their detection of the PMA has a low SNR=3.78 (the SNR is even much lower for
the 1991.25 epoch). The PMA predicted from the relative orbit depends on the
assumed period (which is not well determined); since our observations were
obtained not far from the reference epoch, the main effect is the contribution
of the orbital motion to the estimate of the long term trend by Kervella et
al. (2019). The orbital solution that gives the best agreement with Kervella
et al. (2019) results is for a semi-major axis of 86 mas (=11.0 au) and a
period of 20.3 yr, longer though within the error of the orbit fitting our
astrometry alone. In this case the values we obtain for the primary motion
anomaly at this epoch are about 5.88 mas/yr in RA and 4.33 mas/yr in
declination. While the direction of motion is well reproduced, we would still
expect a larger motion anomaly (by a factor of $\sim 1.5$) for the system
photocenter than observed by Kervella et al. (2019). This difference might be
due to an underestimate by us of the correction from the photocenter to
primary motion anomaly, e.g. because the luminosity difference between the two
components in the Gaia photometric system is smaller than we assumed ($\sim
0.7$ mag rather than $\sim 0.9$ mag). On the whole, we conclude that these
comparisons support the orbit determination presented in this paper.

Figure 9: Posterior distribution of the orbital parameters (from left to
right: semi-major axis, eccentricity, inclination, longitude of node,
longitude of periastron, time of periastron passage) for TYC 6872-1011-1
obtained with the Monte Carlo method.

####  4.4.3 Statistical constraints on orbits for systems with few
observations

For the binary systems with multiple relative astrometric observations
covering only a small fraction of the orbital motion, and therefore unsuitable
for a derivation of orbital parameters using Orbit, we ran a Monte Carlo
simulation to explore the possible families of orbits allowed. The procedure
(hereafter MC) follows Zurlo et al. (2018): we created $2*10^{7}$ orbits with
random orbital parameters and selected only the ones that fitted the
astrometric points. The fitting procedure is based on the visual binaries
constants of Thiele-Innes. In this way we can understand whether the posterior
distributions of the orbital parameters for a given target are uniform, or
whether certain orbits are preferred. Figure 8 shows that useful constraints
were obtained for 7 systems where roughly 10% of the orbit is covered using
this approach, though in favourable conditions even a shorter coverage may
give some hints. An example of constraints used in this procedure is shown in
Figure 9. The results of the MC for the orbits selected are listed in Table
LABEL:tab:orbits. To estimate the error bars of the parameters summarised in
that table, we calculate the 0.16 and 0.84 quantiles of each posterior
distribution. The median value (0.5 quantile) is assumed as the most probable
value. Note that some posterior distributions are more stringent, while others
permit a wide range of values for the parameters. That is normal for a very
small coverage of the orbit.

###  4.5 Sensitivity to additional companions

Figure 10: Probability of detecting additional companions around the stars in
our sample, calculated using the Exoplanet Detection Map Calculator (Exo-DMC,
see Bonavita 2020 for details) and the median mass limits shown in Fig. 3.

We used the Exoplanet Detection Map Calculator (Exo-DMC: Bonavita
2020)666https://ascl.net/2010.008 to obtain a first estimate of the
completeness of our sample in terms of additional stellar companions within
the IFS field of view. The Exo-DMC is the latest (and for the first time in
Python) rendition of the MESS (Multi-purpose Exoplanet Simulation System
Bonavita et al. 2012), a Monte Carlo tool for the statistical analysis of
direct imaging survey results. In a similar fashion to its predecessors, the
DMC combines the information on the target stars with the instrument detection
limits to estimate the probability of detection of companions in a given mass
and semi-major axis range, ultimately generating detection probability maps.

For each star in the sample the DMC produces a grid of masses and physical
separations of synthetic companions, then estimates the probability of
detection given the provided detection limits. In order to account for the
chances of each synthetic companion to be in the instrument’s field of view, a
set of orbital parameters is generated for each point in the grid, which
allows an estimation of the range of possible projected separations
corresponding to each value of semi-major axis. The default setup uses a flat
distribution in log space for both the mass and semi-major axis with all the
orbital parameters uniformly distributed except for the eccentricity, which is
generated using a Gaussian eccentricity distribution with $\mu=0$ and
$\sigma=0.3$, following the approach by Hogg et al. (2010) (see Bonavita et
al. 2013, for details). The detection probability at a given semi-major axis
is then calculated as the fraction of orbital sets that, for a given mass,
allows for the companion to be detected.

Fig. 10 shows the median detection probability for companions with masses over
70 $M_{\rm Jup}$ and separations between 1 and 100 au, obtained using the
limits shown in Fig. 3. Instead of the default setup, for this specific case
we used the mass and semi-major axis distributions from Raghavan et al.
(2010). Separate runs were performed for the targets for which both
coronographic and non-coronographic images were available, then combined
considering the best performance at each point in the grid.

##  5 Discussion

###  5.1 Survey completeness and binary frequency

Figure 11: Mass ratio ($M_{B}/M_{A}$) VS separation for the binaries in our
sample, marked with circles for the young moving group members, triangles for
the young field objects and squares for the members of the Scorpius Centaurus
region (US, LCC and UCL in Tab. 6). The new systems discovered in this paper
are marked by filled symbols. The plus signs show the position of the
additional objects listed in Tab. LABEL:tab:known_bin. Different colours
represent different ranges of primary masses.

We assume our sample to be reasonably complete at separations between
0.05$\arcsec$ and 0.5$\arcsec$, thus including systems that would have been
too close to be resolved in past observations (but see later discussion), but
still sufficiently wide that any stellar companion should have been detected
in our data. This is confirmed by the limits in Fig. 3 and 10, which show that
our sensitivity in this separation range is well below the Hydrogen Burning
Limits for most of our targets. As a consequence, even at the lowest value of
the mass ratio compatible with a stellar secondary, we expect that only a
small fraction of the stellar secondaries should have been missed in this
range of separations. Given the high level of completeness, this reduced
sample could then be used to draw some preliminary conclusion on the impact of
our results on the frequency and properties of young binaries.

In order to properly do so, however, it is first necessary to correct for the
effect of the initial selection biases of the SHINE Survey. As previously
mentioned, in fact, we removed any objects with known visual companions within
the field of view at the time of selection. Several additional systems were
also removed mid-way through the survey, following the publication of
dedicated works characterising the binarity of stars in young associations
(see e.g. Elliott et al. 2016).

We were able to retrieve the information about the objects which were part of
the original list of members of nearby moving groups ($\beta$ Pic=BPIC,
Tucana=TUC, Columba=COL, Carina=CAR, AB Doradus=ABDO, Argus=ARG: see Tab. 6
for a full description and age references.), and excluded from the final
target list because of their binarity. This list includes 62 systems, of which
20 have $0.05\arcsec<\rho<0.5\arcsec$ and are listed in Tab.
LABEL:tab:known_bin. We note that this list also includes a handful of objects
that were observed within SHINE with special status (TWA 5) or as part of the
binary filler program (HIP 25647 and GJ 2060). Without the bias against
binaries all these objects would have all been part of the SHINE statistical
sample, but not all of them would have been observed and therefore included in
our sample of new binaries, because not all the original targets of the survey
were actually observed. To take this into account, we derived for each of
these rejected objects the value of the merit function described in Desidera
et al. (2021) that was used to assign each target to the four priority bins
considered in the survey; and then derived their probability of being
observed, based on the fraction of objects observed within SHINE for each of
those bins (75.5% for P1, 36.5% for P2, 26.5% for P3 and 18.5% for P4). Most
of the excluded systems were originally classified as P1 (53 out of 62), and
only 6, 1 and 2 were marked as P2, P3 and P4 respectively. They would
therefore count as 42.84 additional detected binaries, of which 13.91 have
$0.05\arcsec<\rho<0.5\arcsec$.

The SHINE statistical sample also includes a number of objects belonging to
the Upper Scorpius (USco), Lower Centaurus Crux (LCC) and Upper Centaurus-
Lupus (LCC) regions (ScoCen). A lot more information about the multiplicity of
these objects would have been available at the time of the target selection,
making the bias against binarity a lot more effective. For example, most of
the early type stars in this region have Hipparcos observations, and therefore
a higher sensitivity to similar-luminosity binaries down to small separation,
compared to late type objects in this and other regions. Several dedicated
surveys for multiplicity were also performed (Kouwenhoven et al. 2005; Janson
et al. 2013; Lafrenière et al. 2014). However, given that a fixed number of
stars from ScoCen were added to the SHINE initial sample (40 for each priority
bin, see Desidera et al. 2021), evaluating the probability that any of the
excluded binaries would have been observed is extremely difficult. We
therefore chose to not re-include those in our sample, and we will also
exclude all the ScoCen systems (squares in Fig. 11) from the following
analysis.

Finally, there is a number of young field objects included in the SHINE
sample, which cannot be associated with any of the moving groups mentioned
above. The information on the binarity in this case is rather incomplete and
coming from scattered sources, and evaluating the impact on the initial bias
would be quite complicated. Even if we could retrieve the full list of
excluded binaries, to estimate the expected priority we would need to simulate
the information available at the time of the selection, in particular
regarding the age of the system. Then a full re-assessment of all the systems
characteristics would be needed to properly consider them in our analysis. As
this would be well beyond the scope of this paper, and will be presented in
the final statistical analysis of the SHINE survey, we decided to also exclude
the field binaries from the following discussion.

Taking into account all the caveats discussed above, we then construct a
complete moving group sample, limiting the census to targets belonging to
young moving groups and re-introducing the excluded objects mentioned above.
This sample includes a total of 231.84 stars (189 stars observed within SHINE
and 42.84 that would have been observed but actually were not because known
binaries), of which 32.91 are binaries with a separation of
$0.05\arcsec<\rho<0.5\arcsec$. This leads to a first estimate of the binary
frequency in this separation range of $14.2\pm 2.9\%$ (the error being given
by Poisson statistics). This value appears to be slightly higher than what
reported in previous studies (Duquennoy & Mayor 1991; Raghavan et al. 2010)
that suggest a frequency of companions with period comparable to those in our
statistical sample close to $\sim 11\%$. This difference is significant only
at slightly more than 1-$\sigma$ and it may then be an artefact of low number
statistics. If confirmed by more data, it would resemble the case of the
Taurus star-forming region that shows a slightly higher multiplicity frequency
with respect to field objects (Kraus et al. 2011). This is attributed to the
young age and the fact that these are low-density environments.

Figure 12: Primary mass ($M_{A}$) vs mass ratio ($M_{B}/M_{A}$) for binaries
in our reduced sample. The light grey squares and triangles show the position
of the binaries in ScoCen and in the field excluded from the reduced sample,
respectively. The plus signs show the position of the additional systems from
Table LABEL:tab:known_bin. The dashed line shows the position of the Hydrogen
burning limit ($M_{B}=0.08M_{\odot}$). Figure 13: Primary mass ($M_{A}$) VS
secondary mass ($M_{B}$) for binaries with separation between 0.05 and 0.5
arcsec in our sample. The light grey squares and triangles show the position
of the binaries in ScoCen and in the field excluded from the reduced sample,
respectively. The plus signs show the position of the additional systems from
Table LABEL:tab:known_bin. The dashed line indicates mass equality

###  5.2 Mass ratio distribution

Using the reduced sample defined in Sec. 5.1, we performed a tentative
analysis of the mass ratio distribution of the binaries in the SHINE
statistical sample. In the following discussion we will consider two entries
for triple (hierarchical) systems, one with the mass ratio obtained by summing
the masses of the closer binary, and another with the mass ratio within the
closer binary, that is, our $q$ values are $q_{\rm sys}$ according to the
notation of Tokovinin (2014b). However, at variance from the definition used
there, we forced $q$ to be $\leq 1$ for triple systems: that is, in cases
where the mass ratio for triple systems obtained with the optically brighter
primary would be larger than one, our $q$ value is the reciprocal of it.

Figures 12 and 13 show the mass ratio $q=M_{B}/M_{A}$ and the secondary mass
$M_{B}$ as a function of primary mass $M_{A}$ for the objects with companion
in the reduced sample (once again with the additional binaries from Tab.
LABEL:tab:known_bin shown with a different symbol).

As shown in Fig. 12, there seems to be no clear trend with the primary mass,
except that equal mass binaries seem to be slightly more common among the
stars with $M_{A}<1\leavevmode\nobreak\ M_{\odot}$. While our data are not
robust enough to warrant any solid conclusion in this respect, we note that a
similar trend has been obtained independently by Moe & Di Stefano (2017) (see
their Figure 35) extending over a much wider mass range but with a smaller
statistics in this particular mass range; and by El-Badry et al. (2019) with a
much larger statistics but wider separations using Gaia data (similar result
but with smaller statistics was also previously obtained by Söderhjelm 2007
using Hipparcos data). This seems then a consolidated effect and might be
related to, e.g., differences in the migration efficiency within disks as a
function of stellar mass, so that in massive systems equal mass binaries end
up closer to the star than the region we are considering (see Pinsonneault &
Stanek 2006; Moe & Di Stefano 2017; Tokovinin & Moe 2020).

Figure 14: Mass ratio distribution obtained using the Kernel Density Estimate
(KDE) method (see e.g. Silverman 1986) and a Gaussian kernel with
$\sigma=0.1$. The blue solid line shows the result obtained for the systems in
the complete sample described in Sec. 5.1 including the additional objects
from Table LABEL:tab:known_bin, weighted according to their probability of
being observed (see Sec. 5.1 for details). The grey shaded area shows the
distribution obtained using the full sample of SHINE binaries from this work,
while the gray dashed line shows the distribution obtained using only the
SHINE young moving group systems included in the complete sample. Figure 15:
Mass ratio distribution obtained using the Kernel Density Estimate (KDE)
method (see e.g. Silverman 1986) and a Gaussian kernel with $\sigma=0.1$. The
blue shaded area corresponds to the solid blue line in Fig. 14, while the
coloured solid lines show the distributions for solar type (gold) and low-mass
(red) primaries.

Figure 14 shows the distribution of the values of $q$ obtained with the Kernel
Density Estimate (KDE) method (see e.g. Silverman 1986) and a Gaussian kernel
with $\sigma=0.1$. Given the wide range of primary masses in our sample it is
difficult to properly assess our completeness in a fixed mass ratio range, as
the upper and lower bounds can correspond to very different companion masses.
Fig. 15 shows the mass ratio distributions obtained considering sub-samples
selected according to the primary mass, separating between solar-mass stars
($0.75<M_{A}/M_{\odot}<1.5$) and low-mass stars ($M/M_{\odot}<0.75$). Given
that the exclusion of the ScoCen members (shown with light grey squares in
Fig. 12 and 13) effectively removed most of the systems with high mass
primaries from the reduced sample, we did not consider primaries with
$M_{A}/M_{\odot}>1.5$. While the resulting sub-samples are too small to draw
any conclusion on the single distributions (the low mass primaries bin only
includes 10 systems), this should at least clarify how each group contributes
to the different peaks in the full distribution shown in Fig. 14.

Our distribution show two distinct peaks, one including nearly equal mass
systems, and the second peaking at $q\sim 0.21$. While the overall shape of
the distribution apparently contradicts earlier results showing flat
distributions (see e.g. Moe & Di Stefano 2017; El-Badry et al. 2019), it is
difficult to say whether such difference can be explained by the difference in
terms of completeness between our sample and those used to obtain such
results.

