11institutetext: INAF – Osservatorio Astrofisico di Catania, via S. Sofia 78,
95123 Catania, Italy
11email<EMAIL_ADDRESS>22institutetext: European Southern
Observatory, Karl-Schwarzschild-Strasse 2, 85748 Garching bei München, Germany
33institutetext: INAF – Osservatorio Astronomico di Capodimonte, via
Moiariello, 16, 80131 Napoli, Italy 44institutetext: Konkoly Observatory,
Research Centre for Astronomy and Earth Sciences, Eötvös Loránd Research
Network (ELKH), Konkoly-Thege Miklós út 15-17, H-1121 Budapest, Hungary
55institutetext: ELTE Eötvös Loránd University, Institute of Physics, Pázmány
Péter sétány 1/A, H-1117 Budapest, Hungary 66institutetext: Purple Mountain
Observatory, Chinese Academy of Sciences, 10 Yuanhua Road, Nanjing 210023,
China 77institutetext: INAF – Osservatorio Astronomico di Roma, via di
Frascati 33, 00078 Monte Porzio Catone, Italy 88institutetext: Kavli Institute
for Astronomy and Astrophysics, Peking University, Yiheyuan 5, Haidian Qu,
100871 Beijing, China 99institutetext: Max Planck Institute for Astronomy,
Königstuhl 17, D-69117 Heidelberg, Germany 1010institutetext: SETI Institute,
189 Bernardo Ave, Suite 200, Mountain View, CA 94043, USA 1111institutetext:
Stony Brook University, Stony Brook, NY 11794, USA 1212institutetext: Crimean
Astrophysical Observatory, 298409 Nauchny, Crimea 1313institutetext: Univ.
Grenoble Alpes, CNRS, IPAG, 38000 Grenoble, France 1414institutetext:
Observatoire de Paris, PSL University, Sorbonne Université, CNRS, LERMA,
F-75014 Paris

# PENELLOPE II. CVSO 104: a pre-main sequence close binary
with an optical companion in Ori OB1††thanks: Based on data obtained within
ESO programme 106.20Z8.

A. Frasca PENELLOPE II. CVSO 104: a pre-main sequence close binary  with an
optical companion in Ori OB1††thanks: Based on data obtained within ESO
programme 106.20Z8.PENELLOPE II. CVSO 104: a pre-main sequence close binary
with an optical companion in Ori OB1††thanks: Based on data obtained within
ESO programme 106.20Z8. H. M. J. Boffin PENELLOPE II. CVSO 104: a pre-main
sequence close binary  with an optical companion in Ori OB1††thanks: Based on
data obtained within ESO programme 106.20Z8.PENELLOPE II. CVSO 104: a pre-main
sequence close binary  with an optical companion in Ori OB1††thanks: Based on
data obtained within ESO programme 106.20Z8. C. F. Manara PENELLOPE II. CVSO
104: a pre-main sequence close binary  with an optical companion in Ori
OB1††thanks: Based on data obtained within ESO programme 106.20Z8.PENELLOPE
II. CVSO 104: a pre-main sequence close binary  with an optical companion in
Ori OB1††thanks: Based on data obtained within ESO programme 106.20Z8. J. M.
Alcalá PENELLOPE II. CVSO 104: a pre-main sequence close binary  with an
optical companion in Ori OB1††thanks: Based on data obtained within ESO
programme 106.20Z8.PENELLOPE II. CVSO 104: a pre-main sequence close binary
with an optical companion in Ori OB1††thanks: Based on data obtained within
ESO programme 106.20Z8. P. Ábrahám PENELLOPE II. CVSO 104: a pre-main sequence
close binary  with an optical companion in Ori OB1††thanks: Based on data
obtained within ESO programme 106.20Z8.PENELLOPE II. CVSO 104: a pre-main
sequence close binary  with an optical companion in Ori OB1††thanks: Based on
data obtained within ESO programme 106.20Z8.PENELLOPE II. CVSO 104: a pre-main
sequence close binary  with an optical companion in Ori OB1††thanks: Based on
data obtained within ESO programme 106.20Z8.PENELLOPE II. CVSO 104: a pre-main
sequence close binary  with an optical companion in Ori OB1††thanks: Based on
data obtained within ESO programme 106.20Z8. E. Covino PENELLOPE II. CVSO 104:
a pre-main sequence close binary  with an optical companion in Ori
OB1††thanks: Based on data obtained within ESO programme 106.20Z8.PENELLOPE
II. CVSO 104: a pre-main sequence close binary  with an optical companion in
Ori OB1††thanks: Based on data obtained within ESO programme 106.20Z8. M. Fang
PENELLOPE II. CVSO 104: a pre-main sequence close binary  with an optical
companion in Ori OB1††thanks: Based on data obtained within ESO programme
106.20Z8.PENELLOPE II. CVSO 104: a pre-main sequence close binary  with an
optical companion in Ori OB1††thanks: Based on data obtained within ESO
programme 106.20Z8. M. Gangi PENELLOPE II. CVSO 104: a pre-main sequence close
binary  with an optical companion in Ori OB1††thanks: Based on data obtained
within ESO programme 106.20Z8.PENELLOPE II. CVSO 104: a pre-main sequence
close binary  with an optical companion in Ori OB1††thanks: Based on data
obtained within ESO programme 106.20Z8. G. J. Herczeg PENELLOPE II. CVSO 104:
a pre-main sequence close binary  with an optical companion in Ori
OB1††thanks: Based on data obtained within ESO programme 106.20Z8.PENELLOPE
II. CVSO 104: a pre-main sequence close binary  with an optical companion in
Ori OB1††thanks: Based on data obtained within ESO programme 106.20Z8. Á.
Kóspál PENELLOPE II. CVSO 104: a pre-main sequence close binary  with an
optical companion in Ori OB1††thanks: Based on data obtained within ESO
programme 106.20Z8.PENELLOPE II. CVSO 104: a pre-main sequence close binary
with an optical companion in Ori OB1††thanks: Based on data obtained within
ESO programme 106.20Z8.PENELLOPE II. CVSO 104: a pre-main sequence close
binary  with an optical companion in Ori OB1††thanks: Based on data obtained
within ESO programme 106.20Z8.PENELLOPE II. CVSO 104: a pre-main sequence
close binary  with an optical companion in Ori OB1††thanks: Based on data
obtained within ESO programme 106.20Z8.PENELLOPE II. CVSO 104: a pre-main
sequence close binary  with an optical companion in Ori OB1††thanks: Based on
data obtained within ESO programme 106.20Z8.PENELLOPE II. CVSO 104: a pre-main
sequence close binary  with an optical companion in Ori OB1††thanks: Based on
data obtained within ESO programme 106.20Z8. L. Venuti PENELLOPE II. CVSO 104:
a pre-main sequence close binary  with an optical companion in Ori
OB1††thanks: Based on data obtained within ESO programme 106.20Z8.PENELLOPE
II. CVSO 104: a pre-main sequence close binary  with an optical companion in
Ori OB1††thanks: Based on data obtained within ESO programme 106.20Z8. F. M.
Walter PENELLOPE II. CVSO 104: a pre-main sequence close binary  with an
optical companion in Ori OB1††thanks: Based on data obtained within ESO
programme 106.20Z8.PENELLOPE II. CVSO 104: a pre-main sequence close binary
with an optical companion in Ori OB1††thanks: Based on data obtained within
ESO programme 106.20Z8. J. Alonso-Santiago PENELLOPE II. CVSO 104: a pre-main
sequence close binary  with an optical companion in Ori OB1††thanks: Based on
data obtained within ESO programme 106.20Z8.PENELLOPE II. CVSO 104: a pre-main
sequence close binary  with an optical companion in Ori OB1††thanks: Based on
data obtained within ESO programme 106.20Z8. K. Grankin PENELLOPE II. CVSO
104: a pre-main sequence close binary  with an optical companion in Ori
OB1††thanks: Based on data obtained within ESO programme 106.20Z8.PENELLOPE
II. CVSO 104: a pre-main sequence close binary  with an optical companion in
Ori OB1††thanks: Based on data obtained within ESO programme 106.20Z8. M.
Siwak PENELLOPE II. CVSO 104: a pre-main sequence close binary  with an
optical companion in Ori OB1††thanks: Based on data obtained within ESO
programme 106.20Z8.PENELLOPE II. CVSO 104: a pre-main sequence close binary
with an optical companion in Ori OB1††thanks: Based on data obtained within
ESO programme 106.20Z8. E. Alecian PENELLOPE II. CVSO 104: a pre-main sequence
close binary  with an optical companion in Ori OB1††thanks: Based on data
obtained within ESO programme 106.20Z8.PENELLOPE II. CVSO 104: a pre-main
sequence close binary  with an optical companion in Ori OB1††thanks: Based on
data obtained within ESO programme 106.20Z8. S. Cabrit PENELLOPE II. CVSO 104:
a pre-main sequence close binary  with an optical companion in Ori
OB1††thanks: Based on data obtained within ESO programme 106.20Z8.PENELLOPE
II. CVSO 104: a pre-main sequence close binary  with an optical companion in
Ori OB1††thanks: Based on data obtained within ESO programme 106.20Z8.

