average dispersion bars are showed in the upper left corner.

##  4 Discussion

In order to test the accretion scenario of $\lambda$ Boo stars, we compare the
chemical abundances of the stars in our sample with those of $\lambda$ Boo
stars. The three multiple systems with candidate $\lambda$ Boo stars are
discussed separately, while other binary or multiple systems with candidate
$\lambda$ Boo stars are discussed in the Appendix.

###  4.1 The average pattern of $\lambda$ Boo stars

To derive an average $\lambda$ Boo pattern is not an easy task. Few literature
works obtain homogeneous abundances of many species for $\lambda$ Boo stars
(e.g., Stürenburg, 1993; Andrievsky et al., 2002; Heiter et al., 2002).
Stürenburg (1993) derived abundances for 16 A-type stars classified, in
principle, as $\lambda$ Boo stars. They performed NLTE corrections for some
elements including C. However, they included stars that were subsequently
considered non-members or uncertain members, such as HD 38545 and HD 193281
(Murphy et al., 2015). Paunzen et al. (1999) and Kamp et al. (2001) derived
light-element NLTE abundances for a sample of $\lambda$ Boo stars. Then,
Andrievsky et al. (2002) derived elemental abundances for 20 candidate
$\lambda$ Boo stars basically selected from classification-resolution
spectroscopy. They performed primarily a LTE approach and included NLTE
effects for Na. They were able to confirm the membership of only nine objects
to the $\lambda$ Boo class, while other stars were ruled out or present an
unclear membership. Paunzen et al. (2002) collected abundance values for 26
candidate $\lambda$ Boo stars (see their Table 5), although using different
literature sources. Also, Heiter et al. (2002) reported LTE abundance values
for 12 candidate $\lambda$ Boo stars, four of them belonging to SB systems.
Then, it would be highly desirable a homogeneous abundance determination
including more candidate $\lambda$ Boo stars, newer laboratory data for the
lines and including NLTE effects especially for the light-elements.

In order to test the accretion scenario of $\lambda$ Boo stars, we compare the
chemical abundances of the stars in our sample with those of $\lambda$ Boo
stars. In this work, we used the data derived by Heiter et al. (2002), who
homogeneously determined abundances for a number of $\lambda$ Boo stars. We
excluded from the average those stars without CNO values and the stars
analyzed here.

###  4.2 The triple system HD 15164/65/65C

This remarkable triple system is composed by two early-type stars (HD 15165
and HD 15164, the stars A and B) and a late-type companion (HD 15165C). A
number of studies suggest that the spectrum of HD 15165 resembles that of
metal-deficient star, but the companion HD 15164 has a near solar abundance
(Mechler, 1974, 1976; Abt, 1980). Then, as explained in the Introduction, some
works suggest that the A star belong to the $\lambda$ Boo class (Andrievsky et
al., 1995; Chernyshova et al., 1998), while the B star seems to display a
solar composition (Andrievsky et al., 1995). To our knowledge, there is no
abundance determination for the C component.

We present in Fig. 4 the chemical pattern of the stars HD 15164, HD 15165 and
HD 15165C (black), compared to an average pattern of $\lambda$ Boo stars
(blue). For each star we present two panels, corresponding to elements with
atomic number z$<$32 and z$>$32\. The error bars of the $\lambda$ Boo pattern
show the standard deviation derived from different stars, while the error bars
for our stars correspond to the total error etot. As we can see in the Fig. 4
the chemical pattern of the primary (HD 15165) is similar to the pattern of
$\lambda$ Boo stars, showing subsolar abundances of most metals (Mg, Al, Ca,
Sc, Ti, Cr, Fe) together with near solar values of C and O. The abundances of
Sr and Ba present a less marked deficiency, although still showing subsolar
values. On the other hand, the chemical pattern of the secondary star (HD
15164) shows a slight deficiency in some metals (for instance
[Fe/H]=-0.36$\pm$0.15 dex), although closer in general to the solar pattern
than to the $\lambda$ Boo stars. In this sense, a primary showing a $\lambda$
Boo pattern and a secondary showing near solar abundances, verify the early
result of Andrievsky et al. (1995): the early-type stars A and B present
different chemical compositions.

Figure 4: Chemical pattern of the stars HD 15164, HD 15165 and HD 15165C
(black), compared to an average pattern of $\lambda$ Boo stars (blue).

To our knowledge, there is no abundance determination of $\lambda$ Boo stars
that belong to a triple or multiple system. In particular, a late-type star
that belong to such system, could be used as a proxy of the initial
composition of the material where the $\lambda$ Boo star formed (under the
hypothesis that they born from the same molecular cloud). This could be
important as an additional constrain for any model trying to explain the
$\lambda$ Boo phenomena. We present in the Fig. 4 the chemical pattern of HD
15165C, the late-type component of the triple system. The chemical pattern is
compatible with a solar-like composition (for instance, [FeI/H]=0.04$\pm$0.02
dex). This is in agreement with the idea that $\lambda$ Boo stars are
Population I objects and originate (following any internal or external
mechanism) starting from a solar-like composition.

Notably, the three stars that belong to the triple system present different
chemical patterns. The star A present a $\lambda$ Boo pattern, while the stars
B and C present abundances closer to the Sun. However, the stars B and C are
also slightly different between them: the late-type star C present the closest
abundances to the Sun, while the early-type star B shows a slightly
deficiency. Most abundance values between stars B and C agree within
$\sim$0.30 dex, with a possible exception: the lithium content. The Li I
6707.8 Å line is clearly present in the spectra of the star B (HD 15164) as we
can see in the Fig. 5, while it is not detected in the spectra of stars A nor
C. It is interesting to note that this line is commonly used as a proxy of
recent accretion onto the atmosphere of the stars. For instance, Saffe et al.
(2017) attributed a notable difference in the refractory abundances and in the
Li content between the stars of the binary system HAT-P-4 to a possible
accretion event of a rocky planet onto the primary. However, although HD 15164
shows clearly the Li line, its refractory content is slightly lower than the
star HD 15165C, which would be difficult to explain with the accretion of
refractory species.

Figure 5: Observed spectra (black line) and synthetic spectra (blue dotted
line) near the Li line 6707.8 Å in the star HD 15164. Synthetic lines are
indicated showing the wavelength, atomic number and intensity.

Is it possible that the supposed different abundances between stars A, B and C
are only due to different Teff? The question makes sense because the stars A
and C present Teff of 7150 K and 4960 K, a difference of 2190 K. However, the
total error etot in abundances includes the error e2, which measure the change
in the abundances when varying Teff by their corresponding uncertainty. Then,
we do not expect a strong change in the derived abundances due to Teff (in any
case, the possible change is contained within the total error etot).