####  5.2.1 Dynamical masses

The values of the masses used so far are derived from a comparison of the
position of the stars in the colour-magnitude diagrams with isochrones,
assuming their ages (evolutionary masses; see Sec. 2.4). However, for a few
objects we are also able to estimate the dynamical mass using the amplitude of
the observed radial velocity curve $K$777When using Orbit, $K$ is a free
parameter that fits the radial velocity curve. Hence, we could use it to
derive the masses of the secondaries in spite of the fact that we usually
assumed the total mass of the system as derived from evolutionary models. and
assuming it as the velocity of the primary. This is true for systems with a
large contrast between the components, where and the secondary is very faint
in the optical. For these systems, we derived the mass of the secondaries
$M_{B}$ under the assumption that the total system mass of the system
$M_{A}+M_{B}$ is the one obtained from evolutionary considerations. The
relation between $K$ and $M_{B}$ requires knowledge of the orbit inclination;
this is very poorly determined for systems seen close to face-on, so we did
not consider such cases. We have only three systems for which all the needed
requisites are satisfied (HIP 36985, HIP 95149, and HIP 97255). Results are
given in Table LABEL:tab:orbits. Fig. 16 shows the comparison between the
photometric mass $M_{B}^{phot}$ and the dynamical mass ($M_{B}^{dyn}$) for the
3 systems for which we performed the analysis (blue dots). We found that
dynamical and evolutionary masses agree within their errors for all of our
three objects. This is also true for both components of the additional systems
from Tab. LABEL:tab:known_bin for which an orbital solution was available in
Tokovinin & Briceño (2018) (shown in Fig. 16 as grey and blue crosses. This
also further supports the goodness of the orbit derivations discussed in Sec.
4.4.1 and 4.4.3.

Figure 16: Dynamical mass ($M^{dyn}$) VS photometric mass ($M^{phot}$) for the
companions from our sample for which we were able to obtain an estimate of
$M^{dyn}$ (blue dots, see Tab. LABEL:tab:orbits), as well as both components
of the additional systems from Tab. LABEL:tab:known_bin for which an orbital
solution was available in Tokovinin & Briceño (2018).

###  5.3 Orbital parameters

Figure 17: Eccentricity VS orbital period for the systems in our sample with
orbit determination (blue open circles). The blue filled circle is the average
value ($e=0.416\pm 0.043$). The values for the additional objects from Tab.
LABEL:tab:known_bin are shown with plus signs. The grey squares and triangles
are the average values for spectroscopic binaries obtained by Udry et al.
(1998) and Griffin (2012), respectively. The grey diamonds are estimates of
average values for long period visual binaries by Tokovinin & Kiyaeva (2016).

In Sec. 4.4 we presented the orbital analysis for 25 of the systems presented
in this paper. The same kind of analysis was not possible for systems with
very long periods, for which the orbital coverage was too small even when
multiple epochs were available, nor for those with very short periods because
of the lack of observations (only the SHINE epochs were available). This
obviously introduces several biases in the derived period and eccentricity
distributions, which are difficult to assess and correct for, also due to the
difference in accuracy of our orbital determinations discussed in Sec. 4.4.
Therefore, while we could identify some interesting trends emerging in our
sample, the uncertainties and biases that affect our orbit determinations do
not allow for an accurate analysis of the orbital parameters distributions of
our targets, making our results only tentative.

Fig. 17 shows the eccentricity vs orbital period for the 25 SHINE targets with
orbital solutions (blue open circles) and the additional targets from Tab.
LABEL:tab:known_bin (plus signs), as well as the mean values from several
previous works (grey symbols). The mean eccentricity of our targets
($e=0.416\pm 0.043$, with an r.m.s. of 0.21; marked as a blue filled circle in
Fig. 17) appears to be lower than the typical value of $e=0.498\pm 0.044$ for
spectroscopic binaries in the Hyades studied by Griffin (2012) (blue square in
Fig. 17) and usually considered as the reference in this range of periods (see
e.g. Tokovinin & Kiyaeva 2016). Our value seems also very close to the average
value obtained for shorter period field spectroscopic binaries by Duquennoy &
Mayor (1991) and Udry et al. (1998).

The orbital elements of visual binaries have a preference to lower
eccentricities owing to observational limitations (Finsen 1936). For the same
reason, we also expect a trend for the time of passage at periastron ($T0$) to
be within the range of the observed epochs; this is not obvious in our data.
However, given that our errors on the eccentricity estimates vary quite
strongly due to the partial (and in some case very poor) orbital coverage, the
amplitude of this effect on our sample is not trivial to assess.

Keeping in mind the caveats discussed above, it is also worth pointing out
that our data seem to show a lower mean eccentricity for higher order systems
(about half of those discussed here) compared to the simple binaries.
Confirming this result (moderate eccentricities that are fairly independent of
semi-major axis up to a few tens of au) with a larger sample would be
important to clarify the origin of these systems as we expect that disk
fragmentation would favour lower eccentricity orbits with respect to core
fragmentation (see discussion in Tokovinin & Kiyaeva 2016) that might be
responsible for the formation of wider binaries. More observations and
analysis is required to clarify this issue.

Figure 18: Comparison between the values of the stellar rotational velocity
expected under the assumption that the spin axis of the primary is aligned
with the orbital axis of the binary ($V\leavevmode\nobreak\ \sin{i}_{\rm
exp}$) and the values measured from the spectra ($V\leavevmode\nobreak\
\sin{i}_{\rm obs}$). The error bars are estimated as described in Sec. 5.4.
All the values are reported in Tab. LABEL:tab:vsini. The black dashed line
denotes equality.

###  5.4 Spin-orbit alignment

If confirmed, the behaviour of the mass ratios and eccentricities discussed in
the previous sections would suggest that a significant fraction of the systems
with projected separation below a few tens of au ($<0.5$ arcsec in our sample)
actually formed through disk fragmentation, rather than core fragmentation
followed by capture. An additional expectation of such a scenario is that the
spin axis of the star should be roughly parallel to the direction of the
orbital angular momentum, though as shown by e.g. Bate (2018), late accretion
can considerably affect this prediction possibly leading sometimes to
misaligned disks. On the other hand, the situation is less clear in case of
binaries forming by core fragmentation along filaments, depending on the
relative role of magnetic fields (leading to alignment: Galli & Shu 1993;
Planck Collaboration et al. 2016a, b and turbulent fragmentation (leading to
misalignment: Offner et al. 2016), with a quite confuse picture emerging from
observations with a predominance of alignments over random distribution (Kumar
et al. 2011; Kong et al. 2019; Ménard & Duchêne 2004; Davis et al. 2009; Lee
et al. 2016; Stephens et al. 2017; Ansdell et al. 2016, 2018). Spin-orbit
alignment in binaries has been studied over the last 50 years (Weis 1974; Hale
1994; Harding et al. 2013 and references therein). In general, it has been
found that the majority of systems with separation $<40$ au are reasonably
well aligned (Hale 1994), although this claim have been recently challenged by
Justesen & Albrecht (2020), whose results suggest that the reported trend that
binaries with separations below 30 au are preferentially aligned could be
spurious. Furthermore, the fraction of systems with aligned spin-orbit seems
to decrease with separation (Jensen et al. 2004; Monin et al. 2006).
Most of these past studies are based on statistical arguments or make use of
the Rossiter-McLaughlin effect for eclipsing binaries. Our sample is too small
for a statistical study and the binaries are too wide to be eclipsing. To
verify the spin-orbit alignment for the binary of our sample we would need
measures of the orbital inclination, of the projected rotational velocity of
the star $(V\leavevmode\nobreak\ \sin{i}_{\rm obs})$ from spectroscopy and of
the rotation period from photometry; note that any result given by this
comparison is still statistical, because having the same value of $\sin{i}$
does not necessarily imply alignment of the axes. The required quantities are
actually available only for 12 systems (see Table LABEL:tab:vsini). The
comparison we considered is that between the observed value of
$(V\leavevmode\nobreak\ \sin{i}_{\rm obs})$, and the one expected
$(V\leavevmode\nobreak\ \sin{i}_{\rm exp})$ given the radius of the star, the
rotational period measured from photometry, and the orbital inclination,
assuming that the latter represents the inclination of the stellar rotation
axis. To estimate the radius (of the primaries) we compared its absolute $J$
magnitude - corrected for the contribution of the secondary - to that of main
sequence stars of the same spectral type (a proxy for the temperature) in the
tables by Pecaut & Mamajek
888http://www.pas.rochester.edu/~emamajek/EEM_dwarf_UBVIJHK_colors_Teff.txt.
The result of this comparison is shown in Figure 18. While the sample is
clearly very limited, we obtain a reasonably good agreement between the
predicted and observed value for $V\leavevmode\nobreak\ \sin{i}$, supporting
the spin-orbit alignment for the systems considered in this analysis. The most
discrepant case is TYC 6872-1011-1, which is a very fast rotator with a
photometric period of only 0.503 d but has only a moderately high value of
$V\leavevmode\nobreak\ \sin{i}=33.8$ km/s (Torres et al. 2006); the
discrepancy would be largely resolved if the real rotational period is a
multiple of the photometric period, but there is no clear evidence of this in
the light curve. Obviously more data are required to confirm this result with
an extended sample.

##  6 Conclusions and final remarks

In this paper we present 78 multiple systems with separation $<5$ arcsec
observed within the context of the SHINE survey, including fifteen triple
systems and six quadruple, 56 of which are new discoveries. Given the
extremely heterogeneous nature of the data sets used, several different
methods were used to detect these companions in the high contrast imaging data
provided by SPHERE. In particular, newly developed dedicated routines were
used to deal with very close systems.

We combined the SHINE data set with literature and archival data, trying to
better characterise these systems. In particular, we used TESS data to extract
stellar rotation periods for a large part of our sample; and we carefully
derived ages of all the systems, finding that they are generally young,
following the selection criteria of the SHINE survey. Finally, masses for all
components were derived from a comparison with model isochrones.

Given the strong selection bias against binarity applied while selecting the
SHINE sample (excluding any object with stellar companions within SPHERE FoV
at the time of selection, see Desidera et al. 2021, for details), it was not
possible to draw any definitive conclusion regarding the impact of our results
on the overall frequency of binaries. In an attempt to compensate for the
original bias while estimating the completeness of our sample, we limited our
analysis to the young moving group members included in the original SHINE
list, which allowed us to retrieve and re-introduce the excluded targets, also
taking into account their probability of being observed, based on the SHINE
priority bins. This extended sample of SHINE young moving group members counts
231.84 targets including 76.84 multiple systems (34 of which are among the 189
stars in moving groups observed within SHINE and 42.84 were originally
excluded from the selection because already known to be binaries). Note that
the census of binaries in this sample of young moving stars is still
incomplete because it does not include systems with very small or very wide
separation.

Restricting the analysis to companions with $0.05\arcsec<\rho<0.5"$, where our
sample is more likely to be complete, given the quality of our observations,
we find a binary frequency of $14.2\pm 2.9\%$. This indicates a possible
excess of binaries in our sample with respect to that with similar periods in
Duquennoy & Mayor (1991) and Raghavan et al. (2010), that is about 11%; this
difference is however only marginally significant.

A full discussion of all properties of these systems was beyond the purposes
of this paper, however, we note some very interesting trends. A very tentative
discussion of the mass ratio distribution of the binaries with
$0.05\arcsec<\rho<0.5"$ highlighted two distinct peaks, one including nearly
equal mass systems, and the second peaking at $q\sim 0.21$. While this result
seems to disagree with previous claim of smooth mono-modal distributions
resulting from previous studies such as those by Moe & Di Stefano (2017) or
El-Badry et al. (2019), it is not clear whether such difference can be
explained as a result of selection effects and possible differences in the
samples.

We derived orbital parameters for 25 systems by combining our data with
literature ones. To this purpose, we not only considered high contrast imaging
and speckle interferometry data, but also separate detection of the components
on the Gaia DR2 archive, the derivation of the proper motion anomalies (PMA)
at the epochs of Hipparcos and Gaia by Kervella et al. (2019), and radial
velocities whenever available. While additional constraints might possibly be
derived in the future for additional targets (for instance with a more
systematic combination of PMA with high contrast imaging), there are some
interesting preliminary results. For instance, we derived generally moderate
eccentricities for most of our targets, but once again it is unclear if this
result is actually robust against selection effects. Exploiting the knowledge
of the rotation period of the stars, we could compare the observed values for
the projected rotational velocity of the primaries ($V\sin{i}$) with that
expected if the stellar and orbital spins are aligned for 12 systems, in
general finding a good agreement (but with one possible exception) and
confirming on a star-by-star comparison the results found using statistical
arguments by Hale (1994).

Overall, the properties of the binaries in our sample seem to favour a disk
fragmentation scenario (Bate et al. 2002) for the formation of most of the
systems we studied. However, we should once more stress that these conclusions
are highly speculative, given the biases affecting our sample.

An example of the possible constraints that could be obtained using systems
like those discovered in our survey is given by HIP 107948 (Elliott et al.
2015). This is the only triple system for which we were able to derive
constraints on the orbits of all components, thanks to the combination of our
data with a rich literature999As mentioned in Appendix B, an orbital solution
for both components of this system was already available in Tokovinin et al.
(2020), showing values in good agreement with the ones presented in Tab.
LABEL:tab:orbits. Very intriguingly, we found that the orbits of components Aa
and B (the latter around the barycenter of the Aab system) are both prograde,
roughly co-planar (within $16\pm 13$ degrees), have very similar eccentricity
($e\sim 0.4$) and similar longitude of the periastron. This argues for a
common origin within a disk and for important dynamical interactions; this is
similar to the majority (but not all) of cases of triples with both orbits
determined studied by Tokovinin (2021).

We conclude by stressing the importance of combining different techniques when
tackling complex problems such as the formation of binaries. Even restricting
ourselves to a limited range in separation, such as the one considered here,
the most exciting results from our study required the combination of high
contrast imaging, speckle interferometry, space astrometry and photometry, and
high precision radial velocities and high resolution spectra. While our
analysis is essentially serendipitous, because SHINE was focused on substellar
companions around single stars, we demonstrate how large samples are often
useful for purposes different from those originally considered.