(Received / Accepted )

We present results of our study of the close pre-main sequence spectroscopic
binary CVSO 104 in Ori OB1, based on data obtained within the PENELLOPE legacy
program. We derive, for the first time, the orbital elements of the system and
the stellar parameters of the two components. The system is composed of two
early M-type stars and has an orbital period of about 5 days and a mass ratio
of 0.92, but contrarily to expectations does not appear to have a tertiary
companion. Both components have been (quasi-)synchronized, but the orbit is
still very eccentric. The spectral energy distribution clearly displays a
significant infrared excess compatible with a circumbinary disk. The analysis
of He i and Balmer line profiles, after the removal of the composite
photospheric spectrum, reveals that both components are accreting at a similar
level. We also observe excess emission in H$\alpha$ and H$\beta$, which
appears redshifted or blueshifted by more than 100 km s-1 with respect to the
mass center of the system depending on the orbital phase. This additional
emission could be connected with accretion structures, such as funnels of
matter from the circumbinary disk. We also analyze the optical companion
located at about 2$\aas@@fstack{\prime\prime}$4 from the spectroscopic binary.
This companion, that we named CVSO 104 B, turns out to be a background Sun-
like star not physically associated with the PMS system and not belonging to
Ori OB1.

###### Key Words.:

stars: pre-main sequence – stars: binaries: spectroscopic – stars: low-mass –
accretion, accretion disks – protoplanetary disks – stars: individual: CVSO
104

##  1 Introduction

The formation of stars and planets is strongly influenced by the conditions in
their environment (circumstellar disks, jets and winds) during their early
life. The majority of stars are formed in binary and multiple systems, which,
for sufficiently small separations, allows obtaining information on their
components, such as dynamical masses. It is therefore very important to detect
and study young pre-main sequence (PMS) close binaries for deriving their
orbital and stellar properties and to understand how the binarity affects the
planet formation. Moreover, well-characterized early-PMS binary systems are
critical to constrain theoretical evolutionary tracks, which are used to
derive fundamental stellar properties, but are often affected by offsets of up
to 50% in the predicted masses compared to the dynamical mass estimates (e.g.,
Covino et al. 2004; Hillenbrand & White 2004; Stassun et al. 2014). PMS
binaries are also very important for studying the mass accretion process in a
different field geometry from that of single stars. Both observations and
numerical simulations show that quasi-periodic bursts of accretion are
expected in close binaries, with a different pattern for circular and
eccentric systems (e.g., Muñoz & Lai 2016; Gillen et al. 2017). For eccentric
systems, the accretion bursts, originating from the impact of nearly free-
falling matter on the high atmospheric layers of each of the two components,
are more frequently observed near the periastron passages (e.g., Kóspál et al.
2018; Tofflemire et al. 2019). The simulations also show gas streams from the
circumbinary disk flowing to each component through the $L_{2}$ and $L_{3}$
Lagrangian points (e.g., de Val-Borro et al. 2011; Muñoz & Lai 2016).

We present here the spectroscopic orbit of CVSO 104 (Haro 5-64), which is a
classical T Tauri star located in the $\sim$ 5 Myr old Orion OB1b association.
It was discovered as a star with H$\alpha$ emission in the region of the
Horsehead Nebula (IC 434) by Haro & Moreno (1953) with objective-prism
observations. It was detected in the Kiso survey of emission-line objects by
Wiramihardja et al. (1989) who confirmed a strong H$\alpha$ emitter. Briceño
et al. (2005) observed this star as part of the CIDA Variability Survey of
Orion OB1 (CVSO) and reported it as a variable source with H$\alpha$ and Li i
$\lambda$6708 equivalent widths of $-62.90$ Å and 0.30 Å, respectively.

The spectra analyzed in the present paper have been obtained as part of the
PENELLOPE program running at the Very Large Telescope (VLT). PENELLOPE is a
ground-based follow-up large program (see Manara et al. 2021, hereafter Paper
I) of the ULLYSES HST program111ullyses.stsci.edu (Roman-Duval et al. 2020).

The Gaia EDR3 release lists a parallax $\varpi=2.73\pm 0.03$ mas for the
object, putting it at a distance of 366$\pm$4 pc, i.e., inside the Orion star
forming region. Its magnitude, $G=14.45$, is too faint for Gaia to provide
radial velocity measurements. Based on a limited set of APOGEE-2 spectra,
Kounkel et al. (2019) identified it as a double-lined binary, with a mass
ratio of 0.988$\pm$0.063. It is thus a perfect target for binarity follow-up
studies. As short-period binaries may be induced by the gravitational torque
from an additional companion (e.g., Eggleton & Kiseleva-Eggleton 2001; Naoz &
Fabrycky 2014), it is also interesting to know if the system is in fact
triple.

Gaia lists a slightly fainter ($\Delta G=0.14$) and redder visual companion at
a separation of 2.39$\arcsec$. With a parallax of $\varpi=1.49\pm 0.03$ mas,
corresponding to a distance of about 670 pc, this optical companion is,
however, unrelated to the much closer Ori OB1, as also confirmed by the
analysis of its spectra presented in Sect. 3.3. From now on, we will refer to
this background star as CVSO 104 B, while we will refer to our target as CVSO
104 A.

The paper is organized as follows. Section 2 reports the observations used in
this work. The results of the analysis of the data are then described in
Section 3 and discussed in Section 4. We finally summarize our conclusions in
Section 5.

##  2 Observations

###  2.1 Spectroscopy

High-resolution spectroscopy ($R\simeq 70,000$) was performed with the
Ultraviolet and Visual Echelle Spectrograph (UVES, Dekker et al. 2000) within
the PENELLOPE Large Program (ESO Prog. ID. 106.20Z8; see Paper I). Three UVES
spectra were acquired within two days of the HST observation of CVSO 104 A+B
and a fourth one was purposely requested by us a few days later, after the
clear detection of the two components in the spectra of this spectroscopic
binary. All these spectra were taken with the same position angle of 109∘for
the slit, aligning it with the optical pair. Medium-resolution ($R\sim
10,000-20,000$) broad-wavelength coverage spectroscopy was obtained using the
X-Shooter instrument (Vernet et al. 2011). The strategy of spectroscopic
observations and details on data reduction are explained in Paper I.
Particular attention was paid to the order trace and extraction to separate
the spectra of the two stars of the visual pair.

CVSO 104 A+B was also observed during 5 visits (2 in February 2017, 3 in
October 2017) in the framework of the APOGEE survey (Ahumada et al. 2020),
which obtains $R\sim 22,500$ spectra in the $H$-band. Since the radial
velocities measured on APOGEE spectra and reported in the SDSS Data Release 16
(Jönsson et al. 2020) were derived assuming the object as a single star, they
are not usable. We therefore downloaded the APOGEE spectra222Available at
http://skyserver.sdss.org/dr16/ and derived the velocities of both components,
in the same way as we did for the UVES and X-Shooter spectra (see Sec. 3.2).

Figure 1: TESS light curves of CVSO 104 A+B in 2018 (top) and 2020 (bottom).

###  2.2 Photometry

Space-born accurate photometry was obtained with NASA’s Transiting Exoplanet
Survey Satellite (TESS; Ricker et al. 2015). CVSO 104 A+B was observed in
sector 6 with 1426s exposure times between 2018-12-15 and 2019-01-06, as well
as in sector 32 with 475s exposure times between 2020-11-20 and 2020-12-16 –
this latter data set is contemporaneous to our spectroscopic observations.
Given the pixel size of TESS, the optical companion will contribute to the
light curve. We downloaded the data from the MAST archive and created light
curves using the Lightkurve333docs.lightkurve.org package by using a mask that
contained all the neighboring pixels having S/N$>6$ (see Fig. 1). The 2018
data had a mean flux of 1244$\pm$51 e-/s, while the 2020 data set had a
similar mean flux of 1221$\pm$58 e-/s.

To get the color information that is lacking in TESS data and to separate the
contribution of the two visual components of the optical pair, several ground-
based facilities were involved in the observation of this object during the
HST and VST observations. In the present work we will make use of data taken
with four different instruments.

We observed CVSO 104 at the M. G. Fracastoro station (Serra La Nave, Mt. Etna,
1750 m a.s.l.) of the Osservatorio Astrofisico di Catania (OACT, Italy) from
25 November to 16 December 2020. We used the facility imaging camera at the
0.91 m telescope with a set of broad-band Bessel filters ($B$, $V$, $R$, $I$,
$Z$) as well as two narrow-band H$\alpha$ filters centered on the line core
($H\alpha_{9}$) and on the redward continuum ($H\alpha_{18}$). The index
$H\alpha_{18}$-$H\alpha_{9}$ is basically a measure of intensity of the
H$\alpha$ emission in units of the continuum that can be converted into
equivalent width (EW) of H$\alpha$ (Frasca et al. 2018). Details on
photometric observations and data reduction can be found in Paper I. We note
that the two visual components are clearly resolved only in the nights with
good seeing (see Fig. 2), with the exception of the $Z$ filters in which the
image is slightly out of focus. We have therefore extracted the photometry of
the stars in the field of CVSO 104 from the calibrated images using apertures
of 5″ of radius, which include both components of the visual pair. To measure
the magnitude difference of the two visual components in the images taken with
the best seeing, we have used an approach similar to that of Covino et al.
(2004).

Figure 2: Portion of images acquired at OACT on 16 Dec 2020 in $BVRI$ bands
centered on the visual pair CVSO 104 A+B. North is up and East to the left and
the two components are indicated in the $R$-band image. Note how the relative
brightness of the two stars changes in the different bands, with A becoming
brighter than B in the $I$ band.

CVSO 104 was also observed from 6 November to 26 December, 2020 with the 0.8 m
RC80 telescope of Konkoly Observatory (Hungary) using Bessel $BV$ and Sloan
$r^{\prime}i^{\prime}$ filters. Details about the instrument, data reduction,
and photometry are provided in Ábrahám et al. (in prep.). In most of the
images the binary and the visual companion could be well separated. We
performed PSF photometry separately for CVSO 104 A and B in each image. As we
will show in Sect. 3.1, CVSO 104 B is non-variable, therefore we calculated
its average flux in each filter. Then we added these values to the fluxes of
CVSO 104 A, to obtain a full light curve of the visual pair CVSO 104 A+B
during the run to be used together with other data sets. To this aim the
$r^{\prime},i^{\prime}$ magnitudes were converted to $R_{C},I_{C}$ using the
prescription given by Lupton
(2005)444http://www.sdss3.org/dr8/algorithms/sdssUBVRITransform.php.