###  4.3 The binary system HD 193256/281

HD 193256 was classified as $\lambda$ Boo by Gray (1988) and then as uncertain
$\lambda$ Boo by Renson (1990). It is separated by $\sim$27.5 arcsec from HD
193281, which was classified as $\lambda$ Boo by Gray & Garrison (1987). Both
stars HD 193256 and HD 193281 show approximately solar abundances of C and
subsolar Fe in the study of Stürenburg (1993), who analyzed them separately.
However, they also found near solar values for other elements such as Mg and
Si in both stars, which is different from what found in average $\lambda$ Boo
stars. Kamp et al. (2001) found solar values in HD 193281 for N, O and S,
although for C they found -0.61 dex, similarly to Paunzen et al. (1999).
However, more recent classification spectra suggest that only HD 193256 could
belong to the $\lambda$ Boo class (see Tables 1 and 4 of Murphy et al., 2015;
Gray et al., 2017), while HD 193281 display a normal spectra.

In this work, we analyzed the spectra of HD 193256 and HD 193281 considered
both as single, for which the abundances of HD 193281 should be taken with
caution. We present in Fig. 6 the chemical pattern of the stars HD 193256 and
HD 193281 (black), compared to an average pattern of $\lambda$ Boo stars
(blue). The colors, panels and error bars used are similar to those of Fig. 4.
HD 193256 shows solar or suprasolar values for C and O, together with subsolar
values (between 0.5-0.9 dex) of Ca, Cr, Fe and Sr. However, we also found near
solar values of Mg, Si and Ti, which is not common in $\lambda$ Boo stars.
Then, this object seem to present a mix of metals with solar and subsolar
abundances. On the other hand, HD 193281 present the chemical pattern of a
slightly metal-deficient star in general, showing subsolar values for C and O
($\sim$0.3 dex) similar to Fe I (-0.36$\sim$0.13 dex). However, the results of
HD 193281 should be taken with caution, due to a possible contamination of the
nearby K2 III star.

Figure 6: Chemical pattern of the stars HD 193256 and HD 193281 (black),
compared to an average pattern of $\lambda$ Boo stars (blue).

In short, the solar abundances of some metals of HD 193256 (Mg, Si and Ti) are
different of $\lambda$ Boo stars. The chemical pattern of HD 193281
(considered as single) shows a slightly metal deficient star. In addition,
there is evidence for a possible contamination of HD 193281, where the
components A and B display spectral types A2 III and K2 III. Then, current
evidence does not support the presence of two bonafide $\lambda$ Boo stars in
this binary (or triple) system. It would be desirable an analysis of HD 193281
separately for the components A and B, in order to more properly determine the
individual abundances.

###  4.4 The binary system HD 198160/161

HD 198160 form a visual binary system with HD 198161, separated by $\sim$2.4
arcsec. HD 198160 was classified ”A2 Vann wk4481” and ”A2 Vn” (Gray, 1988;
Corbally & Garrison, 1980), while HD 198161 was classified as ”A3 Vn”
(Corbally & Garrison, 1980). Both stars were studied separately by Stürenburg
(1993) considering them as twins (same Teff and log g). He derived near solar
values for C in both stars and subsolar values for Fe (-0.8$\pm$0.2 dex),
however he also obtained solar values for Mg and Si (0.0$\pm$0.1 dex and
-0.2$\pm$0.2 dex for both stars). Then, Paunzen et al. (1999) estimated near
solar NLTE values for C and O, although quoted for HD 198160/1 (not
separated). More recently, Murphy et al. (2015) caution that individual NLTE
volatile abundances for HD 198160 and HD 198161 are not confirmed (such as
those reported in this work) and tentatively adopt for HD 198160 a
classification ”A2 Vann $\lambda$ Boo”. However, its companion HD 198161 was
classified as a normal star, with spectral type ”A3 V” and ”A3 IV(n)” (Murphy
et al., 2015; Gray et al., 2017).

We present in Fig. 7 the chemical pattern of the stars HD 198160 and HD 198161
(black), compared to an average pattern of $\lambda$ Boo stars (blue). The
colors, panels and error bars used are similar to those of Fig. 4. In both
stars, most Fe-peak metals show a deficiency around 0.7-0.8 dex, similar to
$\lambda$ Boo stars. However, C and O also show subsolar values, being
possibly low compared to other $\lambda$ Boo stars. When comparing C with Fe
abundances, the group of $\lambda$ Boo stars present [C/Fe]$\sim$1.21$\pm$0.35
dex (excluding stars without CNO values and the stars analyzed here, Heiter et
al., 2002) with minimum and maximum values of 0.70 and 1.74 dex. However, the
stars HD 198160 and HD 198161 present [C/Fe] values of $\sim$0.54 and
$\sim$0.48 dex, being low compared to the average [C/Fe] and even lower than
the minimum of 0.70 dex. Then, we consider that these low [C/Fe] values
possibly correspond to mild-$\lambda$ Boo stars, rather than to an average
$\lambda$ Boo object. It is important to note that our C and O abundances were
corrected by NLTE, with average corrections of -0.15 dex and -0.81 dex for
both stars. In other words, if we only adopt LTE values without correction,
the C and O abundances would result closer to those of $\lambda$ Boo stars.

Figure 7: Chemical pattern of the stars HD 198160 and HD 198161 (black),
compared to an average pattern of $\lambda$ Boo stars (blue).

###  4.5 On the physical association of the stars

The stars studied in this work were previously reported as (possible) members
of binary/multiple systems, for the case of HD 15164/65/65C (Andrievsky et
al., 1995; Chernyshova et al., 1998; Murphy et al., 2015), HD 193256/281
(Paunzen et al., 2012a; Murphy et al., 2015; Gray et al., 2017) and HD
198160/161 (Paunzen et al., 2012a; Murphy et al., 2015).

The coordinates, proper motions and parallax of the stars (see Table 1)
suggest that they are, at least, common proper motion objects. We searched our
targets stars in different binary catalogs from literature (Shaya & Olling,
2011; Tokovinin & Lepine, 2012; Andrews et al., 2017). In particular, Andrews
et al. (2017) performed a search of binaries through a Bayesian formulation in
the Tycho-Gaia catalogs and derived likelihoods of Keplerian orbits. For HD
15164/65, they reported a probability greater than 99% that they form a
physical system. Shaya & Olling (2011) developed a Bayesian method to discover
non-random pairs using Hipparcos data. They include HD 198160/161 in their
catalog of binaries, however there is no probability quoted for this pair.
Finally, we find no record for HD 193256/181 in these binary catalogs.