Sample characteristics |  |  |  |  |  |  |  |  |
---|---|---|---|---|---|---|---|---|---
ID | RA (2000) | DEC(2000) | Parallax | Proper Motion |  $J$ mag |  $H$ mag |  $K$ mag | SpType
| (h m s) | (d m s ) | (mas) | RA (mas/yr) | DEC (mas/yr) |  |  |  |
(continued) |  |  |  |  |  |  |  |  |
ID | RA (2000) | DEC(2000) | Parallax | Proper Motion |  $J$ mag |  $H$ mag |  $K$ mag | SpType
| (h m s) | (d m s ) | (mas) | RA (mas/yr) | DEC (mas/yr) |  |  |  |
ID | RA (2000) | DEC (2000) | parallax | pmRA | pmDEC | JMag | Hmag | Kmag | Sp.Type
HIP 2729 | 00 34 51.2019 | -61 54 58.129 | 22.512 $\pm$ 0.021 | 88.69 $\pm$ 0.04 | -52.66 $\pm$ 0.04 | 7.34 | 6.53 | 6.72 | K4Ve
AF Hor | 02 41 47.3054 | -52 59 30.645 | 23.075 $\pm$ 0.044 | 93.52 $\pm$ 0.14 | -11.58 $\pm$ 0.17 | 8.48 | 7.64 | 7.85 | M2Ve
TYC 8491-0656-1 | 02 41 46.8356 | -52 59 52.395 | 22.937 $\pm$ 0.046 | 96.77 $\pm$ 0.06 | -14.16 $\pm$ 0.07 | 7.58 | 6.76 | 6.93 | K6Ve
TYC 8497-0995-1 | 02 42 33.0255 | -57 39 36.830 | 20.109 $\pm$ 0.010 | 84.95 $\pm$ 0.77 | -9.24 $\pm$ 0.82 | 8.56 | 7.78 | 7.97 | K5Ve
HIP 17157 | 03 40 29.3861 | -47 55 30.550 | 39.127 $\pm$ 0.063 | 91.54 $\pm$ 0.06 | 102.45 $\pm$ 0.11 | 7.13 | 6.33 | 6.52 | K7V
HIP 17797 | 03 48 35.8772 | -37 37 12.541 | 18.809 $\pm$ 0.222 | 74.44 $\pm$ 0.71 | -9.09 $\pm$ 0.87 | 3.9 | 4.626 | 4.824 | B9.5
HIP 18714 | 04 00 31.9831 | -41 44 54.293 | 19.265 $\pm$ 0.032 | 69.08 $\pm$ 0.07 | -7.02 $\pm$ 0.08 | 7.20 | 6.88 | 6.94 | G3V
HD 25284 | 04 00 03.8327 | -29 02 16.443 | 19.771 $\pm$ 0.205 | 84.141 $\pm$ 0.198 | -13.207 $\pm$ 0.214 | 7.20 | 6.88 | 6.94 | K6V
HIP 19183 | 04 06 41.5333 | +01 41 02.078 | 17.519 $\pm$ 0.026 | 36.57 $\pm$ 0.06 | -94.59 $\pm$ 0.04 | 6.89 | 6.58 | 6.70 | F6V
2MASS J05195513 | 05 19 55.1413 | -07 23 40.061 | 18.696 $\pm$ 0.143 | 58.949 $\pm$ 0.144 | -39.213 $\pm$ 0.108 | 6.89 | 6.58 | 6.70 | M1
HIP 25434 | 05 26 22.9679 | -43 22 36.279 | 10.830 $\pm$ 0.011 | 20.332 $\pm$ 0.012 | 10.950 $\pm$ 0.015 | 8.056 | 7.891 | 7.76 | G0
TYC 7059-1111-1 | 05 28 56.5176 | -33 28 15.537 | 15.811 $\pm$ 0.029 | 10.50 $\pm$ 0.70 | -32.30 $\pm$ 0.70 | 8.43 | 7.76 | 7.92 | K3Ve
HIP 26369 | 05 36 55.0960 | -47 57 47.952 | 40.758 $\pm$ 0.040 | 28.75 $\pm$ 0.04 | 3.44 $\pm$ 0.04 | 7.45 | 6.61 | 6.83 | K6Ve
HIP 28036 | 05 55 43.1622 | -38 06 16.258 | 18.951 $\pm$ 0.016 | 20.69 $\pm$ 0.04 | 9.96 $\pm$ 0.04 | 6.49 | 6.21 | 6.31 | F7V
TYC 7079-0068-1 | 06 08 33.8679 | -34 02 54.957 | 11.483 $\pm$ 0.028 | 2.62 $\pm$ 0.07 | -16.13 $\pm$ 0.07 | 8.72 | 8.19 | 8.32 | G9Ve
TYC 7080-0147-1 | 06 11 55.7190 | -35 29 12.949 | 12.213 $\pm$ 0.116 | -12.31 $\pm$ 0.12 | -15.23 $\pm$ 0.12 | 8.83 | 8.18 | 8.34 | K1Ve
TYC 7627-2190-1 | 06 41 18.4930 | -38 20 36.076 | 8.586 $\pm$ 0.074 | -6.12 $\pm$ 0.14 | -15.60 $\pm$ 0.18 | 9.48 | 8.81 | 8.94 | K2Ve
UCAC2 06727592 | 06 47 53.37 | -57 13 32.0 | 8.671 $\pm$ 0.022 | -6.84 $\pm$ 0.03 | 4.00 $\pm$ 0.02 | 9.67 | 8.93 | 9.12 | K4V
TYC 9493-0838-1 | 07 30 59.4435 | -84 19 27.687 | 14.537 $\pm$ 0.020 | -19.59 $\pm$ 0.04 | 57.88 $\pm$ 0.05 | 8.40 | 7.91 | 8.02 | G9V
HIP 36985 | 07 36 07.0798 | -03 06 38.791 | 70.275 $\pm$ 0.131 | 68.24 $\pm$ 0.36 | -289.99 $\pm$ 0.19 | 6.79 | 5.93 | 6.16 | M1.0V
TYC 8911-2430-1 | 07 43 42.8741 | -61 07 17.297 | 10.224 $\pm$ 0.028 | -4.28 $\pm$ 0.09 | 25.29 $\pm$ 0.06 | 9.23 | 8.52 | 8.66 | K0V(e)
HIP 37918 | 07 46 14.8430 | -59 48 50.634 | 29.132 $\pm$ 0.286 | -70.52 $\pm$ 0.19 | 150.57 $\pm$ 0.21 | 6.75 | 6.28 | 6.36 | K0IV-V(k)
TYC 7657-1711-1 | 07 48 49.8160 | -43 27 05.672 | 10.616 $\pm$ 0.084 | -21.39 $\pm$ 2.04 | 34.17 $\pm$ 1.23 | 8.92 | 8.19 | 8.31 | K4Ve
TYC 7133-2511-1 | 08 08 39.2692 | -36 05 01.597 | 2.016 $\pm$ 0.374 | -4.55 $\pm$ 0.77 | 7.81 $\pm$ 0.78 | 9.53 | 8.71 | 8.88 | K4IVe
TYC 6004-2114-1 | 08 09 38.6112 | -20 13 50.015 | 8.917 $\pm$ 0.043 | -52.29 $\pm$ 0.05 | 28.66 $\pm$ 0.04 | 7.38 | 7.15 | 7.22 | F0V
TYC 8577-1672-1 | 08 50 08.1037 | -57 45 59.243 | 10.267 $\pm$ 0.117 | -35.68 $\pm$ 0.34 | 35.38 $\pm$ 0.34 | 8.61 | 7.96 | 8.13 | K3Ve
TYC 8582-1705-1 | 08 58 48.7127 | -53 03 25.034 | 18.058 $\pm$ 0.060 | -64.20 $\pm$ 0.06 | 95.20 $\pm$ 0.05 | 8.70 | 8.02 | 8.18 | K4V
TYC 8944-1516-1 | 09 13 30.2735 | -62 59 09.357 | 8.423 $\pm$ 0.048 | -38.90 $\pm$ 0.10 | 28.10 $\pm$ 0.09 | 8.88 | 8.32 | 8.45 | K0V
GSC 08584-01898 | 09 31 25.4223 | -53 14 36.608 | 7.477 $\pm$ 0.019 | -20.20 $\pm$ 0.06 | 9.71 $\pm$ 0.06 | 10.19 | 9.50 | 9.66 |
TYC 4895-1137-1 | 09 46 15.0293 | -00 24 53.571 | 17.005 $\pm$ 0.527 | -63.73 $\pm$ 1.04 | -47.60 $\pm$ 0.91 | 8.40 | 7.65 | 7.83 | K5
TYC 8182-1315-1 | 10 07 25.1550 | -46 21 49.587 | 11.986 $\pm$ 0.855 | -49.86 $\pm$ 1.59 | -30.35 $\pm$ 1.54 | 8.77 | 8.07 | 8.19 | K4V
HIP 49767 | 10 09 31.7653 | -32 50 47.993 | 23.049 $\pm$ 0.028 | -189.31 $\pm$ 0.06 | 37.45 $\pm$ 0.08 | 6.40 | 6.12 | 6.22 | G0V
TYC 7191-0707-1 | 10 17 39.5087 | -34 51 53.496 | 8.052 $\pm$ 0.039 | 17.41 $\pm$ 0.21 | -19.01 $\pm$ 0.24 | 8.03 | 7.59 | 7.67 | G0V
HIP 54477 | 11 08 43.9910 | -28 04 50.347 | 18.314 $\pm$ 0.075 | -70.00 $\pm$ 0.26 | -22.57 $\pm$ 0.23 | 5.31 | 5.24 | 5.35 | A1V
HIP 55334 | 11 19 52.7435 | -70 37 06.482 | 10.763 $\pm$ 0.036 | -39.34 $\pm$ 0.06 | 12.69 $\pm$ 0.05 | 7.32 | 7.08 | 7.15 | F2V
HIP 56128 | 11 30 18.6475 | -34 30 01.355 | 29.801 $\pm$ 0.107 | -407.87 $\pm$ 0.14 | 31.48 $\pm$ 0.12 | 7.77 | 7.03 | 7.18 | K4.5Vk:
HIP 56963 | 11 40 38.5816 | -49 30 33.125 | 7.513 $\pm$ 0.023 | -32.83 $\pm$ 0.05 | -5.88 $\pm$ 0.05 | 7.54 | 7.46 | 7.54 | A3V
HIP 58859 | 12 04 16.7045 | -51 28 21.051 | 8.310 $\pm$ 0.035 | -33.22 $\pm$ 0.06 | -10.15 $\pm$ 0.05 | 6.54 | 6.55 | 6.59 | B9.5V
TWA 24 | 12 09 41.8588 | -58 54 45.059 | 9.483 $\pm$ 0.013 | -37.43 $\pm$ 0.04 | -10.80 $\pm$ 0.04 | 8.32 | 7.66 | 7.81 | K1Ve
HIP 59603 | 12 13 22.3420 | -56 53 35.647 | 8.795 $\pm$ 0.019 | -34.78 $\pm$ 0.05 | -9.85 $\pm$ 0.04 | 7.70 | 7.43 | 7.52 | F2V
HIP 60577 | 12 24 59.6432 | -50 11 30.679 | 8.284 $\pm$ 0.024 | -34.66 $\pm$ 0.07 | -12.11 $\pm$ 0.06 | 7.40 | 7.11 | 7.22 | A8-9V
HIP 61087 | 12 31 12.6419 | -61 54 31.491 | 9.335 $\pm$ 0.028 | -38.26 $\pm$ 0.03 | -10.76 $\pm$ 0.03 | 7.012 | 6.819 | 6.74 | F6V
HIP 61796 | 12 39 55.8513 | -66 30 40.091 | 8.806 $\pm$ 0.036 | -34.61 $\pm$ 0.06 | -5.71 $\pm$ 0.07 | 6.33 | 6.37 | 6.36 | B8V
HIP 62171 | 12 44 26.5928 | -54 20 48.133 | 8.366 $\pm$ 0.024 | -35.60 $\pm$ 0.05 | -14.42 $\pm$ 0.05 | 8.05 | 7.75 | 7.83 | F3V
HIP 62428 | 12 47 39.2067 | -58 17 51.253 | 8.962 $\pm$ 0.111 | -36.58 $\pm$ 0.07 | -15.00 $\pm$ 0.07 | 6.42 | 6.29 | 6.38 | F0III
HIP 63041 | 12 55 03.908 | -63 38 26.86 | 10.049 $\pm$ 0.028 | -39.46 $\pm$ 0.05 | -11.66 $\pm$ 0.06 | 7.34 | 7.14 | 7.21 | F0V
HIP 65219 | 13 21 57.1323 | -51 16 56.141 | 8.426 $\pm$ 0.239 | -36.51 $\pm$ 0.40 | -19.14 $\pm$ 0.44 | 6.65 | 6.55 | 6.57 | A3 4III-IV
HIP 66908 | 13 42 43.6680 | -43 11 08.521 | 10.216 $\pm$ 0.032 | -37.48 $\pm$ 0.10 | -27.24 $\pm$ 0.12 | 6.91 | 6.85 | 6.91 | A4V
HIP 67036 | 13 44 16.0029 | -51 00 44.828 | 7.442 $\pm$ 0.053 | -28.84 $\pm$ 0.13 | -18.88 $\pm$ 0.18 | 6.61 | 6.71 | 6.70 | ApSiCr
HIP 70350 | 14 23 37.8770 | -43 57 42.857 | 8.46 $\pm$ 0.92 | -29.41 $\pm$ 0.75 | -29.20 $\pm$ 0.72 | 7.01 | 6.68 | 6.72 | F7V
HIP 70697 | 14 27 33.5926 | -46 12 48.665 | 7.394 $\pm$ 0.039 | -22.02 $\pm$ 0.13 | -21.07 $\pm$ 0.14 | 7.20 | 7.20 | 7.25 | A0V
HIP 70833 | 14 29 07.15141 | -43 21 42.8118 | 3.986 $\pm$ 0.124 | -28.98 $\pm$ 0.94 | -20.16 $\pm$ 1.04 | 8.02 | 7.79 | 7.84 | F3V
HIP 71321 | 14 35 05.3500 | -43 33 16.379 | 8.195 $\pm$ 0.032 | -25.47 $\pm$ 0.10 | -22.22 $\pm$ 0.11 | 7.32 | 7.15 | 7.25 | A9V
HIP 73913 | 15 06 17.9369 | -35 24 22.352 | 6.546 $\pm$ 0.027 | -20.62 $\pm$ 0.11 | -19.60 $\pm$ 0.09 | 8.00 | 7.83 | 7.84 | A9-F0V
HIP 75367 | 15 24 04.2604 | -41 09 41.688 | 7.030 $\pm$ 0.020 | -22.85 $\pm$ 0.10 | -20.57 $\pm$ 0.08 | 9.12 | 8.82 | 8.89 | F9V
HIP 77388 | 15 47 51.1608 | -38 15 36.032 | 7.787 $\pm$ 0.026 | -20.45 $\pm$ 0.11 | -31.27 $\pm$ 0.06 | 7.55 | 7.40 | 7.49 | A6V
HIP 77813 | 15 53 20.8950 | -19 23 53.515 | 5.655 $\pm$ 0.342 | -15.85 $\pm$ 0.25 | -28.65 $\pm$ 0.19 | 7.78 | 7.30 | 7.42 | F8V
HIP 78581 | 16 02 44.1532 | -30 40 01.460 | 9.580 $\pm$ 0.043 | -24.21 $\pm$ 0.08 | -39.88 $\pm$ 0.04 | 7.92 | 7.56 | 7.64 | G1V
HIP 79124 | 16 09 02.6012 | -18 59 44.039 | 7.372 $\pm$ 0.030 | -7.76 $\pm$ 0.12 | -24.15 $\pm$ 0.09 | 7.19 | 7.01 | 7.03 | A0V
HIP 79156 | 16 09 20.8883 | -19 27 25.958 | 6.706 $\pm$ 0.030 | -7.79 $\pm$ 0.13 | -23.43 $\pm$ 0.09 | 7.59 | 7.48 | 7.58 | A0V
HIP 82688 | 16 54 08.1435 | -04 20 24.660 | 22.126 $\pm$ 0.023 | -38.23 $\pm$ 0.07 | -114.17 $\pm$ 0.05 | 6.70 | 6.36 | 6.48 | F8-G0V
TYC 7364-0911-1 | 17 02 27.7996 | -32 04 36.184 | 10.703 $\pm$ 0.025 | -20.57 $\pm$ 0.10 | -54.18 $\pm$ 0.06 | 8.82 | 8.34 | 8.44 | G8V
TYC 8332-2024-1 | 17 04 23.4699 | -48 15 09.408 | 8.924 $\pm$ 0.112 | -12.20 $\pm$ 4.18 | -32.00 $\pm$ 3.93 | 8.83 | 8.11 | 8.26 | K5Ve
TYC 6820-0223-1 | 17 15 03.6088 | -27 49 39.737 | 12.00 $\pm$ 0.285 | -4.53 $\pm$ 0.45 | -45.82 $\pm$ 0.30 | 8.17 | 7.38 | 7.54 | K5Ve
TYC 7379-279-1 | 17 28 55.6135 | -32 43 56.802 | 17.804 $\pm$ 0.014 | -24.42 $\pm$ 1.54 | -116.95 $\pm$ 1.26 | 8.34 | 7.67 | 7.82 | K3V
HIP 87386 | 17 51 25.7589 | -53 50 19.579 | 14.991 $\pm$ 0.147 | -22.23 $\pm$ 0.24 | -68.71 $\pm$ 0.22 | 6.09 | 5.91 | 5.96 | A9V
TYC 6872-1011-1 | 18 58 04.1531 | -29 53 04.661 | 13.452 $\pm$ 0.044 | 13.41 $\pm$ 0.12 | -48.62 $\pm$ 0.11 | 8.86 | 8.02 | 8.16 | M0Ve
HIP 93580 | 19 03 32.2470 | +01 49 07.437 | 17.780 $\pm$ 0.085 | 17.68 $\pm$ 0.24 | -65.29 $\pm$ 0.19 | 5.46 | 5.32 | 5.36 | A4IV/V
HIP 95149 | 19 21 29.7277 | -34 59 00.356 | 50.286 $\pm$ 0.396 | 68.54 $\pm$ 1.94 | -98.78 $\pm$ 1.24 | 5.32 | 4.93 | 5.05 | G2V
TYC 6299-2608-1 | 19 40 45.4980 | -16 17 01.737 | 12.861 $\pm$ 0.079 | 59.43 $\pm$ 0.18 | -55.22 $\pm$ 0.15 | 7.08 | 6.74 | 6.81 | F7V
HIP 97255 | 19 45 57.3402 | +04 14 54.416 | 34.332 $\pm$ 0.115 | 75.87 $\pm$ 0.08 | -11.56 $\pm$ 0.06 | 5.90 | 5.52 | 5.62 | G0V
TYC 5164-0567-1 | 20 04 49.3613 | -02 39 20.389 | 14.877 $\pm$ 0.015 | 25.06 $\pm$ 0.08 | -72.38 $\pm$ 0.04 | 8.48 | 7.92 | 8.03 | K1V
TYC 8400-0567-1 | 20 14 14.3342 | -50 06 49.597 | 20.253 $\pm$ 0.530 | 16.33 $\pm$ 0.41 | 2.57 $\pm$ 0.30 | 8.18 | 7.55 | 7.67 | K3V
HIP 107948 | 21 52 10.4179 | +05 37 35.859 | 32.79 $\pm$ 5.65 | 119.17 $\pm$ 7.95 | -150.29 $\pm$ 4.89 | 8.25 | 7.39 | 7.65 | M2.4V
HIP 109285 | 22 08 23.01477 | -32 59 18.378 | 27.089 $\pm$ 0.423 | 71.31 $\pm$ 0.67 | -33.00 $\pm$ 0.66 | 4.38 | 4.31 | 4.55 | A1.5IV
HIP 109427 | 22 10 11.9852 | +06 11 52.307 | 35.34 $\pm$ 0.85 | 282.18 $\pm$ 0.95 | 30.46 $\pm$ 0.67 | 3.38 | 3.33 | 3.38 | A1Va
HIP 113201 | 22 55 27.2132 | -52 18 09.747 | 42.493 $\pm$ 0.221 | 139.41 $\pm$ 0.12 | -88.61 $\pm$ 0.15 | 8.31 | 7.42 | 7.66 | M4V