We used $BVri$ photometry collected by
AAVSOnet555https://www.aavso.org/aavsonet, which is a set of robotic
telescopes operated by volunteers for the American Association of Variable
Star Observers (AAVSO). Stars in the AAVSO Photometric All-Sky Survey (APASS,
Henden et al. 2018) were used to calibrate the AAVSO photometric data for our
targets. The $r,i$ magnitudes were converted to $R_{C},I_{C}$ as was done for
the Konkoly photometry. We incorporated data taken by the amateur observers of
the AAVSO666Available at https://www.aavso.org/data-download, in response to
AAVSO Alert Notice 725. Finally, some additional $BVRI$ photometry was
obtained at the Crimean Astrophysical Observatory (CrAO) on the AZT-11 1.25 m
telescope.

The photometry obtained with the latter instruments includes both components
in the aperture. The multiband light curve of CVSO 104 A+B during 30 nights
including the ULLYSES and PENELLOPE campaigns is displayed in Fig. 14.

##  3 Results

###  3.1 Photometry

The TESS light curves (Fig. 1) show rather stochastic variations, albeit
quasi-periodic. The period analysis provides different results depending on
the epoch of observations and the different techniques. Periodograms reveal
that the most likely period for the 2018 data set is 4.73 d, while for the
2020 data, this is 4.91 d, similar to the photometric variability period of
4.68 d in the $R$-band reported by Karim et al. (2016), as well as to the
orbital period we derive (see below). Using the CLEAN deconvolution algorithm
(Roberts et al. 1987), we found the maximum power in 2018 at 2.31 d, i.e.
about the half period of the periodogram and the second highest peak at about
4.53 d. In 2020, the highest peak corresponds to 4.73 d. However, for both
data sets, folding the data with any of the above periods does not reveal a
convincing phase diagram.

The ground-based multiband light curve (Fig. 14) shows a stochastic behaviour
with at least two bursts fully observed in $BVRI$ bands. The intensity of the
bursts is clearly stronger in the bluer bands. However, no intensification of
the H$\alpha$ emission is observed in the OACT photometry contemporaneous to
the stronger burst and the variations of H$\alpha$ EW do not seem to be
correlated or anticorrelated with the brightness variations.

While the contribution of CVSO 104 B is included in the photometric data
presented in Fig. 14, our analyses suggest that this source is non-variable or
that its variability is negligible in comparison to CVSO 104 A. As a first
check, we compared the Gaia EDR3 $G$-band magnitude uncertainty of component B
with the corresponding Gaia magnitude uncertainties of other stars of the same
brightness (we used Fig. 5.15 from the Gaia Early Data Release 3 Documentation
V1.1777gea.esac.esa.int/archive/documentation/GEDR3/ for this comparison). We
concluded that the relative flux error of component B, 0.02%, agrees with the
representative numbers from similarly bright non-variable stars for which
similar number of observations were taken. We also plotted the light curves of
CVSO 104 B using the spatially resolved photometry obtained at OACT and
Konkoly observatories (Sect. 2.2). The plot shown in Fig. 15 also confirms
that the source was constant within the measurement uncertainties, if we
neglect very small systematic differences between Konkoly and OACT. In order
to check this claim on more quantitative grounds, we calculated the $\chi^{2}$
and the probability $P(\chi^{2})$ that the magnitude variations have a random
occurrence (e.g., Press et al. 1992). We found, for the Konkoly dataset,
values of $P(\chi^{2})$ of 0.21, 0.22, 0.35, and 0.25 for $B$, $V$, $R_{\rm
C}$, and $I_{\rm C}$ bands, respectively, which indicates these variations as
non significant.

###  3.2 Radial velocity

The radial velocity was measured by cross-correlating the spectra with late-
type templates. For the X-Shooter spectrum we selected only the regions with
the best S/N in the VIS arm – the one with the highest resolution – and
adopted as template a BT-Settl synthetic spectrum (Allard et al. 2012) with
[Fe/H]=0, $T_{\rm eff}$=4000 K, and $\log g$=4.0. For the UVES spectra we
adopted the same template and an HARPS spectrum of GJ 514 (M1V, RV=14.606 km
s-1, Jönsson et al. 2020) obtaining the same results within the errors. This
analysis was carried out with the IRAF888IRAF is distributed by the National
Optical Astronomy Observatory, which is operated by the Association of
Universities for Research in Astronomy, Inc. task fxcor excluding emission
lines and very broad features that can blur the peaks of the cross-correlation
function (CCF). For a better measure of the centroids and full width at half
maximum (FWHM) of the CCF peaks of the two components, we applied a two-
Gaussian fit. The RV error, $\sigma_{\rm RV}$, was computed by fxcor according
to the fitted peak height and the antisymmetric noise as described by Tonry &
Davis (1979). We used the same BT-Settl template to measure RVs in the APOGEE
spectra by cross-correlation. The individual values of RV measured in our and
APOGEE spectra are listed in Table 1.

Table 1: Heliocentric radial velocities of the two components of CVSO 104 A. HJD | RVa | $\sigma_{\rm RV_{a}}$ | RVb | $\sigma_{\rm RV_{b}}$ | Instrument
---|---|---|---|---|---
(2 400 000+) | (km s-1) | (km s-1) |
57794.66675 | 6.59 | 1.49 | 42.67 | 2.29 | APOGEE
57795.66700 | $-10.39$ | 1.40 | 62.56 | 1.36 | APOGEE
58033.96831 | 33.52 | 2.37 | 15.05 | 2.74 | APOGEE
58037.98405 | $-8.64$ | 1.37 | 57.49 | 1.39 | APOGEE
58039.00019 | 35.16 | 1.71 | 8.92 | 1.79 | APOGEE
59178.64280 | $-12.17$ | 0.68 | 63.68 | 0.56 | UVES
59179.62813 | 24.61 | 0.67 | 17.32 | 2.80 | UVES
59180.61009 | 61.23 | 0.19 | $-16.39$ | 0.23 | UVES
59180.63589 | 64.36 | 2.16 | $-14.54$ | 2.40 | XSHOO
59243.55343 | $-13.01$ | 0.61 | 64.25 | 0.59 | UVES
Figure 3: Heliocentric radial velocity curve (circles = UVES, squares =
APOGEE, diamonds = X-Shooter data) of CVSO 104 A. Filled and open symbols have
been used for the primary (more massive) and secondary component,
respectively. The blue and red lines represent the orbital solution (Table 2)
for the primary and secondary component, respectively.

We searched for the best orbital period by applying periodogram analysis
(Scargle 1982) and the CLEAN deconvolution algorithm (Roberts et al. 1987) to
the RVs of the primary and secondary components. The data folded with the best
period (maximum amplitude in the power spectrum) display a smooth variation
with an asymmetrical shape, typical for an eccentric RV orbital motion. Then,
we fitted the observed RV curve with the curvefit routine (Bevington &
Robinson 2003) to determine the orbital parameters and their standard errors.
This also allowed us to improve the determination of the orbital period,
finding $P_{\rm orb}=5.025$ days (see Fig. 3). The orbital parameters are
reported in Table 2. We note that the high eccentricity, $e\simeq 0.39$, is in
line with the simulations of Zrake et al. (2021), who show that equal-mass
binaries accreting from a circumbinary disk evolve toward an orbital
eccentricity of e$\sim$0.45, unless they start with a nearly circular orbit
($e\lesssim 0.08$).

In addition, we searched for signatures of an eventual third companion, yet
unresolved, in UVES spectra of CVSO 104 A, independently on Gaia. For this
purpose we applied the Broadening Function (BF) method (Rucinski 2012), which
is a linear deconvolution operation. Prior to the analysis, all emission lines
were carefully removed from the spectra. However, no signatures of a third
stellar body were found in the BF profiles. Furthermore, no residual peak at
the velocity of the optical companion, CVSO 104 B (RV $\simeq 42.9\pm 0.5$ km
s-1), was found either in the BFs or in the CCFs. This confirms that the
extraction of the spectra allowed us to separate A and B components without
significant contamination.

Table 2: Orbital and stellar parameters of CVSO 104 A.

Orbital parameters | Stellar parameters
---|---
HJD0∗ | 788.156$\pm$0.023 |  | Primary | Secondary
$P_{\rm orb}$ ($d$) | 5.0253$\pm$0.0001 | SpT | M0 | M2
$e$ | 0.395$\pm$0.048 |  $T_{\rm eff}$(K) | 3770$\pm$180 | 3590$\pm$190
$\omega$ (∘) | 160.9$\pm$1.9 |  $v\sin i$ (km s-1) | 7.5$\pm$1.0 | 6.0$\pm$1.0
$\gamma$ (km s-1) | 24.51$\pm$0.17 | $w_{\rm a,b}$ | 0.63$\pm$0.10 | 0.37$\pm$0.10
$K_{\rm a}$ (km s-1) | 45.93$\pm$0.27 |  $W_{\rm Li}$ (mÅ) | 550$\pm$40 | 680$\pm$70
$K_{\rm b}$ (km s-1) | 48.85$\pm$0.27 |  $A$(Li) | 3.2$\pm$0.2 | 3.6$\pm$0.3
$M_{\rm a}\sin^{3}i$ ($M_{\sun}$) | 0.185$\pm$0.013 |  $R$ ($R_{\sun}$) | 1.07$\pm$0.02 | 0.99$\pm$0.02
$M_{\rm b}\sin^{3}i$ ($M_{\sun}$) | 0.171$\pm$0.012 |  $L$ ($L_{\sun}$) | 0.21$\pm$0.04 | 0.15$\pm$0.04
$M_{\rm b}$/$M_{\rm a}$ | 0.921$\pm$0.008 |  $M$ ($M_{\sun}$) | 0.57$\pm$0.15 | 0.43$\pm$0.15
$a\sin i$ ($R_{\sun}$) | 8.73$\pm$0.04 |  |  |

Notes. ∗ Heliocentric Julian Date (HJD-2,457,000) of the periastron passage;
the subscripts a and b refer to the primary (more massive) and secondary
component, respectively.