In this work, we assume that the stars form physical binary/multiple systems.
In the case of showing that the stars are not gravitationally bound, then
these stars would not be useful to test the accretion scenario.

###  4.6 Are there two bonafide $\lambda$ Boo stars in binary systems?

There is evidence in the literature supporting the accretion scenario. For
example, the anticorrelation of C and O with Si (Paunzen et al., 1999), first
noted by Holweger & Sturenburg (1993) for C. It is expected that refractory
elements like Fe and Si are condensed in dust, while the more volatile CNO and
S remain in the gaseous phase. Then, the selective accretion of gas will
produce ratios [C/Si] or [O/Si] larger than solar and reduced metallicity
(Paunzen et al., 1999). Kamp et al. (2001) reached a similar conclusion
comparing the volatile species N and S with the more refractory Ca. We should
also expect that in stars with large $v\sin i$, the meridional circulation
mixes material of solar composition from the stellar interior into the
convection zone so that any surface contamination due to accretion of
circumstellar material should vanish. This observation seems to be weakly
verified (see e.g., Solano et al., 2001), and would require a larger sample of
$\lambda$ Boo stars. As we can see, the accretion scenario could be tested by
different methods.

In this work, we focus on the presence of $\lambda$ Boo stars as members of
binary systems (e.g., Stürenburg, 1993; Paunzen et al., 2002; Heiter et al.,
2002; Paunzen et al., 2012a, b). These are the following 12 systems (see
Appendix): HD 15164/65/65C, HD 38545, HD 64491, HD 84948, HD 111786, HD
141851, HD 148628/638, HD 171948, HD 174005, HD 193256/281, HD 198160/161, and
HD 210111. Following the accretion scenario, two early-type stars in a binary
system should display, in principle, a similar $\lambda$ Boo pattern after
passing through a diffuse cloud. However, a binary or multiple system having a
$\lambda$ Boo star together with a ”normal” early-type component would be
difficult to explain under the accretion scenario. This test of the accretion
scenario would require a detailed analysis of both stars. As explained in the
Introduction, some stars that belong to these 12 systems were recently
classified as non-members or uncertain members of the $\lambda$ Boo class,
such as HD 141851, HD 148638 and HD 193256 (see, e.g., Murphy et al., 2015;
Gray et al., 2017). Then, we wonder if any of these 12 systems really include
two stars with bonafide $\lambda$ Boo chemical patterns.

It would be desirable a detailed abundance analysis in order to verify the
true $\lambda$ Boo nature of a star, initially suggested (for instance) by its
classification spectra (see, e.g., Andrievsky et al., 2002; Heiter et al.,
2002). To our knowledge, only 5 out of the 12 systems present an abundance
determination of both components: HD 15164/65, HD 84948, HD 171948, HD
193256/281 and HD 198160/161 (three of them were analyzed in this work). Some
works present an abundance study only of the brighter component, such as in
the case of HD 38545 (Stürenburg, 1993) or HD 64491 (Kamp et al., 2001), while
other systems only have a spectral classification, such as HD 174005 (Gray et
al., 2017; Murphy et al., 2015).

An inspection of the abundance values reported in the literature (see
Appendix) shows that, in our opinion, there is no binary system having two
stars with bonafide $\lambda$ Boo chemical patterns. The same is valid for the
three systems analyzed in this work (HD 1564/65/65C, HD 193256/281 and HD
198160/161). In fact, we cannot find even one binary system where the two
stars present bonafide $\lambda$ Boo abundance patterns. We consider that the
closer candidates to show both stars a $\lambda$ Boo pattern are possibly the
binary systems HD 84948, HD 171948 and HD 198160. These three systems show
[C/Fe] values lower than 0.7 dex (the minimum [C/Fe] of $\lambda$ Boo stars,
see Sect. 4.4 and Appendix), being perhaps mild-$\lambda$ Boo systems rather
than clear $\lambda$ Boo objects. Then, we find no clear evidence for the
presence of two $\lambda$ Boo stars as members of binary systems. However,
this fact (if confirmed) do not rule out the accretion scenario.

On the other hand, a challenge for the accretion scenario, would be the
presence of a bonafide $\lambda$ Boo star and a normal early-type object,
together in the same multiple system. By reviewing the 12 systems studied
(including the stars of this work) we found only one candidate: the system HD
15164/65/65C analyzed here. The star A present a $\lambda$ Boo pattern, while
the stars B (early-type) and C (late-type) present abundances closer to the
Sun. The different chemical composition between stars A and B was initially
attributed to a possible stellar capture (Andrievsky et al., 1995). The
probability of a binary capture depends on several factors, such as the number
of stars per cubic parsec, the velocity dispersion and the mass of the stars
(e.g., Clarke & Pringle, 1991; Boffin et al., 1998). The capture is not a
dominant formation process for solar-mass (coeval) binaries in dense clusters
(e.g., Clarke & Pringle, 1991; Heller, 1995; Boffin et al., 1998). To our
knowledge, there is no known binary or triple system with an origin attributed
to a capture. On the other hand, there are multiple observations of young
binaries embedded in dense cores (e.g., Sadavoy & Stahler, 2017), and even an
image of a triple protostar formed via disk fragmentation (Tobin et al.,
2016). Although the capture cannot be totally discarded, most observational
evidence points toward the formation of binary and multiple systems from a
common molecular cloud. Taking up the idea that the three stars are born
together, it is difficult to explain the composition of the stars of HD 15165
under the accretion scenario. Then, there is an urgent need of additional
binary and multiple systems to be analyzed through a detailed abundance
analysis, in order to test the accretion model of $\lambda$ Boo stars.

##  5 Concluding remarks

In the present work, we performed a detailed abundance determination of
selected binary and multiple systems with candidate $\lambda$ Boo stars, in
order to test the accretion scenario. Reviewing abundance values reported in
the literature (see Appendix) shows that, in our opinion, there are no binary
system having two stars with bonafide $\lambda$ Boo chemical patterns. The
same is valid for the three systems analyzed in this work (HD 15164/65/65C, HD
193256/281 and HD 198160/161). We consider that the closer candidates to show
both stars a $\lambda$ Boo pattern are possibly the binary systems HD 84948,
HD 171948 and HD 198160. However, these three binary systems are perhaps
mild-$\lambda$ Boo systems rather than clear $\lambda$ Boo objects. Then, in
our opinion, current evidence of binary/multiple systems does not give strong
support to the accretion scenario of $\lambda$ Boo stars.