Values of the adopted age (Age${}^{max}_{min}$) and rotation period (PRot), derived as described in Sec. 2.3, for all the stars in our sample. |  |  |  |  |  |
---|---|---|---|---|---|---
ID | Age${}^{max}_{min}$ | MG | Agerot | $P_{rot}^{TESS}$ | $P_{rot}^{lit}$ | Notes
| (Myr) |  | (Myr) | (days) | (days) |
*Table 2, continued. |  |  |  |  |  |
ID | Age${}^{max}_{min}$ | MG | Agerot | $P_{rot}^{TESS}$ | $P_{rot}^{lit}$ | Notes
| (Myr) |  | (Myr) | (days) | (days) |
*The ages of the young moving groups (MG) (Lower Centaurus Crux: LCC; Upper Centaurus-Lupus: UCL; Tucana-Horologium: TUC; AB Doradus: ABDO; Columba: COL; Argus: ARG; Carina-Near: CANE; Beta Pictoris: BPIC; Upper Scorpius: US) are adopted from Bonavita et al. (2016) and further discussed in Desidera et al. 2020, submitted. For Octans-Near (OCNE), the age is derived in this work from comparison of the age indicators with respect to those of the groups above and reference open clusters (Pleiades, Hyades). For CG30 group the age is taken from Yep & White (2020). The ages in this work are therefore homogeneous with those adopted in Vigan et al. (2017, 2020). |  |  |  |  |  |
The ages of the objects with no clear moving group membership indications have been estimated following the approach described in Desidera et al. (2020, submitted, see Sec.2.3 for details). |  |  |  |  |  |
For the objects with available TESS data we also include the derived value of (PRot) and the corresponding value of the gyrocronologic age (Agerot). The values of PRot retrieved from the literature are also listed with the corresponding reference (1Kiraga (2012); 2Oelkers et al. (2018), 3Messina et al. (2010), 4Messina et al. (2011), 5Messina et al. (2017), 6Wright et al. (2011), 7D20.). |  |  |  |  |  |
The note in the last column refers to the nature of the estimated period (F: fast rotator; P: retrieved period is likely to be from pulsations; B: retrieved period is likely the one of the listed companion; AB: period retrieved for both components). HIP 2729 | 45${}_{35}^{50}$ | TUC | – | 0.3767 - 0.3263 | 0.3771 | F, AB
AF Hor | 45${}_{35}^{50}$ | TUC | 28 | 3.39 | – | 0.0
TYC 8491-0656-1 | 45${}_{35}^{50}$ | TUC | – | 0.5196 - 0.5603 | 0.560, 05191,2 | F, AB
TYC 8497-0995-1 | 45${}_{35}^{50}$ | TUC | 135 | 7.408 | 7.38, 7.4121,3 | 0.0
HIP 17157 | 150${}_{100}^{200}$ | — | 170 | 7.694 | – | 0.0
HIP 17797 | 45${}_{35}^{50}$ | TUC | 63 | 6.897, 11.766 | – | B
HIP 18714 | 45${}_{35}^{100}$ | TUC ? | 100 | 3.509 | – | 0.0
HD 25284A | 45${}_{35}^{50}$ | TUC | 49 | 0.312 - 4.652 | 4.547 | F, B
HIP 19183 | 149${}_{100}^{180}$ | ABDO | 186 | 2.007 | – | 0.0
2MASS J05195513 | 50${}_{30}^{150}$ | — | 38 | 0.610 - 6.452 | – | F
HIP 25434 | 42${}_{35}^{50}$ | COL | 22 | 1.36 | – | 0.0
TYC 7059-1111-1 | 149${}_{100}^{180}$ | ABDO | Fast | 0.697 - 0.966 | 0.694, 0.9701,4 | F
HIP 26369 | 149${}_{100}^{180}$ | ABDO | 218 | 4.445 - 6.452 | – | 1010104.44 d could be the period of the wide companion UY Pic A
HIP 28036 | 42${}_{35}^{50}$ | COL | 39 | 0.948 | – | 0.0
TYC 7079-0068-1 | 149${}_{100}^{180}$ | ABDO | 220 | – | 3.383 | –
TYC 7080-0147-1 | 149${}_{100}^{180}$ | ABDO? | 32 | – | 1.7761 | –
TYC 7627-2190-1 | 50${}_{30}^{100}$ | — | – | 0.722 - 0.220 | 0.7264 | F, AB
UCAC2 06727592 | 120${}_{70}^{180}$ | — | 132, 234 | 6.230 - 7.860 | 6.053 | AB
TYC 9493-0838-1 | 149${}_{100}^{180}$ | ABDO | 97 | 5.002 - 7.145 | 4.94, 5.1011,3 | AB
HIP 36985 | 500${}_{300}^{700}$ | — | 992 | 11.11 | – | –
TYC 8911-2430-1 | 100${}_{50}^{150}$ | — | 69 | 2.317 | – | –
HIP 37918 | 250${}_{200}^{300}$ | CANE | 190 | 6.25 | – | –
TYC 7657-1711-1 | 50${}_{40}^{70}$ | ARG | – | – | 0.890, 0.4711,4 | –
TYC 7133-2511-1 | 1${}_{0}^{2}$ | CG30 | – | – | 3.741 | –
TYC 6004-2114-1 | 50${}_{40}^{800}$ | ARG ? | – | – | – | –
TYC 8577-1672-1 | 50${}_{40}^{70}$ | ARG | – | – | 52.531 | –
TYC 8582-1705-1 | 120${}_{40}^{300}$ | CANE? | 113 | 3.03 | 3.0081 | 0.0
TYC 8944-1516-1 | 50${}_{40}^{70}$ | ARG | – | – | 0.0 | –
GSC 08584-01898 | 35${}_{20}^{50}$ | — | – | – | – | –
TYC 4895-1137-1 | 150${}_{100}^{200}$ | — | 188 | – | 7.7451 | –
TYC 8182-1315-1 | 150${}_{100}^{200}$ | — | 115 | 3.126 | – | 0.0
HIP 49767 | 600${}_{400}^{800}$ | — | 621 | 5.002 | – | 0.0
TYC 7191-0707-1 | 400${}_{200}^{600}$ | — | – | – | – | –
HIP 54477 | 10${}_{7}^{13}$ | TWA | – | 0.0519 | – | P
HIP 55334 | 16${}_{12}^{40}$ | LCC? | – | 0.797 - 0.531 | – | F
HIP 56128 | 600${}_{300}^{2000}$ | — | 510 | 11.11 | 1.032 | 0.0
HIP 56963 | 16${}_{12}^{460}$ | LCC? | – | 0.0161 | – | P
HIP 58859 | 16${}_{12}^{20}$ | LCC | – | – | – | –
TWA 24 | 16${}_{12}^{20}$ | LCC | 9, 16 | 2.198 - 5.406 | 0.683 | AB
HIP 59603 | 16${}_{12}^{20}$ | LCC | 0.0 | 0.485 | – | F
HIP 60577 | 16${}_{12}^{20}$ | LCC | 9 (B) | 0.583 - 2.151 - 0.341 | – | B
HIP 61087 | 16${}_{12}^{20}$ | LCC | 9 | 1.539 - 1.631 | 0.0 | 0.0
HIP 61796 | 16${}_{12}^{20}$ | LCC | – | – | – | –
HIP 62171 | 16${}_{12}^{20}$ | LCC | 3 (B) | 2.326 | – | B
HIP 62428 | 16${}_{12}^{20}$ | LCC | 0.0 | 0.0438 | – | P
HIP 63041 | 16${}_{12}^{20}$ | LCC | – | – | – | 0.0
HIP 65219 | 17${}_{12}^{20}$ | LCC/UCL | – | – | – | 0.0
HIP 66908 | 17${}_{15}^{20}$ | UCL | – | 0.0183 | – | P
HIP 67036 | 16${}_{12}^{20}$ | LCC/UCL | – | 2.632 | – | 0.0
HIP 70350 | 17${}_{15}^{20}$ | UCL | 34 | 4.168 - 5.557 - 2.042 | – | –
HIP 70697 | 17${}_{15}^{20}$ | UCL | – | – | – | 0.0
HIP 70833 | 300${}_{200}^{1600}$ | — | 977 | 15.386 - 6.897 | – | 0.0
HIP 71321 | 17${}_{15}^{20}$ | UCL | – | 7.145 - 0.389 | – | 0.0
HIP 73913 | 17${}_{15}^{20}$ | UCL | – | 0.466 | – | F
HIP 75367 | 17${}_{15}^{20}$ | UCL | – | – | – | 0.0
HIP 77388 | 17${}_{15}^{20}$ | UCL | – | – | – | 0.0
HIP 77813 | 11${}_{4}^{12}$ | US | – | – | – | 0.0
HIP 78581 | 17${}_{15}^{20}$ | UCL | 7 | 1.137 | 1.1321 | 0.0
HIP 79124 | 11${}_{4}^{12}$ | US | – | – | – | 0.0
HIP 79156 | 11${}_{4}^{12}$ | US | – | – | – | 0.0
HIP 82688 | 149${}_{100}^{180}$ | ABDO | 74 | – | 2.776 | –
TYC 7364-0911-1 | 17${}_{15}^{50}$ | UCL? | 43 | 3.227 | 3.1921 | 0.0
TYC 8332-2024-1 | 17${}_{15}^{20}$ | UCL | 38 | 4.35 | 1.2091 | 0.0
TYC 6820-0223-1 | 24${}_{10}^{30}$ | BPIC? | 12 | – | 2.8571 | –
TYC 7379-279-1 | 149${}_{100}^{180}$ | ABDO? | 116, 415 | 5.884 - 4.001 | – | AB
HIP 87386 | 470${}_{100}^{840}$ | — | 424 | 0.0388 - 4.879 | – | P, B
TYC 6872-1011-1 | 24${}_{19}^{29}$ | BPIC | – | 0.505 - 0.685 | 0.505, 0.5031,5 | F
HIP 93580 | 570${}_{290}^{850}$ | — | – | – | – | 0.0
HIP 95149 | 130${}_{80}^{180}$ | OCNE? | 111 | 3.847 | – | 0.0
TYC 6299-2608-1 | 4000${}_{3000}^{5000}$ | — | 0.0 | – | – | –
HIP 97255 | 130${}_{80}^{180}$ | OCNE | 105 | – | 3.511 | –
TYC 5164-0567-1 | 149${}_{100}^{180}$ | ABDO? | 78 | – | 4.68, 4.691,3 | –
TYC 8400-0567-1 | 180${}_{100}^{300}$ | — | – | – | – | –
HIP 107948 | 149${}_{100}^{180}$ | ABDO | – | – | – | –
HIP 109285 | 240${}_{110}^{370}$ | — | 133 | 5.557 - 2.326 | – | 0.0
HIP 109427 | 560${}_{360}^{760}$ | — | – | – | – | 0.0
HIP 113201 | 750${}_{500}^{1000}$ | — | 394 | 16.67 | – | 0.0