###  3.3 Stellar parameters

In the PENELLOPE survey we used the code ROTFIT to determine the atmospheric
parameters, $v\sin i$, and veiling for single objects (Frasca et al. 2015,
2017, Paper I). We used this code for determining the parameters of CVSO 104 B
and found $T_{\rm eff}$$=5750\pm 100$ K, $\log g$$=4.40\pm 0.12$ dex,
[Fe/H]$=0.08\pm 0.07$ dex, RV$=42.9\pm 0.5$ km s-1, and $v\sin i$$<2$ km s-1,
i.e. CVSO 104 B turns out to be a slowly-rotating Sun-like star. Furthermore,
there is no lithium $\lambda 6708$ Å absorption line and no chromospheric
emission is visible in the H$\alpha$ line core (see Fig. 16). Therefore, we do
not expect significant brightness variations, compared to CVSO 104 A, from
this background star, in agreement with the results from the ground-based
photometry (Sect. 3.1).

In the case of double-lined spectroscopic binaries (SB2), we cannot use
ROTFIT, therefore we used COMPO2, a code developed in the IDL999IDL
(Interactive Data Language) is a registered trademark of Harris Corporation.
environment by Frasca et al. (2006), which has been adapted to the UVES
spectra. COMPO2 adopts a grid of templates to reproduce the observed spectrum,
which is split into segments of 100 Å each that are independently analyzed. As
grid of templates, we used ELODIE spectra of 34 low-active slowly rotating
stars with a spectral type in the range K2–M5. The resolution of UVES was
degraded to that of the ELODIE templates ($R=42,000$) by convolution with a
Gaussian kernel with the proper width. COMPO2 does not derive the projected
rotation velocities of the two components, which are instead estimated as
$v\sin i$${}_{\rm a}=7.5$ km s-1 and $v\sin i$${}_{\rm b}=6.0$ km s-1 from the
FWHM of the peaks of the CCF and kept as fixed parameters in the code. The RV
separation of the two components is well-known from the CCF analysis and was
used to build the composite “synthetic” spectrum. The flux ratio between the
components, which has been expressed in terms of the flux contribution of the
primary component in units of the continuum, $w_{\rm a}$, is instead let free
to vary until a minimum $\chi^{2}$ is attained for each combination of spectra
(1156 different combinations with the adopted grid). We note that the
combination of two spectra with a relevant velocity separation reduces the
intensity of the photospheric absorption lines of each component in a similar
way as the veiling does. However, no combination of late-type spectra was able
to fairly fit the observed spectrum, unless we included a veiling. After
several trials, we found as the best veiling $r=0.4$ for the UVES segments in
the analyzed red region. An example of the application of COMPO2 is shown in
Fig. 4 for three spectral segments, around 6200 Å, 6400 Å and 6700 Å, of the
UVES spectrum acquired on JD = 2 459 180.

Figure 4: UVES spectrum of CVSO 104 A (black dots) observed on JD = 2 459 180
in three spectral regions around 6200 Å (top), 6400 Å (middle), and 6700 Å
(bottom). In each box the synthetic spectrum, which is the weighted sum of two
standard star spectra mimicking the primary and secondary component, is
overplotted with a full red line. Prominent lines from the primary and
secondary component are marked with a and b, respectively. The hatched green
regions denote the level of the veiling. The residual spectra obtained by the
subtraction of the photospheric template are displayed in the lower panel of
each box with a blue line.

To evaluate the atmospheric parameters (APs) and the flux contribution we kept
only the best 25 combinations (based on the $\chi^{2}$), i.e. about the top
2%, of primary and secondary spectra per each spectral segment and calculated
the averages by weighting with the corresponding $\chi^{2}$. These parameters
are listed in Table 2. The spectral types (SpT) of the components are taken as
the mode of the spectral-type distributions (see Fig. 5).

Figure 5: Distribution of spectral types for the components of CVSO 104 A.

Equivalent widths of the Li i 6708Å line, $W_{\rm Li}$, for the components of
CVSO 104 A have been measured on the residual spectrum (blue line in Fig.4).
This procedure offers the advantage to remove any possible contamination from
nearby iron lines. The $W_{\rm Li}$ values were corrected for the veiling by
multiplying them by $(1+r)$, adopting the value of veiling measured with
COMPO2 at red wavelengths $r=0.4$, and for the flux contribution to the
composite spectrum by dividing them by $w_{\rm a}$ or $w_{\rm b}$ for the
primary and secondary component, respectively. The error on the equivalent
width was estimated as the product of the integration range and the mean error
per spectral point, which results from the standard deviation of the flux
values of the residual spectrum measured at the two sides of the line and is
quoted in Table 2 next to the $W_{\rm Li}$ values. From these values we
derived a Lithium abundance, $A$(Li), of 3.2$\pm$0.2 and 3.6$\pm$0.3 for the
primary and secondary component, respectively, by using the curves of growth
of Zapatero Osorio et al. (2002). The abundance difference between primary and
secondary component is within the uncertainties and is likely related to the
uncertainties introduced by the correction for veiling and flux contribution.
These data are also reported in Table 2.

###  3.4 Accretion and wind diagnostics

The residual spectra are also very helpful to measure equivalent widths and
fluxes of emission lines of the two components, when they are well separated
in wavelength. We were successful to separate the contribution of the two
components for He i $\lambda$5876 Å and $\lambda$6678 Å lines in all the UVES
spectra, with the exception of the spectrum taken at reduced Julian date
RJD=59179, when the system was near a conjunction. We fitted two Gaussians to
each spectrum to measure the centroid (RV), the FWHM, and equivalent width,
$W_{line}$, of the line of each component (blue and red dashed lines in Figs.
6 and 7 for the primary and secondary component, respectively). We recover the
line flux at Earth, $f_{\rm line}$, from the equivalent width $W_{\rm line}$
and the extinction corrected flux of the adjacent continuum, namely $f_{\rm
line}=W_{\rm line}f_{\rm cont}^{0}$ (see, e.g., Alcalá et al. 2017; Frasca et
al. 2018), where the continuum flux is measured in the X-Shooter
spectrum101010A small correction was applied to the continuum flux based on
the photometric variations.  and is corrected for the extinction by the factor
$10^{0.4A_{\lambda}}$, scaling $A_{\lambda}$ from the $A_{V}=0.35$ mag (see
Sect. 3.5). The line luminosity is then derived adopting the distance $d$ as
$L_{\rm line}=4\pi d^{2}f_{\rm line}$.

The radial velocities, equivalent widths, and line luminosities are reported
in Table 3, where we used the subscripts $a$ and $b$ for the quantities
related to the primary and secondary component, respectively. A further broad
emission feature, blueshifted with respect to the primary component, is
visible in the He i $\lambda$5876 line at RJD=59178. In this case we fitted
the observed line profile with three Gaussians and report the parameters of
this feature (brown dashed line in Fig. 7) as $RV_{3}$, $W_{3}$, and $L_{\rm
line}^{3}$ in Table 3.

The H$\beta$ emission profiles of the two components (Fig. 8) are much wider
than the He i lines and, as a consequence, they are strongly blended. However,
we could still deblend them with multi-component fits of Lorentzian profiles.
A third emission component is always visible out of the conjunction as a
blueshifted or redshifted feature, depending on the orbital phase. It displays
the maximum intensity in the first epoch (RJD=59178), when we also observed
the excess He i emission. All the three H$\beta$ emission peaks display a much
smaller intensity at RJD=59243, although the system configuration is the same
as the first epoch of UVES observations, suggesting an intrinsic accretion
variation in addition to any potential geometric effect. The RV, equivalent
width and line luminosity of the H$\beta$ emission components are also quoted
in Table 3.

We cannot distinguish the emission contribution of the two stars in the
H$\alpha$ line (Fig. 17), but we always see the blueshifted and redshifted
excess component at a velocity similar to that of H$\beta$. A two-Gaussian fit
has allowed us to separate the contribution of the latter feature from the
integrated emission coming from the two stars, which we also report in Table
3. We think that this broad excess emission feature can be related to the
complex structure of accretion funnels.

The strong intensity variation from the first to the last spectrum is also
observed in H$\alpha$ and suggests a variable accretion rate.

We have calculated the accretion luminosity, $L_{\rm acc}$, with the $\log
L_{\rm acc}$–$\log L_{\rm line}$ linear relations proposed by Alcalá et al.
(2017). The mass accretion rate, $\dot{M}_{\rm acc}$, was then derived from
$L_{\rm acc}$ according to

$\dot{M}_{\rm acc}=(1-\frac{R_{\star}}{R_{\rm in}})^{-1}\leavevmode\nobreak\
\frac{L_{\rm acc}R_{\star}}{GM_{\star}},$ (1)

where $R_{\star}$ and $R_{\rm in}$ are the stellar radius and inner-disk
radius (assumed to be $R_{\rm in}=5R_{\star}$), respectively (see Gullbring et
al. 1998; Hartmann et al. 1998). At each epoch, we have calculated the mean
values of $L_{\rm acc}$ and $\dot{M}_{\rm acc}$ (reported in Table 4), which
have been obtained averaging the individual values derived from the two He i
and H$\beta$ lines. The errors include the relative errors of the line
luminosities and the standard deviation of $L_{\rm acc}$ values from the three
diagnostics.