On the other hand, a binary/multiple system formed by a $\lambda$ Boo star and
an early-type ”normal” object, would be difficult to explain under the
accretion scenario. We found one candidate: the remarkable triple system HD
15164/65/65C. It is composed by two early-type stars (A and B) and a late-type
companion (C). In particular, the late-type component of the system could be
used as a proxy for the initial composition of the system, constraining
formation models of $\lambda$ Boo stars. We found a $\lambda$ Boo pattern for
the A star (HD 15165), while the stars B and C present abundances closer to
the Sun. Then, there is an urgent need of additional binary and multiple
systems to be analyzed through a detailed abundance analysis, in order to test
the accretion model of $\lambda$ Boo stars.

###### Acknowledgements.

We thank the referee Dr. Christopher Corbally for constructive comments that
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. 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.

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##  Appendix A Multiple systems with suspected $\lambda$ Boo components

We review abundance determination of binary or multiple systems with suspected
$\lambda$ Boo components from the literature, in order to determine if two
bonafide $\lambda$ Boo stars can be found. Spectral classification data is
also included whenever available. The data are updated including the results
from the present work.

HD 15164/65/65C: It is a visual triple system, where most works considered
only the two brighter components. Andrievsky et al. (1995) studied spectra of
the stars A and B (HD 15165 and HD 15164) using the LYNX (R$\sim$24000) and
AURELIE (R$\sim$11000) spectrographs. They found subsolar values for two
elements analyzed in the A star (-0.73 dex for [Ca/H] and -0.46 dex for
[Fe/H]). Then, Chernyshova et al. (1998) reanalyzed the data for the A star
and suggest that this object belongs to the $\lambda$ Boo class, showing
$\sim$solar values for C, O and S (0.0 dex, -0.3 dex and 0.0 dex) together
with subsolar values for refractory elements (for example, [Fe/H]=-1.6 dex).
However, Andrievsky et al. (1995) also found solar values for several elements
in the B star. They suggest that the different chemical composition of stars A
and B is probably due to a stellar capture. Murphy et al. (2015) classified
the spectra of the 3 stars as ”F1 V kA7mA6 ($\lambda$ Boo)?” (HD 15164), ”F2 V
kA2mA2 $\lambda$ Boo?” (HD 15165) and ”K2V” (HD15165C). They also claim that
the classification spectrum of HD 15165 does not match solar abundances,
contrary to the result of Andrievsky et al. (1995).

In the present work, we find that the star A present a $\lambda$ Boo pattern,
while the stars B and C present abundances closer to the Sun. In other words,
we find different abundances for the stars A and B, in agreement with
Andrievsky et al. (1995). This is difficult to explain under the accretion
scenario of $\lambda$ Boo stars. Then, current evidence does not support the
presence of two bonafide $\lambda$ Boo components in this system.

HD 38545: Stürenburg (1993) estimated solar abundances for C (-0.1$\pm$0.2
dex) and almost solar values for other metals such as Fe (-0.2$\pm$0.2 dex).
However, it was analyzed as a single object and then considered not reliable
by Heiter et al. (2002). Then, this object was mentioned as a possible visual
binary with a small separation ($<$0.2”, Heiter et al., 2002) and as a
possible SB system (Paunzen et al., 2002). More recently, Prugniel et al.
(2011) reported a low metallicity for this object ([Fe/H]=-0.48 dex)
considered also as single. By inspecting IUE UV spectra, Murphy et al. (2015)
suggest that it is a normal object rather than a $\lambda$ Boo star (”non-
member” of the class), and caution that its high $v\sin i$ ($\sim$191 km/s)
may also have had some role in early identifications as $\lambda$ Boo. We note
that this star is not included in the list of SB $\lambda$ Boo stars of
Paunzen et al. (2012a). To our knowledge, there is no spectral classification
nor abundance determination for the secondary.

HD 64491: Kamp et al. (2001) identified this object as a SB system, a
previously undetected binary, showing high and low $v\sin i$ components. They
estimated abundances for the star with higher $v\sin i$ ($\sim$ 170 km/s) by
directly fitting the composite spectra, obtaining [N/H]=-0.30 dex, [S/H]=-0.09
dex and [Ca/H]=-0.96 dex (using NLTE for C and S). Then, Iliev et al. (2001)
reported that the orbital period of this SB system is between 230 and 760
days, and suggest that a new abundance analysis should be performed taking
into account the binarity of the system. Faraggiana & Gerbaldi (2003) suggest
that this object is composed by two slightly metal-poor objects ($\sim$-0.5
dex) rather than a single object with [M/H]$\sim$-1.5 dex. Murphy et al.
(2015) classified the primary of the system as ”F1 Vs kA3mA3 $\lambda$ Boo”.
To our knowledge, there is no spectral classification nor abundance
determination for the secondary (the object with lower $v\sin i$).

HD 84948: Paunzen et al. (1998) reported this object as a SB system and found
subsolar abundances separately for the stars A and B ([Fe/H]= -1.2$\pm$0.3 dex
and -1.0$\pm$0.2 dex, respectively). Then, Heiter et al. (2002) also performed
a detailed abundance determination separately for components A and B. Both
works reported that that the two stars are metal-poor, however CNO or S
abundances were not reported. Then, Iliev et al. (2002) estimated NLTE
abundances for C and O: they find subsolar values for C (-0.8$\pm$0.4 dex for
both stars) while for O they found -0.6$\pm$0.3 dex and +0.2$\pm$0.3 for stars
A and B. They also reported a period of 7.41 d for this SB2 system.

We present in the Fig. 8 a comparison of an average $\lambda$ Boo pattern777We
excluded from the average stars without CNO values and the stars analyzed
here. taken from Heiter et al. (2002), and literature abundances for the stars
A and B. This plot shows that C abundances seem to be low respect of $\lambda$
Boo stars. When comparing C with Fe abundances, the group of $\lambda$ Boo
stars present [C/Fe]$\sim$1.21$\pm$0.35 dex (excluding stars without CNO
values and the stars analyzed here, Heiter et al., 2002) with minimum and
maximum values of 0.70 and 1.74 dex. However, the stars A and B present [C/Fe]
values of $\sim$0.4 and $\sim$0.2 dex888Using Fe from Heiter et al. (2002)
instead of Paunzen et al. (1998), the values are even lower: $\sim$0.3 and
$\sim$0.1 dex for stars A and B., being low values compared to the average
[C/Fe] and even lower than the minimum of 0.70 dex. These low [C/Fe] values
possibly correspond to an extreme or mild-$\lambda$ Boo star rather than to an
average $\lambda$ Boo object.