|  |  |  |  |  |  |  |  |
---|---|---|---|---|---|---|---|---|---
ID | Obs Date | Mode | DITxNDIT (sec) | ND Filt | FoV Rot. | Seeing | $\tau_{0}$ | Strehl
| (JD - 245000) |  | $IFS_{Science}$ | $IFS_{PSF}$ |  | (deg) | (arcsec) | (ms) |
(continued) |  |  |  |  |  |  |  |  |
ID | Obs Date | Mode | DITxNDIT (sec) | ND Filter | FoV Rot. | Seeing | $\tau_{0}$ | Strehl
| (MJD) |  | $IFS_{Science}$ | $IFS_{PSF}$ |  | (deg) | (arcsec) | (ms) |
HIP 2729 | 57357.01 | IRDIFS | 96 x 48 | 32 x 5 | ND 2.0 | 28.8 | 0.50 | N/A |
HIP 2729 | 58378.19 | IRDIFS | 64 x 64 | 4 x 31 | ND 1.0 | 26.4 | 0.93 | 3.5 | 0.73
AF Hor | 57323.19 | IRDIFS | 64 x 47 | 16 x 8 | ND 1.0 | 25.3 | 0.95 | 1.6 |
TYC 8491-0656-1 | 57322.18 | IRDIFS | 64 x 46 | 16 x 8 | ND 2.0 | 33.4 | 1.62 | 1.0 |
TYC 8491-0656-1 | 58088.09 | IRDIFS | 96 x 32 | 4 x 31 | OPEN | 24.4 | 0.54 | 6.1 |
TYC 8497-0995-1 | 57356.08 | IRDIFS | 64 x 55 | 8 0 | ND 1.0 | 27.5 | 1.62 | 11.5 |
TYC 8497-0995-1 | 58087.07 | IRDIFS | 96 x 32 | 16 x 12 | ND 1.0 | 20.8 | 0.69 | 5.3 |
HIP 17157 | 57675.26 | IRDIFS | 64 x 53 | 2 x 47 | ND 1.0 | 34.3 | 0.42 | 2.8 |
HIP 17797 | 57709.28 | IRDIFS | 32 x 47 | 4 x 21 | ND 2.0 | 7.4 | 3.64 | 0.7 |
HIP 17797 | 58089.12 | IRDIFS | 32 x 160 | 16 x 10 | OPEN | 78.7 | 2.68 | 2.6 |
HIP 18714 | 58062.19 | IRDIFS | 96 x 48 | 2 x 47 | ND 1.0 | 48.7 | 0.58 | 4.8 |
HD 25284A | 57296.00 | IRDIFS | N/A | 16 x 4 | ND 2.0 | 0.0 | 0.61 | 1.3 | 0.72
HIP 19183 | 57407.02 | IRDIFS-EXT | 64 x 59 | 2 x 30 | ND 1.0 | 34.2 | 1.47 | 2.3 |
HIP 19183 | 57735.07 | IRDIFS | 64 x 53 | 2 x 47 | ND 1.0 | 23.7 | 1.07 | 4.8 |
2MASS J05195513 | 58063.27 | IRDIFS | 96 x 48 | 4 x 31 | OPEN | 54.2 | 0.54 | 6.7 |
HIP 25434 | 58145.03 | IRDIFS | 96 x 56 |  |  | 60.4 | 1.31 | 2.2 |
TYC 7059-1111-1 | 58090.22 | IRDIFS-EXT | 96 x 57 | 8 x 18 | ND 1.0 | 83.5 | 2.05 | 2.1 |
HIP 26369 | 57404.07 | IRDIFS | 64 x 16 | 4 x 15 | ND 1.0 | 6.9 | 0.98 | 3.3 |
HIP 28036 | 57405.10 | IRDIFS-EXT | 64 x 64 | 8 x 10 | ND 2.0 | 61.8 | 0.88 | 3.1 |
TYC 7079-0068-1 | 58174.00 | IRDIFS | 96 x 48 | 12 x 1 | OPEN | 91.7 | 0.82 | 7.4 |
TYC 7080 0147 1 | 58122.15 | IRDIFS | 96 x 48 | 16 x 16 | ND 1.0 | 76.1 | 0.59 | 11.3 |
TYC 7627-2190-1 | 58088.27 | IRDIFS | 96 x 42 | 4 x 31 | OPEN | 58.5 | 1.34 | 2.2 |
UCAC2 06727592 | 57404.14 | IRDIFS-EXT | 64 x 53 | 4 x 31 | OPEN | 27.6 | 1.62 | 2.0 |
TYC 9493 0838 1 | 57407.16 | IRDIFS | 64 x 2 | 2 x 35 | ND 1.0 | 0.2 | 1.12 | 2.4 |
HIP 36985 | 58468.23 | IRDIFS | 16 x 10 |  |  | 1.1 | 0.38 | 12.7 | 0.79
HIP 36985 | 58469.27 | IRDIFS | 16 x 50 |  |  | 0.1 | 0.39 | 7.6 | 0.81
HIP 36985 | 58552.02 | IRDIFS | 96 x 48 | 16 x 10 |  | 41.7 | 0.80 | 4.0 | 0.74
HIP 36985 | 58553.02 | IRDIFS | 96 x 30 |  |  | 29.5 | 0.77 |  |
TYC 8911-2430-1 | 58208.98 | IRDIFS | 96 x 2 | 2 x 47 | OPEN | 0.3 | 0.48 | 13.9 |
HIP 37918 | 58173.08 | IRDIFS | 96 x 48 | 16 x 10 | ND 2.0 | 30.1 | 0.55 | 10.1 |
HIP 37918 | 58549.26 | IRDIFS | 96 x 56 | 16 x 10 |  | 35.9 | 0.62 | 7.6 | 0.75
TYC 7657-1711-1 | 57060.61 | IRDIFS | 16 x 4 | 16 x 4 | ND 1.0 | 0.0 | 0.80 | 5.7 |
TYC 7133-2511-1 | 58135.00 | IRDIFS | N/A | 0.84 x 101 | OPEN | 0.0 | 0.75 |  |
TYC 6004-2114-1 | 58121.23 | IRDIFS | 96 x 48 | 4 x 31 | ND 1.0 | 123.4 | 0.66 | 6.4 |
TYC 8577-1672-1 | 57388.26 | IRDIFS | 4 x 224 | 8 x 29 | ND 1.0 | 27.7 | 0.90 | 3.0 |
TYC 8582-1705-1 | 58226.97 | IRDIFS | 96 x 32 | 12 x 13 |  | 24.3 | 0.53 | 12.4 | 0.82
TYC 8582-1705-1 | 58587.98 | IRDIFS | 96 x 36 | 12 x 13 |  | 28.9 | 0.73 | 5.0 | 0.79
TYC 8944 1516 1 | 58123.25 | IRDIFS | 96 x 48 | 16 x 10 | ND 1.0 | 27.9 | 0.39 | 7.9 |
TYC 8944 1516 1 | 58553.06 | IRDIFS | 96 x 48 | 16 x 10 |  | 28.3 | 0.84 | 2.6 | 0.69
GSC 08584-01898 | 57794.17 | IRDIFS | 64 x 64 | 8 x 19 | OPEN | 36.8 | 0.55 | 4.8 |
GSC 08584-01898 | 57875.96 | IRDIFS | 64 x 32 | 8 x 19 | OPEN | 25.4 | 1.35 | 2.0 |
GSC 08584-01898 | 57876.96 | IRDIFS | 64 x 73 | 8 x 19 | OPEN | 36.4 | 0.58 | 3.6 |
TYC 4895-1137-1 | 57792.68 | IRDIFS | N/A | 8 x 19 | ND 1.0 | 0.0 | 0.44 | 7.2 |
TYC 8182-1315-1 | 57060.19 | IRDIFS | 16 x 4 | 32 x 2 | ND 1.0 | 0.0 | 0.98 | 4.0 |
HIP 49767 | 58209.05 | IRDIFS | 96 x 48 | 12 x 13 | ND 2.0 | 1.5 | 0.62 | 12.2 |
TYC 7191-0707-1 | 58249.97 | IRDIFS-EXT | 8 x 18 |  |  | 2.7 | 1.38 | 2.9 |
HIP 54477 | 58227.06 | IRDIFS | 96 x 1 | 4 x 21 | ND 2.0 | 0.0 | 0.61 | 12.7 | 0.89
HIP 55334 | 57497.05 | IRDIFS | 64 x 43 | 4 x 31 | ND 1.0 | 15.1 | 0.73 | 3.2 |
HIP 55334 | 57831.83 | IRDIFS | 64 x 58 | 4 x 31 |  | 12.9 | 1.30 | N/A |
HIP 55334 | 58122.32 | IRDIFS | 12 x 90 | 4 x 31 | ND 1.0 | 27.3 | 0.40 | 17.2 |
HIP 56128 | 58225.06 | IRDIFS | 96 x 2 |  |  | 0.1 | 1.98 | 3.1 | 0.53
HIP 56963 | 58143.32 | IRDIFS-EXT | 32 x 96 | 4 x 31 | ND 1.0 | 36.4 | 0.50 | 8.4 |
HIP 58859 | 58209.17 | IRDIFS-EXT | 96 x 12 | 2 x 47 |  | 1.0 | 0.41 | 8.7 |
TWA 24 | 57177.97 | IRDIFS-EXT | 64 x 50 | 8 x 4 | ND 1.0 | 28.0 | 1.46 | 1.8 |
HIP 59603 | 58588.13 | IRDIFS-EXT | 96 x 48 | 6 x 23 |  | 32.3 | 1.26 | 3.1 | 0.76
HIP 60577 | 57883.10 | IRDIFS-EXT | 64 x 60 | 4 x 31 | ND 1.0 | 36.0 | 1.83 | 3.3 |
HIP 61087 | 58246.08 | IRDIFS-EXT | 48 x 96 |  |  | 29.5 | 1.03 | 1.6 | 0.51
HIP 61796 | 58588.11 | IRDIFS-EXT | 96 x 1 | 1.65 x 1 |  | 0.0 | 0.80 | 5.0 | 0.80
HIP 62171 | 58250.05 | IRDIFS-EXT | 32 x 144 | 8 x 18 |  | 36.7 | 1.00 | 2.8 | 0.71
HIP 62171 | 58548.23 | IRDIFS-EXT | 96 x 50 | 8 x 18 |  | 36.9 | 0.40 | 6.4 |
HIP 62428 | 57919.96 | IRDIFS-EXT | 96 x 64 | 2 x 47 | ND 1.0 | 41.6 | 0.56 | 5.6 |
HIP 63041 | 58208.05 | IRDIFS-EXT | 38 x 64 |  |  | 29.9 | 0.44 |  |
HIP 65219 | 57213.97 | IRDIFS-EXT | 32 x 50 |  |  | 13.0 | 0.96 | 1.7 | 0.61
HIP 65219 | 58175.28 | IRDIFS-EXT | 96 x 48 | 2 x 47 | ND 1.0 | 37.9 | 0.44 | 10.5 |
| 58634.06 | IRDIFS-EXT | 16 x 80 | 1.65 x 1 | ND 2.0 | 12.4 | 0.79 | 2.5 |
HIP 66908 | 57172.05 | IRDIFS | 64 x 4 | 8 x 15 | ND 2.0 | 2.0 | 0.79 | 1.8 |
HIP 67036 | 57154.13 | IRDIFS | 64 x 64 | 8 x 15 | ND 2.0 | 34.0 | 0.83 | 3.5 |
| 57220.97 | IRDIFS-EXT | 32 x 40 | 1.65 x 3 | ND 2.0 | 12.4 | 0.59 | 2.5 |
HIP 70350 | 58246.15 | IRDIFS-EXT | 32 x 1 | 2 x 47 |  | 0.0 | 0.90 | 1.6 |
HIP 70697 | 57153.14 | IRDIFS | 64 x 69 | 16 x 8 | ND 2.0 | 51.0 | 0.53 | 5.6 |
HIP 70833 | 58173.35 | IRDIFS-EXT | 96 x 39 |  |  | 41.7 | 0.65 | 8.4 |
HIP 70833 | 58287.05 | IRDIFS-EXT | 96 x 8 | 8 x 18 |  | 8.0 | 2.82 | 0.6 | 0.45
HIP 70833 | 58289.05 | IRDIFS-EXT | 16 x 68 | 8 x 18 |  | 53.7 | 1.40 | 3.1 | 0.80
HIP 71321 | 57509.60 | IRDIFS | N/A |  | ND 1.0 | 0.0 | 0.60 | 2.2 |
HIP 73913 | 58175.36 | IRDIFS-EXT | 96 x 48 | 8 x 18 | ND 1.0 | 78.2 | 0.46 | 11.1 |
HIP 73913 | 58553.35 | IRDIFS-EXT | 96 x 5 | 8 x 18 |  | 7.7 | 0.47 | 8.4 | 0.78
HIP 75367 | 58223.23 | IRDIFS-EXT | 96 x 1 | 2 x 47 | OPEN | 1.9 | 0.50 | 6.6 | 0.72
HIP 77388 | 58223.00 | IRDIFS-EXT | N/A | 4 x 31 |  | 0.0 | 0.50 | 8.1 |
HIP 77813 | 58256.18 | IRDIFS-EXT | 96 x 48 | 6 x 23 |  | 115.8 | 0.46 | 4.4 | 0.67
HIP 78581 | 57177.10 | IRDIFS | 64 x 63 | 8 x 4 | ND 1.0 | 124.3 | 1.04 | 2.9 |
HIP 78581 | 57177.12 | IRDIFS-EXT | 64 x 62 | 8 x 4 | OPEN | 30.7 | 1.03 | 3.0 |
HIP 78581 | 57832.00 | IRDIFS | 64 x 63 | 8 x 15 |  | 115.7 | 1.10 | 3.1 |
HIP 79124 | 58586.27 | IRDIFS-EXT | N/A | 4 x 31 |  | 1.0 | 1.20 | 2.7 | 0.68
HIP 79156 | 58586.29 | IRDIFS-EXT | 96 x 48 | 4 x 31 |  | 117.5 | 0.60 | 4.6 | 0.73
HIP 82688 | 58251.29 | IRDIFS | 64 x 48 | 2 x 47 | ND 1.0 | 32.8 | 0.71 | 4.4 | 0.78
TYC 7364-0911-1 | 57871.30 | IRDIFS | 64 x 49 | 16 x 11 | ND 1.0 | 64.4 | 0.45 | 5.0 |
TYC 8332-2024-1 | 57147.77 | IRDIFS | N/A | 8 x 15 | ND 1.0 | 0.0 | 0.90 | 1.5 |
TYC 6820-0223-1 | 57147.76 | IRDIFS | N/A | 4 x 4 | ND 1.0 | 0.0 | 0.65 | 1.6 | 0.77
TYC 7379-279-1 | 57173.20 | IRDIFS | 64 x 59 | 16 x 7 | ND 1.0 | 88.0 | 0.97 | 1.3 |
HIP 87386 | 57553.18 | IRDIFS | 64 x 8 | 8 x 19 | ND 2.0 | 4.3 | 1.49 | 1.8 |
HIP 87386 | 58247.27 | IRDIFS | 8 x 12 | 2 x 47 |  | 0.8 | 1.70 | 1.4 |
TYC 6872-1011-1 | 57288.99 | IRDIFS | 64 x 63 | 8 x 4 | ND 1.0 | 49.5 | 1.23 | 2.4 |
TYC 6872-1011-1 | 57648.97 | IRDIFS | 64 x 50 | 16 x 11 | ND 1.0 | 115.3 | 0.54 | 7.5 |
TYC 6872-1011-1 | 58256.31 | IRDIFS | 96 x 50 | 16 x 10 |  | 119.5 | 0.49 | 5.6 | 0.57
HIP 93580 | 57180.26 | IRDIFS | 16 x 60 | 4 x 8 | ND 2.0 | 9.1 | 1.29 | 2.9 |
HIP 93580 | 58290.21 | IRDIFS | 16 x 61 |  |  | 9.3 | 0.95 | 4.4 | 0.85
HIP 95149 | 58378.30 | IRDIFS | 64 x 48 | 4 x 31 | ND 2.0 | 55.6 | 0.75 | 2.7 | 0.82
HIP 95149 | 58621.42 | IRDIFS-EXT | N/A | 4 x 35 |  | 0.0 | 1.02 | 2.4 | 0.77
HIP 95149 | 58646.34 | IRDIFS-EXT | N/A | 4 x 70 |  | 0.0 | 1.07 | 1.9 | 0.73
TYC 6299-2608-1 | 57873.38 | IRDIFS | 64 x 8 | 2 x 47 | ND 1.0 | 5.7 | 0.64 | 4.5 |
HIP 97255 | 57208.21 | IRDIFS | 64 x 4 | 8 x 4 | ND 2.0 | 1.2 | 1.86 | 1.2 |
TYC 5164-0567-1 | 57157.36 | IRDIFS | 64 x 64 | 8 x 15 | ND 1.0 | 38.9 | 1.80 | 2.0 |
TYC 8400-0567-1 | 57873.41 | IRDIFS | 64 x 32 | 8 x 19 | ND 1.0 | 18.0 | 0.78 | 5.6 |
HIP 107948 | 57177.39 | IRDIFS | N/A | 8 x 8 | ND 1.0 | 0.0 | 1.09 | 3.5 |
HIP 109285 | 57204.33 | IRDIFS | 32 x 50 | 2 x 12 | ND 2.0 | 54.3 | 1.00 | 1.2 |
HIP 109285 | 57648.12 | IRDIFS | 16 x 229 | 2 x 15 | ND 2.0 | 87.8 | 0.52 | 19.7 | 0.94
HIP 109285 | 57675.03 | IRDIFS | 32 x 69 | 2 x 47 | ND 2.0 | 90.5 | 0.54 | 2.6 | 0.85
HIP 109427 | 57551.37 | IRDIFS | 8 x 48 | 8 x 19 | NDF3.5 | 3.2 | 1.14 | 2.4 |
HIP 113201 | 57649.13 | IRDIFS | 64 x 70 | 8 x 19 | ND 1.0 | 43.7 | 0.77 | 4.7 |
HIP 113201 | 57676.03 | IRDIFS-EXT | 64 x 51 | 8 x 19 | ND 1.0 | 39.3 | 0.84 | 2.9 |
HIP 113201 | 57709.01 | IRDIFS-EXT | 8 x 55 | 8 x 2 | ND 1.0 | 6.0 | 2.07 | 1.0 |
HIP 113201 | 57906.44 | IRDIFS | 64 x 2 | 8 x 19 |  | 0.8 | 0.71 | 5.0 |
HIP 113201 | 58621.39 | IRDIFS-EXT | N/A | 8 x 35 |  | 0.0 | 0.90 | 2.9 | 0.43