Table 3: Multi-component analysis of permitted emission lines of CVSO 104 A in the UVES spectra. RJD | RVa | RVb | RV3 | FWHMa | FWHMb | FWHM3 | $W_{a}$ | $W_{b}$ | $W_{3}$ | $L_{line}^{a}$ | $L_{line}^{b}$ | $L_{line}^{3}$
---|---|---|---|---|---|---|---|---|---|---|---|---
|  (km s-1) |  (km s-1) |  (Å) | ($10^{-6}L_{\sun}$)
He i $\lambda$6678 |
59178 | -10.9$\pm$1.9 | 68.6$\pm$2.2 | … | 38$\pm$5 | 30$\pm$5 | … | 0.48$\pm$0.07 | 0.30$\pm$0.06 | … | 7.9$\pm$1.1 | 5.1$\pm$1.0 | …
59179 | 26.9$\pm$1.9 | … | 36$\pm$5 | … | 0.59$\pm$0.09 | … | 8.8$\pm$1.5 | …
59180 | 60.1$\pm$5.1 | -16.8$\pm$2.1 | … | 57$\pm$13 | 26$\pm$5 | … | 0.25$\pm$0.05 | 0.24$\pm$0.03 | … | 4.2$\pm$0.8 | 3.9$\pm$0.5 | …
59243 | -12.6$\pm$7.7 | 73.1$\pm$3.0 | … | 43$\pm$18 | 30$\pm$7 |  | 0.22$\pm$0.05 | 0.31$\pm$0.04 | … | 3.6$\pm$0.8 | 6.1$\pm$0.8 | …
He i $\lambda$5876 |
59178 | -7.6$\pm$1.2 | 70.2$\pm$1.6 | -82.4$\pm$21.4 | 46$\pm$4 | 45$\pm$4 | 124$\pm$49 | 1.57$\pm$0.08 | 1.07$\pm$0.08 | 0.76$\pm$0.15 | 16.3$\pm$0.8 | 11.2$\pm$0.8 | 7.8$\pm$1.5
59179 | 29.5$\pm$1.1 | … |  |  |  | 1.76$\pm$0.14 | … | 16.6$\pm$1.4 | …
59180 | 65.3$\pm$2.6 | -12.5$\pm$1.3 | … | 58$\pm$7 | 40$\pm$3 | … | 0.86$\pm$0.09 | 0.90$\pm$0.08 | … | 9.0$\pm$0.9 | 9.3$\pm$0.8 | …
59243 | -9.7$\pm$3.2 | 73.7$\pm$1.5 | … | 49$\pm$8 | 41$\pm$4 | … | 0.40$\pm$0.07 | 0.66$\pm$0.06 | … | 4.1$\pm$0.7 | 6.9$\pm$0.6 | …
H$\beta$ |
59178 | -16.3$\pm$3.3 | 71.5$\pm$2.8 | -158.5$\pm$3.6 | 120$\pm$8 | 111$\pm$4 | 123$\pm$5 | 22.3$\pm$1.8 | 24.9$\pm$1.7 | 18.9$\pm$1.9 | 214.5$\pm$17.3 | 239.5$\pm$16.3 | 182.0$\pm$18.3
59179 | 18.7$\pm$2.1 | … | 190$\pm$40 | … | 21.2$\pm$2.0 | … | 163.1$\pm$19.2 | …
59180 | 59.4$\pm$2.6 | -19.3$\pm$3.3 | 171.9$\pm$4.7 | 69$\pm$5 | 116$\pm$3 | 116$\pm$8 | 8.0$\pm$0.9 | 19.1$\pm$1.6 | 6.1$\pm$1.2 | 77.4$\pm$8.7 | 183.8$\pm$15.4 | 58.6$\pm$11.5
59243 | -12.7$\pm$3.5 | 65.6$\pm$2.3 | -124.8$\pm$18.7 | 70$\pm$13 | 78$\pm$6 | 285$\pm$32 | 6.7$\pm$0.7 | 10.8$\pm$0.9 | 3.6$\pm$1.0 | 64.9$\pm$6.8 | 103.9$\pm$8.7 | 34.8$\pm$9.7
H$\alpha$ |
59178 | 47.1$\pm$3.7 | -168.8$\pm$5.3 | 200$\pm$12 | 150$\pm$15 | 60.0$\pm$1.5 | 26.8$\pm$1.1 | 1177$\pm$29 | 525$\pm$22
59179 | 24.1$\pm$3.8 | … | 250$\pm$19 | … | 50.7$\pm$1.7 | … | 895$\pm$33 | …
59180 | 21.6$\pm$3.3 | 210.8$\pm$9.0 | 188$\pm$10 | 131$\pm$21 | 59.8$\pm$1.0 | 9.9$\pm$0.9 | 1174$\pm$20 | 194$\pm$18
59243 | 51.4$\pm$6.6 | -86.2$\pm$14.5 | 152$\pm$11 | 214$\pm$18 | 21.3$\pm$1.0 | 14.5$\pm$1.2 | 419$\pm$20 | 284$\pm$24
Figure 6: Residual UVES spectra in the He i $\lambda$6678 Å region (black dots). In each panel, the two-Gaussian fit to the emission peaks (green full line) and the individual Gaussians corresponding to the primary (blue dashed line) and secondary (red dashed line) component are overlaid. The blue and red vertical ticks mark the expected position of the lines of primary and secondary component, respectively, according to the photospheric RV. The reduced Julian day and the orbital phase ($\phi$) are marked in the upper right corner of each box.  Figure 7: Residual UVES He i $\lambda$5876 Å profiles (black dots). In each panel, the two-Gaussian fit to the emission peaks (green full line) and the individual Gaussians corresponding to the primary (blue dashed line) and secondary (red dashed line) component are overlaid. The blue and red vertical ticks mark the expected position of the lines of primary and secondary component, respectively, according to the photospheric RV. The dashed brown line in the upper panel represents the Gaussian fitted to the excess blueshifted emission observed only in this spectrum. The reduced Julian day and the orbital phase ($\phi$) are marked in the upper right corner of each box. Figure 8: Residual UVES H$\beta$ profiles (black dots). The meaning of the symbols is as in Fig. 7. Table 4: Accretion luminosity and mass accretion rates of the components of CVSO 104 A. RJD | $\log L_{\rm acc}^{a}$ | $\log L_{\rm acc}^{b}$ | $\log\dot{M}_{\rm acc}^{a}$ | $\log\dot{M}_{\rm acc}^{b}$
---|---|---|---|---
| ($L_{\sun}$) | ($M_{\sun}yr^{-1}$)
59178 |  $-1.70\pm$0.18 |  $-1.8\pm$0.30 |  $-8.82\pm$0.18 |  $-8.86\pm$0.30
59179 |  $-1.45\pm$0.22 | …
59180 |  $-2.08\pm$0.24 |  $-1.95\pm$0.31 |  $-9.20\pm$0.24 |  $-8.98\pm$0.31
59243 |  $-2.27\pm$0.28 |  $-2.01\pm$0.29 |  $-9.39\pm$0.28 |  $-9.04\pm$0.29

The [O i] $\lambda$6300 Å line is always observed as a single and rather
symmetric feature (see Fig. 9). Its radial velocity spans from +19 to +26 km
s-1, therefore it is always close to the barycentric velocity. Its equivalent
width in the UVES spectra does not change very much, being about 0.80, 0.86,
1.09, and 1.20 Å, from the first to the last epoch of UVES observations.
Interestingly, it is slightly stronger when the permitted lines are weaker.
The [O i] $\lambda$6363Å line displays a similar behavior. This does not
necessarily mean a real intensity variation, but it could be instead the
result of a decrease of the excess continuum flux due to the accretion, viz.,
the veiling.

The observed profile of the [O i] $\lambda$6300 Å line displays broad wings
that cannot be reproduced with a single Gaussian. We have therefore
interpreted them as the superposition of a narrow and a broad component, which
have been fitted with Gaussians. The broad component is normally associated
with magneto-hydrodynamic winds from the inner ($<1$ au) disk, while the
narrow one is related to photoevaporated winds originating in a much more
extended region (up to 100 au, e.g. Ercolano & Pascucci 2017). The parameters
of the narrow (N) and broad (B) emission component are reported in Table 5.
The narrow emission component has an average width of FWHM$\simeq$ 44 km s-1,
while it is about 288 km s-1 for the broad one. Following the prescriptions
applied in different studies of forbidden emission lines (Simon et al. 2016;
McGinnis et al. 2018; Fang et al. 2018; Banzatti et al. 2019; Gangi et al.
2020) we can estimate the emission size of these components under the
assumption that the line widths are dominated by Keplerian broadening.
Assuming an inclination of $i=43$ ∘ and a total mass of 1 $M_{\sun}$ (Sect.
3.5), the narrow component should be emitted by a region of $\approx$ 0.84 au
(180 $R_{\sun}$), which is compatible with a circumbinary disk, as also found
from the SED analysis. The size of the source of the broad component should be
$\approx$ 0.02 au (4.3 $R_{\sun}$).

Figure 9: Residual UVES spectra in the [O i] $\lambda 6300$ region (black
dots). The blue and red vertical ticks mark the expected position of the lines
of primary and secondary component, respectively, according to the
photospheric RV. The Gaussian fits to the broad and narrow component of the
line profile are overplotted with purple and magenta lines, respectively,
while their sum is plotted with a green line. The reduced Julian day and the
orbital phase ($\phi$) are marked in the upper right corner of each box.
Table 5: Multicomponent analysis of the [O i] $\lambda$6300 Å emission line.