Paunzen et al. (2001) classified HD 84948 as ”kA7hF1mA6 V (LB)”, while Murphy
et al. (2015) classified HD 84948 as ”F1.5 Vs kA5mA5 $\lambda$ Boo?”, a
”probable member” of the $\lambda$ Boo class using a newer spectra. Given the
low values of [C/Fe] for both stars together with the ”probable” spectral
classification, we prefer to consider them as candidate $\lambda$ Boo stars
(perhaps mild-$\lambda$ Boo stars) rather than bonafide members of the class.
This binary system deserves a verification of the abundance values.

Figure 8: Comparison of an average $\lambda$ Boo pattern (blue, Heiter et al.,
2002) with the abundances from literature for the stars HD 84948 A and B (left
and right panels, black).

HD 111786 (= HR 4881): This star is considered as a classic $\lambda$ Boo
object by different works (e.g., Murphy et al., 2015). Stürenburg (1993)
derived abundances in agreement with the $\lambda$ Boo class (for example,
[C/H]=-0.2$\pm$0.2 dex and [Fe/H]=-1.5$\pm$0.3 dex). However, it was analyzed
as a single object and then considered not reliable by Heiter et al. (2002).
Then, some authors propose an SB nature for this system (Faraggiana et al.,
1997; Paunzen et al., 2012b). We refer the reader to Murphy et al. (2015) for
a more complete discussion about this object. The star was classified as ”F0 V
kA1mA1 $\lambda$ Boo” (Murphy et al., 2015; Gray et al., 2017) and ”F0 Vs
kA1mA1 $\lambda$ Boo” (Murphy et al., 2020). Notably, Faraggiana et al. (2001)
proposed that HD 111786 is in fact a multiple system composed by five members:
one broad-lined star and four narrow-lined stars with similar temperature.
Beyond the multiplicity of this system, to our knowledge there is no spectral
classification nor abundance determination for the secondary (or any other
component) of the system.

HD 141851: Paunzen et al. (1999) found [C/H] and [O/H] NLTE abundances of
-0.81 and -0.21 dex, respectively, showing $v\sin i$ in excess of 200 km
s${{}^{-}1}$. Kamp et al. (2001) derived LTE abundances of [Ca/H]=-1.30 dex,
with typical errors of 0.2 dex. However, Heiter et al. (2002) mention that
this object was analyzed as a single star and then the abundances are not
reliable. Then, different works claim that this object was misclassified and
did not belong to the $\lambda$ Boo class (e.g., Paunzen et al., 2001).
Andrievsky et al. (2002) found [Fe/H] = -0.70, [Si/H] = -0.65 and [Na/H] =
+0.60 dex, however they do not decide if this object is a $\lambda$ Boo star.
Then, Murphy et al. (2015) classified this object as a normal ”A2 IVn” star,
while Gray et al. (2017) as ”A2 IV-Vn”, i.e. non-member of the $\lambda$ Boo
class. To our knowledge, there is no spectral classification nor abundance
determination for the secondary.

HD 148628/638: The primary of this visual pair (HD 148638) was analyzed by
Kamp et al. (2001) obtaining solar values of N and S together with subsolar Ca
(-1.20 dex). However, Murphy et al. (2015, 2020) and Gray et al. (2017)
classified this object as ”A2 IV-n (4481-wk)” and ”A2 IVn” rather than a
member of the $\lambda$ Boo class. To our knowledge, there is no spectral
classification nor abundance study for the companion (HD 148628).

HD 171948: Together with HD 84948, Paunzen et al. (1998) identified this
object as the first SB systems with $\lambda$ Boo components. They reported
very low abundances for Mg, Ti, Cr and Fe separately for the components A and
B. Then, Heiter et al. (2002) derived LTE abundances for this system,
estimating the same values within the errors for both stars. For C they
obtained an upper limit ([C/H]¡-0.5 dex), while O is considered for the same
authors as deficient ([O/H]=-0.6$\pm$0.4 dex) although high compared to heavy
elements ([Fe/H]=-1.6$\pm$0.4 dex). Then, Iliev et al. (2002) reported NLTE
abundances for C and O in this system, estimating the same values for both
stars within the errors ([C/H]=-1.2$\sim$0.4 dex and [O/H]=+0.2$\sim$0.3).
They also derived a period of 21.9 days for the SB system.

We present in the Fig. 9 a comparison of an average $\lambda$ Boo pattern
(Heiter et al., 2002) with the literature abundances of stars A and B, showing
that C values seem to be low respect of $\lambda$ Boo objects. Comparing C and
Fe abundances, both stars A and B present [C/Fe] values of $\sim$0.4 dex
(taking NLTE C values from Iliev et al. 2002 and Fe from Heiter et al. 2002),
being lower than the average [C/Fe] of $\lambda$ Boo stars
($\sim$1.21$\pm$0.35 dex excluding stars without CNO values and the stars
analyzed here, Heiter et al., 2002) and lower than the minimum of 0.70 dex
(Heiter et al., 2002). We consider that these low [C/Fe] values possibly
correspond to an extreme or mild-$\lambda$ Boo star rather than to an average
$\lambda$ Boo object.

Murphy et al. (2015) classified the primary of this binary system as ”A3 Va-
kB8.5 $\lambda$ Boo”, however, there is no spectral classification listed for
the secondary (see their Table 1). Given the low values of [C/Fe] for both
stars and the lack of a spectral classification for the secondary, we prefer
to consider them as candidate $\lambda$ Boo stars (perhaps mild-$\lambda$ Boo
stars) rather than bonafide members of the class. This binary system deserves
a verification of the abundance values.

Figure 9: Comparison of an average $\lambda$ Boo pattern (blue, Heiter et al.,
2002) with the abundances from literature for the stars HD 171948 A and B
(left and right panels, black).

HD 174005: This object was mentioned as a possible SB system with a maximum
separation of $\sim$38 arcsec (Paunzen, 2000; Solano et al., 2001; Paunzen et
al., 2012a). Both Gray et al. (2001) and Murphy et al. (2015) classified this
object as ”A7 V kA2 mA2 $\lambda$ Boo”. To our knowledge, there is no
abundance determination for the components of this system, nor spectral
classification for the secondary. This system would deserve a further
analysis.