IFS contrast limits (expressed as $\Delta mag$) for all the available datasets. Red.Nc shows the reduction method used to obtain the limits from the non-coronographic images (see Sec. 3.4 for details): NcA: no-coro auto; sDIT: single DIT; NfR: no field rotation. |  |  |  |  |  |  |  |  |  |  |  |  |  |
---|---|---|---|---|---|---|---|---|---|---|---|---|---|---
ID | JD | Red.Nc | Contrast (nocoro) | Contrast (nocoro)
|  |  | 45mas | 75mas | 105mas | 135mas | 100mas | 200mas | 300mas | 400mas | 500mas | 600mas | 700mas | 800mas
*(continued) |  |  |  |  |  |  |  |  |  |  |  |  |  |
ID | JD | Red.Nc | Contrast (nocoro) | Contrast (nocoro)
|  |  | 45mas | 75mas | 105mas | 135mas | 100mas | 200mas | 300mas | 400mas | 500mas | 600mas | 700mas | 800mas
HIP 2729 | 57357.01 | NcA | 2.84 | 3.22 | 4.18 | 4.3 | 6.28 | 9.89 | 11.51 | 12.16 | 12.47 | 12.94 | 13.07 | 13.25
HIP 2729 | 58378.19 | NcA | 2.84 | 4.84 | 6.16 | 7.31 | 6.92 | 11.02 | 12.35 | 13.05 | 13.68 | 13.15 | 13.62 | 13.35
AF Hor | 57323.19 | sDIT | 0.6 | 1.68 | 3.55 | 4.41 | 6.94 | 9.27 | 10.67 | 10.88 | 11.26 | 11.31 | 11.54 | 11.53
TYC 8491 0656 1 | 57322.18 | sDIT | 1.95 | 2.2 | 3.59 | 4.02 | 8.07 | 11.44 | 13.13 | 13.97 | 14.36 | 14.82 | 15.45 | 15
TYC 8491 0656 1 | 58088.09 | NcA | 1.5 | 3.95 | 5.12 | 6.42 | 7.62 | 11.35 | 12.61 | 13.01 | 13.2 | 13.18 | 13.37 | 13.55
TYC 8497 0995 1 | 57356.08 | NcA | 3.44 | 3.1 | 5.22 | 5.54 | 6.24 | 8.85 | 9.82 | 10.99 | 11.12 | 11.66 | 11.73 | 11.63
TYC 8497 0995 1 | 58087.07 | NcA | 5.31 | 5.18 | 6.4 | 6.97 | 6.61 | 10.78 | 12.09 | 12.36 | 12.61 | 12.62 | 12.72 | 13.06
HIP 17157 | 57675.26 | NcA | 3.77 | 3.82 | 5.7 | 6.55 | 8.53 | 11.57 | 13.11 | 13.62 | 13.66 | 14.19 | 13.98 | 14.19
HIP 17157 | 57675.26 | NcA | 3.77 | 3.82 | 5.7 | 6.55 | 8.53 | 11.57 | 13.11 | 13.62 | 13.66 | 14.19 | 13.98 | 14.19
HIP 17797 | 58089.12 |  |  |  |  |  | 5.36 | 5.24 | 5.2 | 7.29 | 8.66 | 9.48 | 9.55 | 10.95
HIP 17797 | 57709.28 | sDIT | 2.7 | 3.24 | 3.05 | 3.21 | 9.76 | 12.55 | 14.6 | 14.16 | 15.01 | 14.97 | 15.37 | 15.49
HIP 18714 | 58062.19 | NcA | 4.3 | 5.79 | 6.84 | 7.74 | 9.07 | 12.48 | 13.56 | 13.72 | 14.18 | 14.12 | 14.7 | 14.61
HD 25284A | 57296 |  |  |  |  |  |  |  |  |  |  |  |  |
HIP 19183 | 57407.02 | NcA | 3.17 | 4.02 | 5.62 | 6.68 | 6.69 | 10.84 | 12.84 | 13.38 | 13.51 | 13.69 | 14.21 | 13.7
HIP 19183 | 57735.07 | NcA | 4.63 | 5.4 | 7.14 | 7.37 | 7.28 | 11.02 | 12.66 | 13.7 | 13.87 | 13.89 | 14.08 | 14.57
2MASS J05195513 | 58063.27 | NcA | 3.29 | 3.08 | 4.76 | 5.7 | 6.41 | 7.97 | 10.08 | 8.03 | 11.02 | 7.92 | 10.88 | 11.03
HIP 25434 | 58145.03 | NcA | 4.09 | 4.44 | 6.32 | 6.66 | 10.21 | 12.50 | 13.17 | 13.10 | 13.60 | 13.75 | 13.79 | 13.55
TYC 7059 1111 1 | 58090.22 | NcA | 2.69 | 4.1 | 5.68 | 5.75 | 8.36 | 10.43 | 10.48 | 11.26 | 11.27 | 10.86 | 11.58 | 9.51
HIP 26369 | 57404.07 | NcA | 4.73 | 4.87 | 6.13 | 6.75 | 6.71 | 8.8 | 11.48 | 11.55 | 12.61 | 12.02 | 12.1 | 12.23
HIP 28036 | 57405.1 | NcA | 3.97 | 4.19 | 4.97 | 5.08 | 9.24 | 12.69 | 14.27 | 14.14 | 14.2 | 14.35 | 14.35 | 14.39
TYC 7079 0068 1 | 58174 | NcA | 4.71 | 4.72 | 5.42 | 6.59 | 8.63 | 10.88 | 12.09 | 9.48 | 12.61 | 8.93 | 12.72 | 11.01
TYC 7079 0068 1 | 58174 | NcA | 4.71 | 4.72 | 5.42 | 6.59 | 8.63 | 10.88 | 12.09 | 9.48 | 12.61 | 8.93 | 12.72 | 11.01
TYC 7080 0147 1 | 58122.15 | NcA | 2.5 | 2.03 | 4.7 | 5.46 | 8.86 | 11.41 | 12.42 | 12.77 | 13.03 | 12.74 | 13.1 | 12.33
TYC 7627 2190 1 | 58088.27 | NcA | 1.72 | 2.85 | 4.59 | 6.25 | 8.43 | 10.05 | 11.23 | 11.34 | 11.66 | 11.43 | 11.84 | 11.15
UCAC2 06727592 | 57404.14 | NcA | 1.27 | 1.97 | 4.15 | 4.93 | 5.52 | 7.5 | 8.6 | 9.29 | 9.29 | 9.52 | 9.84 | 9.61
TYC 9493 0838 1 | 57407.16 | sDIT | 1.62 | 2.36 | 4.44 | 5.03 | 5.36 | 7.67 | 8.36 | 8.94 | 9.27 | 9.72 | 10.03 | 9.58
HIP 36985 | 58468.23 | sDIT | 2.26 | 3.01 | 3.12 | 3.24 | 4.05 | 5.05 | 5.46 | 5.47 | 5.86 | 6.53 | 6.97 | 7.53
HIP 36985 | 58469.27 | NfR | 2.89 | 5.05 | 3.24 | 3.44 | 3.44 | 7.29 | 7.05 | 7.39 | 7.39 | 7.52 | 7.48 | 7.31
HIP 36985 | 58552.02 | NcA | 4.91 | 4.4 | 4.89 | 5.8 | 4.58 | 11.45 | 13.13 | 13.42 | 13.57 | 13.82 | 13.86 | 13.92
TYC 8911 2430 1 | 58208.98 | sDIT | 3.79 | 4.53 | 5.47 | 5.34 | 7.54 | 9.08 | 9.99 | 9.78 | 10.19 | 10.9 | 10.96 | 10.75
HIP 37918 | 58173.08 | NcA | 4.17 | 5.19 | 6.25 | 6.23 | 7.5 | 12.33 | 14.16 | 14.45 | 14.7 | 14.44 | 14.69 | 14.92
HIP 37918 | 58549.01 | NcA | 4.42 | 3.94 | 4.81 | 5.02 | 7.08 | 11.09 | 12.99 | 13.24 | 13.62 | 14.03 | 14.13 | 14.08
TYC 7657 1711 1 | 57060.11 | sDIT | 2.08 | 0.51 | 2.23 | 3.84 | 5.61 | 6.23 | 6.06 | 4.02 | 5.94 | 6.06 | 6.21 | 6.4
TYC 7133 2511 1 | 58135 |  |  |  |  |  |  |  |  |  |  |  |  |
TYC 6004 2114 1 | 58121.23 | NcA | 4.03 | 2.61 | 4.77 | 6.44 | 8.25 | 13.15 | 14.34 | 13.9 | 13.99 | 14.98 | 14.66 | 14.26
TYC 8577 1672 1 | 57388.26 | sDIT | 1.74 | 0.32 | 1.31 | 2.45 | 4.19 | 6.76 | 6.76 | 7.17 | 7.38 | 7.66 | 7.73 | 7.93
TYC 8582 1705 1 | 58226.97 | NcA | 5.56 | 5.1 | 6.67 | 6.77 | 6.18 | 10.5 | 11.94 | 12.26 | 12.22 | 12.42 | 12.86 | 12.77
TYC 8582 1705 1 | 58587.98 | NcA | 4.72 | 4.87 | 6.89 | 7.02 | 6 | 10.43 | 11.9 | 11.99 | 12.08 | 12.38 | 12.26 | 12.87
TYC 8944 1516 1 | 58123.25 | NcA | 2.61 | 3.79 | 4.81 | 5.92 | 6.4 | 9.95 | 10.87 | 11.12 | 11.47 | 11.95 | 11.96 | 11.96
TYC 8944 1516 1 | 58553.06 | NcA | 2.32 | 4.27 | 5.37 | 6.32 | 5.16 | 10.26 | 11.02 | 11.15 | 11.58 | 11.6 | 11.86 | 12.13
GSC 08584-01898 | 57794.17 | NcA | 3.73 | 5.02 | 6.05 | 6.5 | 5.62 | 8.51 | 9.97 | 9.87 | 10.19 | 10.44 | 10.73 | 11.05
GSC 08584-01898 | 57875.96 | sDIT | 2.81 | 3.26 | 4.06 | 4.44 | 6.9 | 9.18 | 10.06 | 9.87 | 9.89 | 10.38 | 10.47 | 10.53
GSC 08584-01898 | 57876.96 | NcA | 3.53 | 3.72 | 5.06 | 5.33 | 6.69 | 8.94 | 9.95 | 10.11 | 9.97 | 10.47 | 10.1 | 10.29
TYC 4895 1137 1 | 57792.68 |  |  |  |  |  | 4.57 | 5.15 | 3.81 | 5.75 | 6.93 | 6.62 | 7.32 | 7.32
TYC 8182 1315 1 | 57060.19 | sDIT | 3.27 | 3.76 | 4.08 | 2.94 | 5.16 | 6.11 | 5.98 | 6.53 | 5.88 | 6.07 | 6.24 | 6.62
HIP 49767 | 58209.05 | sDIT | 3.6 | 4.01 | 5.74 | 5.82 | 10.66 | 11.68 | 13.32 | 12.46 | 13.58 | 13.48 | 14.15 | 13.75
TYC 7191 0707 1 | 58249.97 | NfR | 3.01 | 3.67 | 5.73 | 6.02 | 6.1 | 7.33 | 7.27 | 7.43 | 7.71 | 7.67 | 8 | 7.81
HIP 54477 | 58227.06 |  | 3.38 | 3.78 | 5.91 | 5.12 | 7.43 | 8.28 | 7.74 | 6.43 | 6.66 | 7.84 | 8.15 | 8.48
HIP 55334 | 57497.05 | NcA | 4.63 | 5.6 | 6.13 | 6.69 | 6.11 | 8.94 | 10.45 | 11.49 | 12.27 | 11.87 | 12.52 | 12.76
HIP 55334 | 57831.83 | NcA | 3.77 | 3.8 | 5.8 | 6.02 | 5.57 | 9.4 | 12.33 | 12.62 | 13.27 | 13.26 | 13.36 | 13.8
HIP 55334 | 58122.32 | NcA | 5.41 | 5.72 | 7.07 | 7.31 | 5.89 | 10.41 | 12.35 | 13.26 | 13.47 | 13.69 | 13.96 | 13.9
HIP 56128 | 58225.06 | NcA | 6.86 | 6.92 | 7.12 | 6.87 | 5.67 | 7.21 | 7.8 | 7.17 | 7.95 | 7.07 | 8.13 | 8.01
HIP 56963 | 58143.32 |  | 2.69 | 3.78 | 5.33 | 5.84 | 8.96 | 8.91 | 9.5 | 10.17 | 9.87 | 9.4 | 9.93 | 10.36
HIP 58859 | 58209.17 |  | 3.79 | 4.01 | 5.65 | 6.25 | 8.64 | 10.15 | 12.61 | 13.19 | 12.38 | 12.48 | 12.58 | 12.57
TWA 24 | 57177.97 | sDIT | 2.68 | 3.47 | 4.21 | 4.22 | 7.74 | 9.83 | 10.98 | 10.98 | 11.27 | 11.73 | 11.76 | 11.71
HIP 59603 | 58588.13 | NcA | 5.45 | 5.75 | 7.32 | 7.88 | 7.42 | 10.5 | 12.52 | 13.25 | 13.39 | 13.62 | 13.69 | 13.6
HIP 60577 | 57883.06 | NcA | 5.1 | 5.39 | 6.41 | 7.06 | 7.78 | 11.62 | 13.06 | 13.57 | 13.83 | 13.46 | 13.77 | 14.01
HIP 60577 | 57883.06 | NcA | 5.1 | 5.39 | 6.41 | 7.06 | 7.78 | 11.62 | 13.06 | 13.57 | 13.83 | 13.46 | 13.77 | 14.01
HIP 61087 | 58246.08 | NcA | 5.04 | 5.90 | 7.25 | 7.30 | 7.78 | 10.42 | 12.43 | 12.88 | 13.62 | 13.84 | 14.33 | 13.93
HIP 61796 | 58588.11 |  | 3.6 | 4.03 | 5.35 | 6.02 | 6.76 | 7.99 | 8.48 | 8.25 | 7.89 | 7.66 | 8.7 | 9.03
HIP 62171 | 58250.05 | NcA | 4.79 | 6.06 | 5.94 | 6.66 | 3.89 | 9.91 | 11.56 | 12.7 | 12.92 | 13.2 | 13.39 | 12.73
HIP 62171 | 58548.23 | NcA | 4.67 | 4.92 | 6.86 | 7.22 | 4.2 | 9.92 | 11.43 | 12.53 | 13.23 | 13.21 | 13.2 | 13.1
HIP 62428 | 57919.96 | NcA | 1.91 | 3.99 | 5.79 | 6.58 | 8.93 | 9.89 | 13.47 | 13.33 | 14.07 | 14.06 | 14.17 | 13.62
HIP 63041 | 58208.05 | NcA | 4.74 | 5.26 | 7.59 | 7.70 | 7.50 | 12.00 | 13.58 | 14.06 | 14.39 | 14.40 | 14.07 | 13.36
HIP 65219 | 57213.97 | NcA | 3.39 | 4.34 | 4.24 | 4.85 | 5.36 | 9.34 | 10.83 | 11.49 | 11.86 | 12.13 | 12.29 | 11.77
HIP 65219 | 58175.28 | NcA | 3.89 | 4.26 | 4.33 | 6.71 | 9.81 | 12.36 | 14.08 | 14.41 | 14.68 | 14.83 | 14.46 | 14.01
HIP 66908 | 57172.05 | sDIT | 2.97 | 3.43 | 5.29 | 4.77 | 7.08 | 10.43 | 11.37 | 11.53 | 12.34 | 12.2 | 12.1 | 11.58
HIP 67036 | 57154.13 | NcA | 3.21 | 3.08 | 5.11 | 5.1 | 5.32 | 9.73 | 12.18 | 12.87 | 12.94 | 14.05 | 13.93 | 14.13
HIP 70350 | 58246.15 |  |  |  |  |  |  |  |  |  |  |  |  |
HIP 70697 | 57153.14 | NcA | 3.07 | 4.17 | 4.57 | 4.8 | 7.03 | 10.55 | 11.94 | 12.99 | 12.75 | 13.12 | 13.49 | 13.69
HIP 70833 | 58173.36 | NcA | 5.08 | 5.87 | 7.84 | 7.58 | 8.51 | 11.74 | 13.59 | 13.95 | 14.21 | 14.49 | 14.26 | 13.96
HIP 70833 | 58287.05 | NcA | 4.13 | 5.38 | 5.98 | 6.31 | 5.81 | 7.27 | 9.68 | 9.93 | 10.37 | 10.49 | 10.8 | 11.02
HIP 70833 | 58289.05 | NcA | 4.75 | 5.39 | 7.28 | 7.84 | 7.67 | 11.35 | 12.91 | 12.34 | 12.93 | 12.7 | 13.36 | 13.3
HIP 71321 | 57509.6 |  |  |  |  |  | 4.04 | 6.02 | 5.98 | 6.31 | 7.04 | 7.32 | 7.32 | 7.31
HIP 73913 | 58175.36 | NcA | 5.9 | 3.79 | 3.83 | 6.39 | 7.47 | 13.36 | 14.08 | 14.05 | 14.53 | 14.36 | 14.2 | 13.32
HIP 73913 | 58553.35 | NcA | 5.6 | 5.88 | 6.81 | 7.74 | 5.41 | 9.55 | 12.6 | 12.65 | 12.49 | 13.01 | 12.85 | 12.9
HIP 75367 | 58223.23 |  |  |  |  |  | 3.14 | 5.56 | 6.43 | 6.15 | 6.18 | 6.94 | 7.06 | 7.14
HIP 77388 | 58223 |  |  |  |  |  | 3.01 | 5.39 | 6.23 | 6.13 | 6.23 | 6.96 | 7.11 | 7.17
HIP 77813 | 58256.18 | NcA | 3.21 | 5.18 | 6.03 | 7.2 | 9.31 | 11.17 | 12.91 | 13.15 | 13.7 | 14.27 | 13.81 | 13.69
HIP 78581 | 57177.1 | NcA | 3.54 | 4.31 | 5.22 | 5.7 | 9.45 | 12.2 | 12.9 | 13.37 | 13.9 | 13.63 | 13.83 | 13.94
HIP 78581 | 57177.18 | sDIT | 2.95 | 2.78 | 4.72 | 5.54 | 8.11 | 11.4 | 12.67 | 13.52 | 13.65 | 13.69 | 13.31 | 12.61
HIP 78581 | 57832 | NcA | 4.23 | 4.8 | 6.51 | 6.91 | 8.94 | 11.72 | 13.5 | 13.76 | 13.52 | 13.72 | 13.44 | 13.51
HIP 79124 | 58586.27 | NfR | 4.1 | 3.86 | 5.49 | 6.08 | 8.09 | 8.94 | 9.26 | 9.12 | 9.63 | 7.47 | 9.32 | 9.04
HIP 79124 | 58586.27 | NfR | 4.1 | 3.86 | 5.49 | 6.08 | 8.09 | 8.94 | 9.26 | 9.12 | 9.63 | 7.47 | 9.32 | 9.04
HIP 79156 | 58586.29 | NcA | 5.06 | 5.85 | 7.74 | 7.92 | 10.08 | 11.34 | 13.37 | 13.7 | 13.84 | 14.01 | 13.6 | 13.03
HIP 82688 | 58251.25 | NcA | 5.67 | 5.24 | 7.46 | 7.96 | 6.57 | 12.25 | 13.47 | 14.07 | 13.89 | 14.6 | 14.21 | 14.44
TYC 7364 0911 1 | 57871.3 | NcA | 4.36 | 4.06 | 3.64 | 6.18 | 8.3 | 10.02 | 12.47 | 12.63 | 13.06 | 13.39 | 13.38 | 12.9
TYC 8332 2024 1 | 57147.77 |  |  |  |  |  | 2.14 | 5.15 | 5.76 | 6.02 | 6.31 | 6.42 | 6.37 | 6.33
TYC 6820 0223 1 | 57147.76 |  |  |  |  |  | 3.17 | 5.95 | 6.06 | 6.25 | 6.59 | 6.6 | 6.65 | 6.81
TYC 7379 279 1 | 57173.2 | NcA | 3.6 | 4.24 | 4.33 | 4.09 | 9.52 | 12.29 | 12.67 | 13.6 | 13.67 | 13.63 | 13.72 | 13.46
TYC 7379 279 1 | 57173.2 | NcA | 3.6 | 4.24 | 4.33 | 4.09 | 9.52 | 12.29 | 12.67 | 13.6 | 13.67 | 13.63 | 13.72 | 13.46
HIP 87386 | 57553.18 | NcA | 6.3 | 6.77 | 6.72 | 6.88 | 5.84 | 9.36 | 9.95 | 11.33 | 11.66 | 12.01 | 12.18 | 12.34
HIP 87386 | 58247.27 | NcA | 3.74 | 4.73 | 5.31 | 5.48 | 6.99 | 7.34 | 7.45 | 7.47 | 7.94 | 7.46 | 8.01 | 7.2
TYC 6872 1011 1 | 57288.99 | NcA | 1.65 | 2.67 | 2.96 | 3.05 | 6 | 8.45 | 9.46 | 10.22 | 9.94 | 10.34 | 10.43 | 8.95
TYC 6872 1011 1 | 57648.97 | NcA | 1.63 | 3.52 | 4.67 | 5.95 | 9.08 | 11.24 | 11.82 | 12.03 | 12.44 | 12.46 | 12.56 | 11.26
TYC 6872 1011 1 | 58256.31 | NfR | 0.99 | 2.66 | 4.76 | 5.11 | 7.81 | 9.73 | 11.28 | 11.79 | 11.7 | 12.15 | 12.2 | 11.55
HIP 93580 | 57180.26 | sDIT | 2.49 | 3.08 | 5.37 | 5.77 | 5.38 | 7.8 | 8.47 | 11.74 | 12.94 | 13.01 | 13.4 | 13.96
HIP 93580 | 58290.21 | NcA | 4.46 | 5.16 | 6.61 | 6.86 | 6.38 | 11.37 | 10.08 | 13.89 | 13.95 | 14.05 | 14.36 | 14.49
HIP 95149 | 58377.98 | NcA | 5.53 | 6.2 | 6.99 | 6.54 | 7.1 | 10.92 | 12.78 | 13.5 | 14.1 | 14.08 | 14.19 | 14.49
HIP 95149 | 58621.42 | NfR | 3.71 | 5.36 | 6.87 | 7.61 | 1.12 | 4.91 | 5.96 | 7.13 | 7.22 | 7.67 | 7.26 | 7.9
HIP 95149 | 58646.34 | NfR | 4.36 | 5.37 | 6.59 | 7.12 | 2.91 | 7.88 | 8.1 | 8.52 | 8.37 | 8.68 | 8.17 | 8.08
TYC 6299 2608 1 | 57873.38 | sDIT | 3.34 | 4.07 | 5.27 | 3.75 | 6.94 | 5.7 | 12.26 | 11.72 | 12.31 | 12.45 | 12.33 | 11.19
HIP 97255 | 57208.21 | NcA | 6.62 | 7.02 | 7.13 | 6.82 | 6.2 | 9.18 | 10.09 | 10.88 | 10.98 | 10.45 | 11.5 | 11.92
TYC 5164 0567 1 | 57157.36 | NcA | 4.66 | 4.58 | 6.62 | 6.7 | 7.91 | 10.17 | 10.78 | 11.35 | 11.69 | 12.21 | 11.93 | 12.09
TYC 8400 0567 1 | 57873.41 | NcA | 4.15 | 4.63 | 5.74 | 6.83 | 5.66 | 10.32 | 11.86 | 12.07 | 12.63 | 12.79 | 12.76 | 12.82
HIP 107948 | 57177.39 |  |  |  |  |  | 2.41 | 5.34 | 5.79 | 5.9 | 6.19 | 6.29 | 6.23 | 6.49
HIP 107948 | 57177.39 |  |  |  |  |  | 2.41 | 5.34 | 5.79 | 5.9 | 6.19 | 6.29 | 6.23 | 6.49
HIP 109285 | 57204.33 | NcA | 3.2 | 3.25 | 4.84 | 5.38 | 8.35 | 12.24 | 13.32 | 14.09 | 14.5 | 14.53 | 14.88 | 15.21
HIP 109285 | 57648.12 | NcA | 1.86 | 2.21 | 3.31 | 4.68 | 8.8 | 10.68 | 12.47 | 13.28 | 13.06 | 13.33 | 12.92 | 13.05
HIP 109285 | 57675.03 | NcA | 5.47 | 5.86 | 6.78 | 7.13 | 10.74 | 13.88 | 15.03 | 15.02 | 15.51 | 15.82 | 16.32 | 15.12
HIP 109427 | 57551.36 | NfR | 2.65 | 3.44 | 4.33 | 4.14 | 6.44 | 9.51 | 10.21 | 11.46 | 11.58 | 11.06 | 11.36 | 11.19
HIP 113201 | 57649.13 | NcA | 4.56 | 4.72 | 5.27 | 6.46 | 7.57 | 10.1 | 11.97 | 12.34 | 12.81 | 13.01 | 12.85 | 13.26
HIP 113201 | 57676.03 | NcA | 3.17 | 3.1 | 4.65 | 4.91 | 6.18 | 9.05 | 10.39 | 11.25 | 11.95 | 11.93 | 12.05 | 11.67
HIP 113201 | 57709.02 | sDIT | 2.2 | 2.79 | 3.08 | 3.5 | 5.87 | 5.49 | 5.73 | 6.3 | 6.17 | 6.66 | 6.8 | 7.16
HIP 113201 | 57709.01 | sDIT | 2.51 | 2.81 | 3.26 | 3.76 | 5.83 | 5.73 | 6.42 | 6.98 | 7.22 | 7.15 | 6.64 | 7.18
HIP 113201 | 57906.44 | NfR | 2.87 | 3.01 | 4.64 | 5.05 | 5.42 | 7.03 | 9.65 | 10.12 | 10.81 | 11.29 | 11.11 | 11.17