RJD | RVN | FWHMN | $I_{N}$ | RVB | FWHMB | $I_{B}$
---|---|---|---|---|---|---
| (km s-1) |  | (km s-1) |
59178 | 21.0$\pm$3.7 | 46.9$\pm$7.8 | 0.46(6) | 11.1$\pm$18.9 | 355$\pm$84 | 0.12(2)
59179 | 21.6$\pm$3.3 | 38.7$\pm$6.9 | 0.58(8) | 39.6$\pm$12.6 | 267$\pm$51 | 0.19(3)
59180 | 22.0$\pm$2.0 | 43.8$\pm$4.0 | 0.67(5) | 32.8$\pm$13.4 | 289$\pm$50 | 0.12(2)
59243 | 22.8$\pm$1.6 | 37.9$\pm$3.4 | 0.93(8) | 48.3$\pm$10.8 | 258$\pm$43 | 0.18(3)

Notes. $I_{N}$ and $I_{B}$ are the intensities of the narrow and broad
component, respectively, in units of the continuum.

###  3.5 Spectral energy distribution

Figure 10: X-Shooter flux-calibrated spectrum of CVSO 104 A (pink line), CVSO
104 B (cyan line), and the sum of the two (grey line) in the visual (top
panel) and NIR region (bottom panel). The contemporaneous photometry obtained
from OACT in the $BVR_{\rm C}I_{\rm C}z^{\prime}$ bands and the 2MASS
$JHK_{\rm s}$ photometry are displayed with colored dots. The filter
bandpasses and the synthetic photometry (open squares) obtained integrating
the combined spectrum (A+B) over these bandpasses are overlaid with the same
color as the photometric points.

To study the shape of the spectral energy distribution (SED) of CVSO 104 A,
excluding the contribution of the optical companion, we used both the OACT
photometry and synthetic photometry made on the flux-calibrated X-Shooter
spectra of the two stars (see Fig. 10). We extended the SED to mid-infrared
(MIR) and far-infrared (FIR) wavelengths by adding flux values from the
literature. These data are quoted in Table 6. As the contribution of the
visual companion B is strongly decreasing at longer wavelengths (see Fig. 10),
we assign the MIR (WISE) and FIR (Herschel) fluxes to the component A.

Figure 11: Spectral energy distribution of CVSO 104 A based on OACT and
synthetic photometry made on the X-shooter spectrum. Gaia magnitudes, mid- and
far-infrared fluxes are shown with different symbols, as indicated in the
legend. The combination of BT-Settl spectra (Allard et al. 2012) that
reproduce the photospheres of the components of the close binary is shown by a
gray line. The two black bodies with $T=1000$ K and $T=150$ K that fit the MIR
and FIR disk emission are shown by the dotted red and dashed green lines,
respectively. The continuous blue line displays the sum of the smoothed
photospheric template and the two black bodies. The light-grey shaded area in
the background is the median SED of Class II Taurus sources according to
D’Alessio et al. (1999) and Furlan et al. (2006), reddened by $A_{V}$=0.35
mag, and scaled to the $H$ band photometric point.

To reproduce the photospheres of the two components of the binary, we combined
two BT-Settl spectra (Allard et al. 2012) adopting the temperatures and flux
contributions at red wavelengths, $w_{\rm a}$ and $w_{\rm b}$, found with
COMPO2 and reported in Table 2. With this photospheric template, we fitted the
optical-NIR portion (from $B$ to $H$ band) of the SED (Fig. 11) fixing the
Gaia EDR3 parallax and letting the extinction $A_{V}$ and the radius of the
primary component, $R_{\rm a}$, free to vary until a minimum $\chi^{2}$ is
attained. We found $A_{V}=0.35$ mag and $R_{\rm a}=1.07R_{\sun}$. The 3D
extinction map of the Galaxy published by Green et al. (2019) provides
$E(g-r)=0.12$ mag at the position and distance of CVSO 104
A111111http://argonaut.skymaps.info/query, which translates into
$E(B-V)=$0.106–0.120 mag, depending on the conversion used, which would then
correspond to $A_{V}=0.33$–0.37 mag. This value is in close agreement with the
value of $A_{V}$ derived by us and suggests that most (or all) of this
reddening is interstellar in nature, rather than circumstellar. The radius of
the secondary component, $R_{\rm b}=0.99R_{\sun}$, is derived from $R_{\rm a}$
and the flux contributions $w_{\rm a}$ and $w_{\rm b}$. We have then evaluated
the stellar luminosities as $L=4\pi R^{2}\sigma$$T_{\rm eff}$4. The error of
luminosity includes also the error on flux contribution. These stellar
parameters are also quoted in Table 2. The position of the two components of
CVSO 104 A in the Hertzsprung-Russell (HR) diagram is shown in Fig. 12 along
with the pre-main-sequence evolutionary tracks and isochrones by Baraffe et
al. (2015). Both components lie close to the isochrone at 5 Myr and masses of
0.57$\pm$0.15 $M_{\sun}$ and 0.43$\pm$0.15 $M_{\sun}$ can be inferred for the
primary and secondary component, respectively. Comparing these values with the
dynamical masses reported in Table 2, we deduce a system inclination $i\simeq
43^{+8}_{-4}$ degrees. We note that the mass ratio derived from the position
in the HR diagram, $M_{\rm b}/M_{\rm a}=0.75\pm 0.33$, is smaller than, but
still compatible with the dynamical one, if we take the large errors of the
individual masses ($\pm 0.15M_{\sun}$) into account. The latter are mainly the
result of the large $T_{\rm eff}$ uncertainty. Therefore, the radial velocity
curve strongly suggests components with closer masses and effective
temperatures.

Figure 12: Position of the primary (blue dot) and secondary (red dot)
component of CVSO 104 A in the HR diagram. Isochrones and evolutionary tracks
by Baraffe et al. (2015) are overlaid as dashed and solid lines; the labels
represent their age and mass. The green asterisk connected with green dotted
lines to the dots marks the position of CVSO 104 A with the parameters derived
from the combined spectrum in Paper I.  Figure 13: Schematic representation of
the geometry of the CVSO 104 A system in the orbital plane. The dark shaded
surfaces represent the primary (at left) and secondary component, while the
light-grey one is the critical Roche surface. The size and separation of the
stars and Roche lobes as well as the position of the barycenter (MC) and the
Lagrangian point $L_{3}$ are to scale. The source of the excess H$\alpha$ and
H$\beta$ emission is represented by the shaded green area.

The masses and radii derived for the components Aa and Ab of CVSO 104 enable
to trace the configuration of the system, which is depicted in Fig. 13.

The SED also shows that up to the $H$ band the emission is photospheric, while
longward of that it follows the typical emission pattern displayed by Class II
sources in Taurus (e.g., D’Alessio et al. 1999; Furlan et al. 2006), which is
outlined by the light-grey shaded area in Fig 11. Such an IR excess could be
explained by thermal emission from a dusty disk. We can, for example,
reproduce it with two blackbodies with $T_{1}$=1000 K (red dotted line in Fig.
11) and $T_{2}$=150 K (green dashed line). The areas of these sources, assumed
as uniformly emitting, are about 1.0$\times 10^{25}$ cm2 and 2.1$\times
10^{28}$ cm2, respectively. If they are related to an inner and outer disk
with inclination $i=43$∘ and internal radius $R_{\rm in}=1\,R_{\sun}$, their
radii should be $\sim 21$ and 960 $R_{\sun}$. If compared with the semi-major
axis $a\simeq 13^{+1}_{-2}\,R_{\sun}$, this implies that a circumbinary disk
must exist.

##  4 Discussion

Our study confirms that CVSO 104 A is a double-lined spectroscopic binary with
a dynamical mass ratio around 0.92, that is slightly lower but marginally
consistent within the uncertainties with the estimate by Kounkel et al.
(2019). With an orbital period of 5 days, but no physical companion detected
directly, by Gaia nor by Broadening Function method, the system is a rare case
of a close binary without a detected tertiary component. Indeed, Laos et al.
(2020) found that 90% of spectroscopic binaries with orbital periods between 3
and 6 days had a tertiary companion.

Observations with TESS reveal variability, with an irregular, or at most semi-
regular, pattern. The fact that TESS has large pixels, so that the component B
is also contributing to the flux, makes difficult assessing the real amplitude
of the variations and whether they really originate from the spectroscopic
binary. However, according to our classification as a slowly-rotating, low-
activity G2 V star and the photometry displayed in Fig. 15, the brightness
variations of the component B are negligible in comparison with CVSO 104 A.
The observed variations look more like small-amplitude bursts (perhaps linked
to accretion) than dips, especially in the second TESS light curve,
simultaneous with PENELLOPE. The TESS light curves are reminiscent of the
sources in NGC 2264 with “aperiodic accretion variability” or “stochastic
variability” presented by Stauffer et al. (2016) or to the “stochastic” or
“burster” type of variability detected with $K2$ in Upper Sco (Cody &
Hillenbrand 2018). With the aim of distinguishing between a “dipper” or
“bursting” light curve on more quantitative grounds, we have evaluated the
asymmetry of the TESS light curve using the metric $M$ expressed by Cody et
al. (2014) in their Eq. 7. After removing the linear trend, we found $M\simeq$
0.13. According to Cody et al. (2014), values of $M<-0.25$ are typical of
“bursting” light curves, $M>+0.25$ correspond to “dippers”, while
$-0.25<M<+0.25$ indicates symmetric light curves. As a further test, we
calculated the third moment (skewness) as $-0.45$, which also indicates a
rather symmetric light curve. We repeated this analysis with the ASAS-SN
$g^{\prime}$ data, which have a precision and cadence lower than TESS, but a
longer time baseline and we found $M\simeq$ 0.17 and skewness $-0.30$, which
also suggest a symmetric light curve. Therefore, we cannot classify CVSO 104 A
either as a “burster” or a “dipper”. The stochastic variability is probably
related to fluctuating accretion. However, our multiband photometry (Fig. 14)
shows at least two clearly detected bursts simultaneous with the enhancements
of the TESS light curve. The intensification is stronger in the bluer bands,
as usually observed during accretion bursts (see, e.g., Tofflemire et al.
2017).