HD 193256/281: The star HD 193281 resulted with near solar C (-0.2$\pm$0.2
dex) and subsolar Fe (-1.0$\pm$0.2 dex), however also with near solar values
of Mg, Ti, Cr and Sr in the study of Stürenburg (1993). Then, Paunzen et al.
(1999) estimated a NLTE oxygen abundance of -0.61 dex. Kamp et al. (2001)
found solar values in HD 193281 for N, O and S, although for C they found
-0.61 dex, similarly to Paunzen et al. (1999). The star HD 193256 resulted
with near solar C (0.0$\pm$0.2 dex) and subsolar Fe (-0.7$\pm$0.2 dex), but
also near solar values of Mg and Si (0.0$\pm$0.2 and 0.0$\pm$0.3 dex) in the
study of Stürenburg (1993). Then, abundance values for both stars do not seem
to agree with the general pattern of $\lambda$ Boo stars. The spectra of HD
193256 was classified as ”A9 Vn kA2mA2 $\lambda$ Boo” (Murphy et al., 2015)
and similarly as ”A8 Vn kA3mA3 ($\lambda$ Boo)” (Gray et al., 2017). However,
the spectra of HD 193281 was classified as ”A2IVn” (Murphy et al., 2015) and
”A2 IV-V” (Gray et al., 2017). Given a spectral classification in conflict
with the abundances, Murphy et al. (2015) consider HD 193281 as an ”uncertain
member” of the $\lambda$ Boo class.

In this work, we find that HD 193256 present subsolar values of Cr, Mn and Fe,
however also near solar values for Mg, Si and Ti, which is different than
$\lambda$ Boo stars. For HD 193281, we found a chemical pattern compatible
with a a slightly metal deficient star. However, we also caution that HD
193281 is possibly contaminated by a nearby star (see Sect. 3.2). Then,
current evidence does not support the presence of two bonafide $\lambda$ Boo
objects in this system.

HD 198160/161: Both stars were studied separately by Stürenburg (1993)
considering them as twins (same Teff and log g), although Gerbaldi et al.
(2003) criticized this assumption based in their different V and B (0.35 and
0.39 mag). Stürenburg (1993) derived near solar values for C in both stars
(-0.2$\pm$0.3) and subsolar values for Fe (-0.8$\pm$0.2 dex), however he also
obtained solar values for Mg and Si (0.0$\pm$0.1 dex and -0.2$\pm$0.2 dex for
both stars). They also estimated suprasolar values for Na (+0.3$\pm$0.2 dex
and +0.6$\pm$0.2 dex for both stars). Then, Paunzen et al. (1999) estimated
near solar NLTE values for C and O. Murphy et al. (2015) classified the
spectra of both stars as ”A2 Vann $\lambda$ Boo” and ”A3 V” (see their Table
1), respectively, while Gray et al. (2017) classified the spectra of HD 198160
as ”A3 IV(n)”.

In this work, we find a general deficiency of metals around 0.7-0.8 dex for
both stars. However, we also found subsolar values for C and O, possibly low
compared to other $\lambda$ Boo stars. When comparing C with Fe abundances, we
found that the stars HD 198160 and HD 198161 present [C/Fe] values of
$\sim$0.54 and $\sim$0.48 dex, being low compared to the average [C/Fe] of
$\lambda$ Boo stars ($\sim$1.21$\pm$0.35 dex) and even lower than the minimum
of 0.70 dex (see Sect. 4.3). Then, we consider that these low [C/Fe] values
possibly correspond to mild-$\lambda$ Boo stars, rather than to an average
$\lambda$ Boo object. In our opinion, current evidence does not support the
presence of two bonafide $\lambda$ Boo objects in the system.

HD 210111: Stürenburg (1993) analyzed this object as a single star, obtaining
solar abundances for C (0.1$\sim$0.1 dex), a subsolar value for Fe
(-1.1$\sim$0.2 dex), but also obtaining suprasolar and solar values for Sr and
Ba (+0.45$\sim$0.2 and 0.05$\sim$0.2 dex). Solano et al. (2001) obtained
subsolar values for Mg, Cr, Sc and Fe (between -0.8 dex and -1.3 dex), while
Kamp et al. (2001) derived subsolar and near solar abundances for C and O
(-0.45 dex and -0.20 dex, with typical errors of 0.2 dex). Paunzen et al.
(1999) estimated NLTE values for C and O of -0.45 dex and -0.20 dex. We
suppose that the data presented in these abundance works correspond to the
primary of the system, where its binary nature were not reported. This object
was classified as ”kA2hA7mA2 Vas $\lambda$ Boo” with peculiar hydrogen lines
by Gray (1988), and then as ”A9 V kA2mA2 $\lambda$ Boo” by Gray et al. (2017).
A classification spectra for HD 210111 was presented by Paunzen et al.
(2012b), who suggest a SB2 nature for the system. They fitted the observed
data using a composite spectrum with two equal components having [M/H]=-1.0
dex. For a more detailed abundance analysis, the authors suggest additional
spectra for a large separation of the two components. In particular, for the
secondary there is no detailed abundance determination (including for the
volatile species) nor spectral classification.

##  Appendix B Chemical abundances

We present in this section the chemical abundances derived in this work and
their errors. The total error etot was derived as the quadratic sum of the
line-to-line dispersion e1 (estimated as $\sigma/\sqrt{n}$ , where $\sigma$ is
the standard deviation), and the error in the abundances (e2, e3 and e4) when
varying Teff, $\log g$ and vmicro by their corresponding uncertainties999We
adopt a minimum of 0.01 dex for the errors e2, e3 and e4.. For chemical
species with only one line, we adopt as $\sigma$ the standard deviation of
iron lines. Abundance tables show the average abundance and the total error
etot, together with the errors e1 to e4.