##  Appendix B Notes on individual targets

HIP 2729 = HD 3221 The target was included in various surveys looking for
faint companions (SEEDS: Janson et al. 2013; NICI: Biller et al. 2013; VLT-
NACO LP: Vigan et al. 2017; GPIES: Nielsen et al. 2019; WEIRD: Baron et al.
2018) but it was not mentioned as binary. Our data indicates that this is a
close system with two components having nearly the same luminosity. Kinematics
indicates membership to Tuc-Hor group. Grandjean et al. (2020) found a large
range for radial velocities RV ($RV_{\rm amp}=4.79$ km/s) suggesting that this
is a spectroscopic binary. Since apparent separation of the binary is about 1
au, the RV variability may well be due to the newly detected companion. TESS
periodogram shows two very significant short period peaks consistent with both
the components being fast rotators.

AF Hor = GSC 08491-01194 It shares proper motion with TYC 8491 656 1
($\rho>20$ arcsec); both systems are member of the Tuc-Hor group. Binary
nature was discovered by Shan et al. (2017). The target was included in the
NICI (Biller et al. 2013), SEEDS (Janson et al. 2013), and WEIRD (Baron et al.
2018) surveys but no companion was detected. Our data indicates that this is a
close system with two components having nearly the same luminosity. TESS
reveals a period of 3.39 days, that if considered the rotation of one of the
two component agrees with expectation for a member of the Tuc-Hor group.