In Paper I we quote a mass for the star of 0.37 $M_{\odot}$ and a mass
accretion rate of $3.24\times 10^{-9}M_{\odot}$yr-1, but these values assumed
that the object was single. The mass accretion rate reported in Paper I is
higher than the $\dot{M}_{\rm acc}$ values of each component, but it is close
to their sum at RJD=59178, $\dot{M}_{\rm acc}$ $\simeq 2.95\times
10^{-9}M_{\odot}$yr-1.

Using evolutionary tracks and isochrones of Baraffe et al. (2015), we found
that both components are near the isochrone at 5 Myr and derived the masses to
be 0.57$\pm$0.15 $M_{\sun}$ and 0.43$\pm$0.15 $M_{\sun}$, with a binary
inclination of $i\simeq 43^{\circ}\pm 6^{\circ}$ (see Fig. 12). However, apart
from the large mass errors, which mainly stem from the large uncertainties of
$T_{\rm eff}$ values, the masses depend on the adopted set of evolutionary
tracks. To evaluate the impact of different models on the determination of
masses and age, we have also used the SPOTS Models of Somers et al. (2020),
which include the effects of magnetic activity and starspots on the structure
of active stars. The HR diagrams with the Somers et al. (2020) tracks for a
spot covering fraction of 0%, 34%, and 51% are shown in Fig. 18. We note that
the position of the two components with respect to the tracks with a spot
filling factor $F_{\rm spot}=0$ is the same as for the tracks of Baraffe et
al. (2015). A spot coverage of 34% would increase the masses by a factor of
$\sim 1.3$, which is inside the mass errors, while a larger spot filling
factor, $F_{\rm spot}=51\%$, would make the masses larger by a factor $\sim
1.5$ and the inclination of the system would be $i\simeq 37$∘. However, unless
the spots are evenly distributed over the surface of the two stars, there is
no indication of such a high coverage from the light curve which should have
shown a large-amplitude rotational modulation superimposed to the observed
bursts. Moreover, in the latter case, the age of the system would turn out to
be 15–20 Myr, which is hardly compatible with the high mass accretion rate and
with the age of $\sim$ 5 Myr of the Ori OB1 association.

If we use the radii and orbital period listed in Table 2 we find equatorial
rotation velocities of $v_{eq}=10.7$ and 10.0 km s-1for the primary and
secondary components, respectively, which become $v\sin i$=7.3 and 6.8 km s-1,
i.e. in agreement with the measured $v\sin i$, within the errors. We therefore
conclude that the components are synchronized or close to synchronization, but
the system is not yet circularized. This is in agreement with the time scales
for circularization and synchronization, $\tau_{\rm circ}\sim 800$ Myr and
$\tau_{\rm sync}\sim 2$ Myr, which we have calculated according to Zahn
(1989). With an age of $\sim 5$ Myr, this system should have had time enough
to synchronize the rotation of the two components with the orbital period.
However, as pointed out by Hut (1981), in an eccentric orbit the tidal
interaction is stronger at periastron, when the orbital velocity is higher,
with the consequence that the equilibrium is reached at a value of rotation
period, $P_{\rm pseudo}$, which is smaller than $P_{\rm orb}$, leading to a
pseudo-synchronization. The value of $P_{\rm pseudo}$ depends on the
eccentricity of the system and results to be about 2.5 days for both
components of CVSO 104 A, which is not consistent with the $v\sin i$ values
measured by us. The timescale for the pseudo-synchronization can be evaluated
as $\tau_{\rm pseudo}\sim 12$ Myr, following the guidelines of Hut (1981).
This suggest that the pseudo-synchronization equilibrium has not yet been
attained for the components of CVSO 104 A.

With a semi-major axis of about 13 R⊙ and an inclination of about 43$\degr$,
the system should certainly not show any eclipses. In all cases, the Roche
lobe radius of both stars is about 5 R⊙, that is, much larger than the radius
of the individual stars, which is about 1 R⊙ (see also Fig. 13). There is thus
also place for circumstellar accretion discs, in addition to a circumbinary
disc. The existence of the latter is confirmed by our SED analysis.

The accretion luminosities and mass accretion rates of the two components,
calculated from the fluxes of the lines where the profiles of the two
components could be deblended (He i$\lambda$6678, He i$\lambda$5876, and
H$\beta$), are similar. The secondary component seems to accrete slightly more
than the primary, but at the first epoch (RJD= 59178) we see the reverse.
However, these variations are within the errors and cannot be considered as
highly significant.

The numerical simulations of circumbinary accretion onto eccentric and
circular binaries made by Muñoz & Lai (2016) show that for circular binaries
one expects accretion bursts with a period $\sim 5P_{\rm orb}$, while for
eccentric orbits it is mostly modulated at $\sim 1P_{\rm orb}$. This seems to
be the case with CVSO 104 A, according to what is observed in the TESS light
curve (Fig. 1). Moreover, based on the above simulations, the two components
should have similar accretion rates in circular orbits, while very different
accretion rates, with a ratio up to 10–20, should be observed in the
components of binaries with eccentric orbits. The asymmetry breaking between
the stars, however, alternates over timescales of the order of 200 $P_{\rm
orb}$ and it can be attributed to a slowly precessing, eccentric circumbinary
disk. Our spectroscopic observations always display very similar $\dot{M}_{\rm
acc}$ for the two components, with some hint of variation, but they span a too
small time range and further observations are needed to search for variations
on 100-day timescales.

Figure 1 of Muñoz & Lai (2016) clearly shows the spiral pattern of matter from
the circumbinary disk to the stars/circumstellar disks. Similar structures are
found in the simulations of Gillen et al. (2017) and de Val-Borro et al.
(2011) for the two gas streams passing the co-linear Lagrangian points on the
way down to the stars. de Val-Borro et al. (2011) also used their model to
estimate the RV of Balmer line emission from the region within the
circumbinary disk, finding maximum velocities of about $\pm 100$ km s-1 around
the quadratures, in agreement with the excess emission observed on V4046 Sgr
by Stempels & Gahm (2004). According to Stempels & Gahm (2004), these excess
emissions are consistent with two concentrations of gas co-rotating with the
stars and moving with a projected velocity of 80 km s-1 around the center of
mass. This suggests that they are located well inside the edge of the
circumbinary disk, and also inside the co-linear Lagrangian points of V4046
Sgr. It is possible that a similar accretion structure that brings matter
towards the primary component is responsible for the extra emission component
observed in H$\alpha$ and H$\beta$. We note that its RV is approximately
symmetric with respect to barycentric velocity $\gamma=24.5$ km s-1. If we are
looking at the same structure at opposite phases (near quadratures) and if we
assume it as quasi-stationary in the reference frame rotating with the system,
it should be located at $\approx 15R_{\sun}$ from the barycenter, i.e. near
the Lagrangian point $L_{3}$ (see Fig. 13). The absence of a similar structure
on the side of the secondary star is not an unexpected result, since the
simulations of de Val-Borro et al. (2011) and Muñoz & Lai (2016) show that the
density distribution in the inner gap can be highly asymmetrical. For
instance, the model of Terquem et al. (2015) applied to CoRoT 223992193, a PMS
binary with $P_{\rm orb}\simeq 3.87$ day and M-type components, shows that the
stream of matter from the circumbinary disk to the primary is much denser than
the one directed toward the secondary component. A similar result is also
reported by Gómez de Castro et al. (2020) for simulations of AK Sco, a short-
period SB2 composed of mid-F stars with a circumbinary disk. They observed an
enhancement of the accretion rate during the periastron passages from UV
tracers and higher resolution COS spectra revealed that the flow was channeled
preferentially into one of the two components.

Contrary to what observed by Gómez de Castro et al. (2020) for AK Sco, Ardila
et al. (2015) do not detect any clear correlation between accretion luminosity
and phase for the two binaries DQ Tau and UZ Tau either from UV continuum or C
iv line flux, suggesting that gas is stored in the system throughout the orbit
and may accrete stochastically. These systems are both composed of early-M
type stars in eccentric orbits with longer orbital periods ($P_{\rm orb}=15.8$
and 19.1 day, respectively) and circumbinary disks. However, optical emission
lines and continuum veiling intensification near the periastron passage have
been reported for DQ Tau (e.g., Mathieu et al. 1997; Basri et al. 1997).
Kóspál et al. (2018) analyzed contemporaneous ground-based, Spitzer, and $K2$
photometry of DQ Tau. The $K2$ light curve displays a clear rotational
modulation which is crossed by short-duration flare-like events but also by
stronger brightening events with a complex behaviour and a longer duration,
which are observed only at phases close to the periastron passage. The latter
are interpreted as accretion bursts. Tofflemire et al. (2017) observed
accretion bursts near periastron passages in the same binary system with
multiband photometry which are stronger in the bluer bands. We observed a
similar color behaviour for the bursts observed in CVSO 104 A, but they are
not seen near periastron passages. Another system for which accretion burst
have been observed near periastron passages is TWA 3A (Tofflemire et al.
2019). This is apparent from their $U$-band light curve, but it is also
indicated by the intensity and FWHM of H$\alpha$, H$\beta$ and He i
$\lambda$5876 lines, which increase near the periastron passages. They also
show that the He i $\lambda$5876 line indicates that the primary is accreting
more than the secondary and suggest that this can be explained by the Muñoz &
Lai (2016) simulations. We note, however, that $P_{\rm orb}$ and eccentricity
are larger for TWA 3A than for CVSO 104 A, with periastron and apastron
separations of 14.7 and 64.0 $R_{\sun}$, respectively. Therefore the effect of
periastron passage may be stronger for the former system. We need time-series
high-resolution spectroscopy, possibly in different epochs, to investigate the
behavior of accretion variability in CVSO 104 A.