Table 3: Chemical abundances for HD 15164. Specie | [X/H] $\pm$ etot | e1 | e2 | e3 | e4
---|---|---|---|---|---
Li I | 1.32 $\pm$ 0.17 | 0.07 | 0.15 | 0.01 | 0.01
C I | -0.30 $\pm$ 0.05 | 0.02 | 0.04 | 0.02 | 0.01
N I | 0.08 $\pm$ 0.10 | 0.07 | 0.05 | 0.01 | 0.04
O I | 0.12 $\pm$ 0.35 | 0.07 | 0.30 | 0.04 | 0.16
Mg I | -0.12 $\pm$ 0.22 | 0.11 | 0.13 | 0.03 | 0.13
Mg II | 0.08 $\pm$ 0.17 | 0.07 | 0.06 | 0.02 | 0.14
Al I | -0.74 $\pm$ 0.29 | 0.02 | 0.08 | 0.02 | 0.28
Si II | -0.30 $\pm$ 0.14 | 0.04 | 0.06 | 0.02 | 0.12
Ca II | -0.26 $\pm$ 0.17 | 0.07 | 0.15 | 0.01 | 0.02
Sc II | -0.30 $\pm$ 0.27 | 0.17 | 0.07 | 0.03 | 0.19
Ti II | -0.20 $\pm$ 0.16 | 0.02 | 0.08 | 0.02 | 0.14
Cr II | -0.35 $\pm$ 0.08 | 0.02 | 0.03 | 0.02 | 0.07
Mn I | -0.38 $\pm$ 0.16 | 0.04 | 0.14 | 0.01 | 0.06
Fe I | -0.36 $\pm$ 0.15 | 0.01 | 0.05 | 0.01 | 0.14
Fe II | -0.37 $\pm$ 0.11 | 0.01 | 0.03 | 0.01 | 0.11
Ni II | -0.50 $\pm$ 0.10 | 0.07 | 0.06 | 0.02 | 0.02
Zn I | -0.53 $\pm$ 0.12 | 0.02 | 0.12 | 0.01 | 0.01
Sr II | 0.37 $\pm$ 0.32 | 0.02 | 0.16 | 0.01 | 0.27
Y II | -0.26 $\pm$ 0.12 | 0.03 | 0.11 | 0.02 | 0.04
Zr II | -0.06 $\pm$ 0.11 | 0.07 | 0.08 | 0.02 | 0.02
Ba II | 0.10 $\pm$ 0.29 | 0.10 | 0.16 | 0.01 | 0.22
Table 4: Chemical abundances for HD 15165. Specie | [X/H] $\pm$ etot | e1 | e2 | e3 | e4
---|---|---|---|---|---
C I | -0.06 $\pm$ 0.07 | 0.02 | 0.04 | 0.05 | 0.02
O I | 0.52 $\pm$ 0.12 | 0.02 | 0.02 | 0.02 | 0.11
Mg I | -1.06 $\pm$ 0.25 | 0.21 | 0.08 | 0.06 | 0.09
Mg II | -1.00 $\pm$ 0.24 | 0.22 | 0.05 | 0.06 | 0.06
Al I | -1.49 $\pm$ 0.28 | 0.10 | 0.17 | 0.09 | 0.18
Ca II | -1.03 $\pm$ 0.28 | 0.26 | 0.09 | 0.01 | 0.04
Sc II | -1.40 $\pm$ 0.28 | 0.22 | 0.11 | 0.06 | 0.12
Ti II | -0.97 $\pm$ 0.16 | 0.06 | 0.04 | 0.06 | 0.13
Cr II | -1.12 $\pm$ 0.08 | 0.02 | 0.07 | 0.03 | 0.01
Fe I | -1.24 $\pm$ 0.16 | 0.06 | 0.09 | 0.01 | 0.12
Fe II | -1.14 $\pm$ 0.07 | 0.04 | 0.04 | 0.03 | 0.04
Sr II | -0.34 $\pm$ 0.34 | 0.07 | 0.13 | 0.01 | 0.31
Ba II | -0.54 $\pm$ 0.26 | 0.15 | 0.08 | 0.03 | 0.19
Table 5: Chemical abundances for HD 15165C. Specie | [X/H] $\pm$ etot | e1 | e2 | e3 | e4
---|---|---|---|---|---
Mg I | -0.25 $\pm$ 0.11 | 0.10 | 0.01 | 0.01 | 0.01
Al I | -0.08 $\pm$ 0.03 | 0.02 | 0.01 | 0.01 | 0.01
Si I | 0.09 $\pm$ 0.10 | 0.08 | 0.06 | 0.01 | 0.01
Ca I | 0.15 $\pm$ 0.07 | 0.04 | 0.05 | 0.01 | 0.01
Sc II | -0.11 $\pm$ 0.06 | 0.06 | 0.01 | 0.01 | 0.01
Ti I | -0.03 $\pm$ 0.06 | 0.03 | 0.05 | 0.01 | 0.01
Ti II | -0.04 $\pm$ 0.05 | 0.05 | 0.01 | 0.01 | 0.01
V I | 0.01 $\pm$ 0.07 | 0.04 | 0.06 | 0.01 | 0.01
Cr I | -0.02 $\pm$ 0.06 | 0.04 | 0.04 | 0.01 | 0.01
Cr II | -0.02 $\pm$ 0.08 | 0.08 | 0.01 | 0.01 | 0.01
Mn I | 0.29 $\pm$ 0.09 | 0.09 | 0.02 | 0.01 | 0.01
Fe I | 0.04 $\pm$ 0.02 | 0.01 | 0.01 | 0.01 | 0.01
Fe II | -0.01 $\pm$ 0.05 | 0.04 | 0.02 | 0.01 | 0.01
Co I | -0.13 $\pm$ 0.05 | 0.04 | 0.01 | 0.01 | 0.01
Cu I | -0.21 $\pm$ 0.18 | 0.18 | 0.01 | 0.01 | 0.01
Zn I | -0.15 $\pm$ 0.24 | 0.24 | 0.01 | 0.01 | 0.01
Sr II | -0.18 $\pm$ 0.13 | 0.13 | 0.01 | 0.01 | 0.01
Y II | 0.04 $\pm$ 0.21 | 0.21 | 0.01 | 0.01 | 0.03
Zr II | 0.24 $\pm$ 0.13 | 0.13 | 0.01 | 0.02 | 0.01
Ba II | 0.53 $\pm$ 0.17 | 0.17 | 0.01 | 0.01 | 0.01
Nd II | 0.12 $\pm$ 0.08 | 0.08 | 0.01 | 0.01 | 0.01
Table 6: Chemical abundances for HD 193256. Specie | [X/H] $\pm$ etot | e1 | e2 | e3 | e4
---|---|---|---|---|---
C I | -0.05 $\pm$ 0.22 | 0.21 | 0.04 | 0.02 | 0.05