##  5 Conclusions

We presented a spectroscopic and photometric study of the PMS object CVSO 104
A, which is located in the Ori OB1 association and has an optical companion
with a similar brightness at about 2$\aas@@fstack{\prime\prime}4$. The latter
star, which we have named CVSO 104 B, is a background Sun-like star not
physically associated with the PMS object and does not belong to Ori OB1.
Thanks to high- and intermediate-resolution spectra taken in the framework of
the PENELLOPE large program and archival APOGEE spectra, we confirmed CVSO 104
A as a double-lined spectroscopic binary and derived, for the first time, its
orbital parameters. We found a dynamical mass ratio of 0.92, an orbital period
of about 5 days, and an eccentric orbit ($e\simeq 0.39$, see Table 2). It is a
rare case of a close binary without a detected tertiary component.

The analysis of the UVES spectra allowed also us to estimate the parameters of
the two components of the binary system, which turn out to be slowly-rotating
($v\sin i$ $\simeq 6-7$ km s-1), early M-type stars. In the HR diagram both
stars are located near the same isochrone at 5 Myr (according to two different
models), in good agreement with the age estimated for Ori OB1. The rotation
rates indicate that the two stars have already attained spin-orbit
synchronization but the orbit is not yet circular. This agrees with the
timescales for synchronization and circularization that we have estimated as
$\tau_{\rm sync}\sim 2$ Myr and $\tau_{\rm circ}\sim 800$ Myr, respectively.

The SED shows a significant infrared excess that can be explained only with a
circumbinary accretion disk with an extension of at least 4.5 au (960
$R_{\sun}$), although the presence of circumstellar disks around the two stars
cannot be ruled out. The kinematic properties (RV and FWHM) of the narrow
component of the [O i] $\lambda$6300Å line are also compatible with emission
from a region of a circumbinary disk of $\approx$ 1 au.

The analysis of permitted lines, such as H$\beta$, He i $\lambda$5876 Å and He
i $\lambda$6678 Å, after the removal of the underlying photospheric spectrum,
clearly reveals emission from both components of the close binary CVSO 104 A
with a similar intensity. This result is in contrast with what is suggested by
some theoretical studies, which predict, for most of the time, very different
accretion rates for the components of an eccentric binary system, even if they
have a similar mass (e.g., Muñoz & Lai 2016). However, since the ratio of
accretion rates of the components is expected to reverse periodically, a
binary system should be observed nearly-continuously for a long time ($\gtrsim
100$ d) to draw firm conclusions. In addition to the emission profiles
corresponding to the velocity of the components and likely produced by
emitting material near the accretion shocks, the H$\alpha$ and H$\beta$
profiles display a broad excess emission component, which appears blueshifted
or redshifted by more than 100 km s-1 at different phases. We think that this
emission could be produced by an accretion funnel from the circumbinary disk
towards the primary component, similar to what found by Stempels & Gahm (2004)
for V4046 Sgr and in agreement with the predictions by numerical simulations
(e.g., de Val-Borro et al. 2011; Terquem et al. 2015; Gillen et al. 2017).

The contemporaneous space- (TESS) and ground-based photometry displays a
stochastic variability pattern with a possible periodicity around 4.7 days
(not far from the orbital period) and some short-duration ($\approx$ 1 day)
peaks. The latter are reminiscent of accretion bursts, based on their shape
and color dependence. However, unlike other binary systems with eccentric
orbits (such as the paradigmatic case of DQ Tau), these accretion bursts do
not seem to occur near the periastron passages.

Future studies with multi-epochs and high-resolution spectroscopy of this and
other PMS binaries are needed for a deeper investigation of the impact of
multiplicity on the mass accretion phenomenon. The synergy between ODYSSEUS
and PENELLOPE large programs can be of great help in this respect and the data
are publicly
available121212https://zenodo.org/communities/odysseus/?page=1&size=20.

###### Acknowledgements.

We thank the anonymous referee for her/his useful comments and suggestions. We
acknowledge the support from the Italian Ministero dell’Istruzione, Università
e Ricerca (MIUR). This work has been partially supported by the project PRIN-
INAF-MAIN-STREAM 2017 “Protoplanetary disks seen through the eyes of new-
generation instruments” and from the project PRIN-INAF 2019 ”Spectroscopically
Tracing the Disk Dispersal Evolution”. This work benefited from discussions
with the ODYSSEUS team (HST AR-16129), https://sites.bu.edu/odysseus/. This
project has received funding from the European Union’s Horizon 2020 research
and innovation programme under the Marie Skłodowska-Curie grant agreement No
823823 (DUSTBUSTERS). This project has received funding from the European
Research Council (ERC) under the European Union’s Horizon 2020 research and
innovation programme under grant agreement No 716155 (SACCRED). This work was
partly supported by the Deutsche Forschungs-Gemeinschaft (DFG, German Research
Foundation) - Ref no. FOR 2634/1 TE 1024/1-1. FMW is grateful to the AAVSO for
the award of AAVSOnet telescope time, and to the Ken Menzies for implementing
the observations. We thank Elizabeth Waagen and the citizen scientists of the
AAVSO for their contributions to this program. K.G. acknowledges the partial
support from the Ministry of Science and Higher Education of the Russian
Federation (grant 075-15-2020-780). This research made use of SIMBAD and
VIZIER databases, operated at the CDS, Strasbourg, France, and of Lightkurve,
a Python package for Kepler and TESS data analysis (Lightkurve Collaboration
et al. 2018).

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##  Appendix A Additional tables and figures

Figure 14: Multiband ground-based light curve of CVSO 104 A+B around the time
of VLT and HST observations. Different symbols, as reported in the legend, are
used for the different data sets. The H$\alpha$ color index, which measures
the intensity of the line, is shown in the upper panel. Synthetic H$\alpha$
colors based on X-Shooter and UVES spectra are overplotted with asterisks in
the same box. The contemporaneous TESS light curve is overplotted with grey
dots to the $R_{\rm C}$ light curve. The arrows in the same box mark the
periastron passages. The epoch of VLT and HST observations are marked in the
lower box.  Figure 15: Multiband ground-based photometry of CVSO 104 B
obtained at OACT and Konkoly observatories in the nights with a good seeing.
Table 6: Photometry of CVSO 104 A and B.

Band | $\lambda_{\rm c}$ | Comp. A | Comp. B | Reference
---|---|---|---|---
| ($\mu$m) | (mag) | (mag) |
$B$ | 0.444 | 16.67$\pm$0.08 | 15.67$\pm$0.06 | Present work
$V$ | 0.550 | 15.57$\pm$0.05 | 14.87$\pm$0.04 | Present work
$R_{C}$ | 0.621 | 14.64$\pm$0.05 | 14.47$\pm$0.05 | Present work
$I_{C}$ | 0.767 | 13.53$\pm$0.04 | 14.10$\pm$0.07 | Present work
$z^{\prime}$ | 0.910 | 13.56$\pm$0.05 | 14.43$\pm$0.11 | Present work
$BP$ | 0.505 | 15.242$\pm$0.042 | 15.024$\pm$0.004 |  Gaia EDR3
$G$ | 0.623 | 14.455$\pm$0.009 | 14.588$\pm$0.003 |  Gaia EDR3
$RP$ | 0.772 | 13.398$\pm$0.018 | 13.932$\pm$0.005 |  Gaia EDR3
$J$ | 1.235 | 12.16$\pm$0.05 | 13.52$\pm$0.08 | Present work
$H$ | 1.662 | 11.24$\pm$0.05 | 13.03$\pm$0.08 | Present work
$K_{\rm s}$ | 2.159 | 10.81$\pm$0.04 | 13.23$\pm$0.08 | Present work
$WISE\,1$ | 3.35 |  9.699$\pm$0.023 | … | WISE
$WISE\,2$ | 4.60 |  9.204$\pm$0.020 | … | WISE
$WISE\,3$ | 11.56 |  7.560$\pm$0.019 | … | WISE
$WISE\,4$ | 22.09 |  5.571$\pm$0.039 | … | WISE
$Herschel\,blue$ | 70 | 84.7$\pm$2.7 mJy | … | H2020
$Herschel\,red$ | 160 | 72.9$\pm$14.0 mJy | … | H2020

Notes Gaia EDR3 = Gaia Collaboration (2020); WISE = Cutri et al. (2021); H2020
= Herschel Point Source Catalogue Working Group et al. (2020).

Figure 16: UVES spectrum of CVSO 104 B (black dots) in three spectral regions
around the Mg i b triplet (top), H$\alpha$ (middle), and 6700 Å (bottom). In
each box the spectrum of the standard star that is best fitting that of CVSO
104 B is overplotted with a full red line.  Figure 17: Residual UVES H$\alpha$
profiles (black dots). The totally blended emission from the two stars is
represented, in each box, by a magenta dashed line, while the brown dashed
line is the Gaussian fitted to the excess blueshifted or redshifted emission.
The sum of the two Gaussians is overplotted as a full green line. The blue and
red vertical ticks mark the expected position of the lines of primary and
secondary component, respectively, according to the photospheric RV.

Figure 18: HR diagram of the primary (blue dot) and secondary (red dot)
component of CVSO 104 A with isochrones and evolutionary tracks by Somers et
al. (2020) for three spot covering factors.