O I | 0.74 $\pm$ 0.16 | 0.15 | 0.04 | 0.05 | 0.02
Mg I | 0.34 $\pm$ 0.25 | 0.21 | 0.05 | 0.04 | 0.12
Mg II | 0.02 $\pm$ 0.24 | 0.21 | 0.07 | 0.04 | 0.08
Si II | 0.08 $\pm$ 0.18 | 0.07 | 0.16 | 0.04 | 0.02
Ca II | -0.47 $\pm$ 0.23 | 0.21 | 0.06 | 0.05 | 0.01
Sc II | -0.60 $\pm$ 0.31 | 0.21 | 0.08 | 0.03 | 0.21
Ti II | -0.18 $\pm$ 0.25 | 0.07 | 0.03 | 0.08 | 0.23
Cr II | -0.61 $\pm$ 0.09 | 0.02 | 0.02 | 0.07 | 0.06
Mn I | -0.53 $\pm$ 0.13 | 0.04 | 0.10 | 0.01 | 0.07
Fe I | -0.92 $\pm$ 0.15 | 0.07 | 0.06 | 0.02 | 0.12
Fe II | -0.69 $\pm$ 0.10 | 0.04 | 0.04 | 0.05 | 0.08
Sr II | -0.61 $\pm$ 0.49 | 0.27 | 0.10 | 0.10 | 0.38
Table 7: Chemical abundances for HD 193281. Specie | [X/H] $\pm$ etot | e1 | e2 | e3 | e4
---|---|---|---|---|---
C I | -0.35 $\pm$ 0.10 | 0.07 | 0.07 | 0.01 | 0.01
O I | -0.30 $\pm$ 0.06 | 0.03 | 0.04 | 0.02 | 0.01
Mg I | -0.16 $\pm$ 0.29 | 0.20 | 0.10 | 0.06 | 0.18
Mg II | -0.54 $\pm$ 0.21 | 0.18 | 0.02 | 0.01 | 0.11
Al I | -0.65 $\pm$ 0.25 | 0.18 | 0.08 | 0.02 | 0.15
Si II | -0.84 $\pm$ 0.11 | 0.07 | 0.08 | 0.05 | 0.01
Ca II | -0.27 $\pm$ 0.21 | 0.18 | 0.11 | 0.01 | 0.01
Sc II | -0.23 $\pm$ 0.32 | 0.18 | 0.10 | 0.01 | 0.25
Ti II | -0.24 $\pm$ 0.15 | 0.05 | 0.05 | 0.04 | 0.13
Cr II | -0.53 $\pm$ 0.04 | 0.02 | 0.02 | 0.02 | 0.01
Fe I | -0.36 $\pm$ 0.13 | 0.05 | 0.09 | 0.02 | 0.07
Fe II | -0.48 $\pm$ 0.13 | 0.03 | 0.07 | 0.01 | 0.10
Sr II | -0.04 $\pm$ 0.47 | 0.01 | 0.16 | 0.01 | 0.44
Y II | -0.09 $\pm$ 0.16 | 0.13 | 0.09 | 0.04 | 0.01
Zr II | -0.02 $\pm$ 0.19 | 0.18 | 0.06 | 0.02 | 0.01
Ba II | 0.20 $\pm$ 0.17 | 0.09 | 0.14 | 0.01 | 0.03
Table 8: Chemical abundances for HD 198160. Specie | [X/H] $\pm$ etot | e1 | e2 | e3 | e4
---|---|---|---|---|---
C I | -0.29 $\pm$ 0.08 | 0.07 | 0.01 | 0.04 | 0.03
O I | -0.43 $\pm$ 0.28 | 0.15 | 0.02 | 0.02 | 0.23
Mg I | -0.91 $\pm$ 0.18 | 0.08 | 0.03 | 0.01 | 0.15
Mg II | -0.50 $\pm$ 0.20 | 0.15 | 0.06 | 0.05 | 0.10
Al I | -1.25 $\pm$ 0.21 | 0.15 | 0.05 | 0.03 | 0.13
Si II | -0.86 $\pm$ 0.23 | 0.15 | 0.09 | 0.05 | 0.13
Ca II | -0.65 $\pm$ 0.16 | 0.15 | 0.05 | 0.03 | 0.01
Sc II | -0.85 $\pm$ 0.18 | 0.15 | 0.03 | 0.04 | 0.09
Ti II | -0.73 $\pm$ 0.18 | 0.02 | 0.02 | 0.06 | 0.17
Cr II | -0.68 $\pm$ 0.08 | 0.07 | 0.01 | 0.04 | 0.01
Mn I | -1.06 $\pm$ 0.11 | 0.01 | 0.11 | 0.01 | 0.02
Fe I | -0.83 $\pm$ 0.10 | 0.04 | 0.07 | 0.01 | 0.06
Fe II | -0.83 $\pm$ 0.10 | 0.03 | 0.04 | 0.04 | 0.08
Sr II | -1.29 $\pm$ 0.30 | 0.18 | 0.08 | 0.07 | 0.22
Ba II | -0.47 $\pm$ 0.17 | 0.15 | 0.08 | 0.01 | 0.02
Table 9: Chemical abundances for HD 198161. Specie | [X/H] $\pm$ etot | e1 | e2 | e3 | e4
---|---|---|---|---|---
C I | -0.32 $\pm$ 0.06 | 0.04 | 0.02 | 0.04 | 0.02
O I | -0.21 $\pm$ 0.28 | 0.15 | 0.02 | 0.02 | 0.23
Mg I | -0.87 $\pm$ 0.18 | 0.07 | 0.03 | 0.01 | 0.17
Mg II | -0.55 $\pm$ 0.20 | 0.15 | 0.06 | 0.05 | 0.10
Al I | -1.01 $\pm$ 0.25 | 0.15 | 0.06 | 0.04 | 0.18
Si II | -0.38 $\pm$ 0.21 | 0.15 | 0.08 | 0.05 | 0.10
Ca II | -0.67 $\pm$ 0.16 | 0.15 | 0.05 | 0.03 | 0.01
Sc II | -0.85 $\pm$ 0.18 | 0.15 | 0.03 | 0.04 | 0.09
Ti II | -0.83 $\pm$ 0.17 | 0.05 | 0.03 | 0.06 | 0.14
Cr II | -0.68 $\pm$ 0.07 | 0.06 | 0.01 | 0.04 | 0.01
Mn I | -0.97 $\pm$ 0.14 | 0.08 | 0.11 | 0.01 | 0.03
Fe I | -0.80 $\pm$ 0.17 | 0.03 | 0.07 | 0.01 | 0.15
Fe II | -0.81 $\pm$ 0.11 | 0.04 | 0.04 | 0.04 | 0.08
Sr II | -1.41 $\pm$ 0.22 | 0.04 | 0.08 | 0.07 | 0.19
Ba II | -0.38 $\pm$ 0.18 | 0.15 | 0.07 | 0.01 | 0.06