A number of constraints can be considered for the orbit of AF Hor, in addition
to the relative astrometric position measured by Shan et al. (2017) and by us.
These include the lack of radial velocity variations (Torres et al. 2006;
Kraus et al. 2014), the difference in the proper motion measured by Gaia DR2
with respect to historical values (SPM4 and UCAC4), from which we can extract
a PMA at epoch 2015.5 (that yields values of 1.0 and -7.8 mas/yr in RA and
Dec, respectively), and the lack of detection in the observations with NICMOS
at HST on 2004-10-24 and 2005-08-02 (proposal ID-10176, PI Song) that we
retrieved from the MAST
archive111111https://archive.stsci.edu/hst/search_retrieve.html, that implies
that separation was $<0.15$ arcsec at the epochs. The constraints from radial
velocity and PMA should take into account the dilution effect due to the
contribution by the secondary that is large; for instance for radial velocity
we estimate an upper limit of 5 km/s on the amplitude of radial velocity
variations for the primary. With these data, we used Orbit to explore the
range of possible solutions assuming the masses from photometry. We found
$P=10\pm 2$ yr (that implies $a=110\pm 20$ mas), $e=0.62^{+0.05}_{-0.14}$, and
$0<i<45$ degree.

TYC 8491-0656-1 Wide companion of AF Hor. Kinematics indicate membership to
Tuc-Hor association. No companion detected in SEEDS (Janson et al. 2013), NICI
(Biller et al. 2013), and Janson et al. (2017). It was previosly resolved by
Speckle (Tokovinin & Horch 2016; Tokovinin et al. 2018, 2019). Our data
indicates that this is a close, nearly equal mass binary. Then this is a
quadruple system. It shows significant $\Delta\mu$ between Tycho, Gaia DR1,
and Gaia DR2. The TESS periodogram shows two very significant short period
peaks consistent with both the components being fast rotators. Similar periods
were also reported by Kiraga (2012); Oelkers et al. (2018).

Combining our astrometry with literature data suggests a binary period of
about 2.3 yr and semi-major axis of about 45-50 mas, yielding a total mass
reasonably consistent with that determined from the photometric analysis of
the two components. When speckle data are taken at face value, residuals from
the best orbit are always very large; the best orbital solution has a very
large eccentricity and a quite large inclination. This might perhaps explain
why the binary is resolved only in some of the Speckle observations. However,
it is possible that the two stars are exchanged in some of the speckle data
set. This opens the possibility to an orbital solution with very small
residuals; this orbit would have low eccentricity (e=$0.27\pm 0.01$) and an
intermediate inclination ($i=127.0\pm 1.7$ degree). A more careful examination
of existing data and/or acquisition of new data might clarify this issue. We
note that interferometric observations might be needed for a full orbital
solution.

TYC 8497-0995-1 = GSC 08497-00995 Our data indicates that this is a close
system with a small luminosity difference between the two components.
Kinematics indicate membership to Tuc-Hor association. No companion was
detected by Janson et al. (2013), Biller et al. (2013), Shan et al. (2017),
Janson et al. (2017). The star has constant RV from SACY (Elliott et al.
2014). TESS reveals a period of 7.408 days, that if considered the rotation of
one of the two component agrees fairly well (it is actually a bit longer) with
expectation for a member of Tuc-Hor. A similar period was reported by Messina
et al. (2010); Kiraga (2012).

HIP 17157 = CD-48 1042 Binary system with both components in Gaia DR2, with
$\Delta$G=3.44. They have similar parallaxes and proper motion although with
formally significant differences. There is also a significant Gaia-Dr1 vs
Tycho2 $\Delta\mu$. Notably, our observation reveals a third object not much
fainter than the secondary. Since this further object is only detected in our
observation, it is not sure that it is a physical companion to the other two
objects; however, given the large distance of this star from the galactic
plane, this is highly probable. This third object is far enough from the two
other components so that the Gaia astrometry is unlikely to be affected
significantly; however, the impact of this source on the dynamics of the
system is unclear.

The star is not associated with known moving groups, but all the indicators
consistently indicate a young age. Periodogram of TESS data indicates a period
of 7.694 d that nicely overlaps the sequence of the Pleiades. Lithium and
activity are consistent with an age similar or slightly older then the
Pleiades.

Gaia DR2 gives two separate entries for the components of this system with
full astrometric solution. If we combine this data with our position, the
radial velocities from Trifonov et al. (2020) (that uses the same observations
described in Grandjean et al. 2020), and we assume the masses from photometry,
we derive that the period should be $>160$ yr (semi-major axis $a>1.09$
arcsec, that is $a>29$ au), the eccentricity is between 0.70 and 0.78, and the
orbit is seen with inclination $i>135$ degree, under the assumption that this
result is not affected significantly by the third component. If we add the
constraints on the PMA by Kervella et al. (2019), that are however not well
matched (perhaps because of the third component), the best result is at the
lower edge of the period range ($P=160$ yr, $e=0.78$ and $i=167$ degree, that
is an orbit seen almost face-on. Given the complex and unclear dynamics of
this system, we do not use it in the statistical discussions of Section 5.

HIP 17797 Triple system member of Tuc-Hor association. The secondary at about
8 arcsec is HD 24071, also an early type star. Both components are in Gaia
DR2, with similar parallax but with large offsets in proper motion (16.6
mas/yr in RA and 3.6 mas/yr in Dec) between the two components in Gaia DR2. We
found that the primary is a tight binary, possibly explaining the offset in
proper motion. Nothing found at large separation by WEIRD (Baron et al. 2018).
The TESS periodogram has a peak at 6.897 days that is consistent with the
expected rotation of the companion revealed by our data.

HIP 18714 = HD 25402A This star has a wide companion at 8.6 arcsec (Tokovinin
2014a) confirmed by Gaia DR2 parallax (HD 25402B). There are large offsets in
proper motion (2.6 mas/yr in RA and 6.6 mas/yr in Dec) and in RV (3.7 km/s)
between the two components in Gaia DR2. The primary was shown to be a tight
binary by our SPHERE observation; this might explain the offset in proper
motion with the wide companion. This is a new discovery from this study (the
star was previously observed by Baron et al. 2018, with nothing found). The
star was classified as a member of the Tuc-Hor association by Kiss et al.
(2011) and classified as bona fide member in several successive studies.
BANYAN kinematic analysis with Gaia parameters yield 22.0% membership
probability on Tuc-Hor for the primary and 99.9% for the secondary.
Independently on the kinematics, the age indicators are somewhat ambiguous:
the Li EW by Kiss et al. (2011) is compatible with Tuc-Hor but also with the
locus of older systems as the Pleiades and AB Dor MG, while clearly above the
Hyades. No rotation period was detected by Messina et al. (2011) and the star
has no X-ray emission from ROSAT. However, analysis of TESS light curve shows
a photometric modulation with possible period of 3.58 days, slightly slow for
a Tuc-Hor member. The position on the CMD of the secondary is well above the
ZAMS, indicating a young age of few tens Myr, unless it is itself a binary
with a companion contributing significantly to the integrated flux. Nordström
et al. (2004) noted significant RV variability (rms 2.8 km/s from 2 epochs
over 829 days). Considering the projected separation of 44 mas = 2.3 au, the
newly detected companion can be the responsible for the observed variability.

HD 25284A This multiple system is thought to be a member of Tuc-Hor
association. There is a wide companion (HD25284B) at 11 arcsec. HD 25284B was
flagged as 0.5” binary in WDS from Tycho, but this companion was not retrieved
in our SPHERE observations (Langlois et al. 2021) and then the target was kept
in the SHINE statistical analysis as an isolated object. Nothing found at
large separation by WEIRD (Baron et al. 2018). The companion we found around
the primary of this system is a new discovery. The TESS photometry reveals
both a short (0.312 d) and a long (4.652 d) period. Considering the
$V\leavevmode\nobreak\ \sin{i}$ of the targets (69.8 and 6.9 km/s for A and B,
respectively, Kraus et al. 2014), it is plausible that the two periods are
both real and belong to the two components.

HIP 19183 = HD 25953 This is an F6 member of AB Doradus. The secondary is in
Gaia DR2 without parallax and proper motion, likely due to the large $\Delta
G=7.75$. SPHERE second epoch confirms that the object is bound. The estimated
mass for the companion is $0.14$ $M_{\odot}$. The target is in GPIES (Nielsen
et al. 2019) and WEIRD (Baron et al. 2018), with no special comment. The
period from the TESS periodogram (2.007 d) agrees with expectation for a
member of AB Dor.

2MASS J05195513-0723399 = UCAC4 414-008464 Our detection confirms that the
companion detected with Astralux (Janson et al. 2012) is physically bound to
the primary. There is another star at 15 arcsec (2MASS J05195412-0723359),
which is itself a close visual binary, which has similar proper motion, albeit
with significant differences, likely due to orbital motion. This is then a
quadruple system of 4 M-type objects: the AB pair (observed with SHINE) with a
separation of 0.5 arcsec (28 au); and the CD pair (not observed with SHINE)
with a separation 0.8 arcsec (46 au). Spectral type M1 was reported for A
component and M4 for the C component. 2MASS J05195412-0723359 was flagged as
Columba member by Malo et al. (2013); Gagné et al. (2015), while 2MASS
J05195513-0723399 was flagged as Tuc-Hor member by Schlieder, priv. comm. With
Gaia DR2 parameters of the A component, there is not association with known
groups. The C component is in Gaia DR2 but without astrometric solution,
likely because of the effect its close companion. Astrometric solution for the
A component has large errors for the same reason.

We observed the primary with FEROS 121212Prog. ID 090.A-9010(A), PI J.
Schlieder on 22 and 25 Feb 2013. We classified the star as M1 and we measured
RV 27.2$\pm$1.0 km/s, and $V\leavevmode\nobreak\ \sin{i}$=4 km/s. The RV is
fully consistent with that measured by Schneider et al. (2019) ($27.36\pm
0.41$ km/s) and slightly discrepant with respect to Gaia DR2 (24.78$\pm$0.44).
Indications of high level of magnetic activity were detected in our spectra,
supporting the youth of the system (age $<300Myr$). We did not detect Lithium
in the spectrum of the primary, which implies an age older than about 30-40
Myr. The TESS periodogram gives a short (P=0.610 d) and a long (P=6.452 d)
period. If the long period is due to rotation of the most massive (and
brightest) object in the system, the age would be as short as 38 Myr (with an
uncertainty of 50%), that is in agreement with membership to Tuc-Hor. The
position on CMD of the primary after correction for the flux of the companion
is above main sequence, slightly below but compatible within errors with the
locus of Tuc-Hor members. We then adopt an age of 50 Myr, with minimum limit
at 30 Myr from Lithium and maximum at 150 Myr. Improved astrometric parameters
with forthcoming Gaia releases will allow improved placement on CMD and
kinematic analysis and then to refine the system age. For an age of 50 Myr and
the Gaia distance to the primary, the masses of the components result of 0.60,
0.15, 0.20, 0.13 $M_{\odot}$, respectively. Some constraints on the orbit are
obtained using the MC method (see Sect. 4.5.3 and Table LABEL:tab:orbits).

HIP 25434 Quadruple system, formed by two pairs (HIP 25434 and HIP 25436)
separated by 12.0”. The close companion to HIP 25434 is a new discovery from
the present paper, HIP 25436 was shown to be a SB by Desidera et al. (2015).
The star has a large PMA from Hipparcos and Gaia DR2 (Kervella et al. 2019).
The system is member of the Columba association (Moór et al. 2013; Desidera et
al. 2015), confirmed after Gaia DR2. The stellar rotation period of 1.36 d
obtained from TESS data agrees with membership in Columba association, it if
belongs to HIP 25434.

TYC 7059-1111-1 = UX Col The binarity of the object was originally discovered
in 1930 (WDS J05290-3328AB), but in WDS there are no report of observation
after 1942. Beside our recent SPHERE observations, both components are in Gaia
DR2 with similar parallax, some differences in proper motion likely due to
orbital motion, and small magnitude difference ($\Delta G=0.14$ mag). The
larger magnitude difference in the NIR suggests the possibility of an
additional component around the primary.

The star was flagged as AB Dor MG member by Torres et al. (2008). Our
kinematic analysis yields probabilities larger than 95% for the individual
values of each component, adopting Gaia DR2 results. The indirect age
indicators are also compatible with membership. TESS shows a large variability
(PTV $\sim$0.07 mag). The periodogram shows two strong peaks (at 0.6947 and
0.9712 d), plus a third one at 0.347 d, that is half the period of the highest
peak. Periods similar to these ones were reported by Messina et al. (2011)
(0.694 days) and Kiraga (2012) (0.970 days). It is likely that these two
periods are both real and correspond to the rotation periods of two components
in the system. arge differences in proper motion between available measurement
are reported. There is also a quite large difference between the Gaia and
SPHERE astrometry; the system was also observed optically in the thirties, at
a much smaller separation than observed by Gaia and SPHERE (0.3 arcsec with
respect to 1 arcsec at present epoch). Attempts to derive orbits adopting the
masses derived from the photometric analysis suggest an orbit seen almost
edge-on ($i\sim 89.95$ and $\Omega=268.45$) with a small or moderate
eccentricity ($e<0.4$). The semi-major axis and period cannot be determined
from existing observations alone; the only constrain is that semi-major axis
is $>1.7$ arcsec and the period is $>900$ yr.

An additional object with similar parallax and proper motion was identified by
Gaia DR2: Gaia DR2 4823262708795280384 = 2MASS J05284233-3326596 at 193”
$\sim$ 12000 au projected separation. It was flagged as a possible AB Dor
member by Gagné et al. (2018), with an expected M3 spectral type. This star is
also likely the counterpart of the X-ray source 1RXS J052843.5-332646 (nominal
separation 21”, well within ROSAT positional errors). The system is then at
least triple.

HIP 26369 = UY Pic B = CD-48 1893 Triple system formed by the isolated K0 star
HIP 26373 = UY Pic at 18” from the newly discovered close binary HIP 26369 =
UY Pic B. The system is a bona-fide member of the AB Dor MG. There is a 2.2
km/s RV difference between Gaia DR2 and SACY (Elliott et al. 2014). Nothing
was detected by NICI (Biller et al. 2013) and WEIRD (Baron et al. 2018). TESS
photometry yields a period of P=4.445 d. This is very close to the measured
period of HIP 26373 by Messina et al. (2010) and Kiraga (2012) (4.52 and 4.54
days, respectively). It is then plausible that it belongs to HIP 26373, which
is the brightest component in the system.

HIP 28036 = HD 40216 Member of Columba MG. The secondary is in Gaia DR2 but
without parallax and proper motion, and $\Delta$G=5.37. It was previously
detected by Galicher et al. (2016) and Song et al. (2003). Small RV difference
between SACY (Elliott et al. 2014: 24.38$\pm$0.09 km/s) and Gaia DR2
(25.02$\pm$0.45 km/s). Nothing detected at large separation in WEIRD (Baron et
al. 2018). TESS photometry yields a period of P=0.948 d, consistent with the
expected rotation of the primary if member of the Columba moving group.

TYC 7079-0068-1 = CD-34 2676 New triple system, with a first component at 100
mas with $\Delta H$=5.1 and a second one at 560 mas with $\Delta H$=4.2
(physical association not confirmed but very likely). The primary is a G9V
star, identified as member of AB Dor MG by Torres et al. (2008). The
membership probability results 52%, when adopting (Elliott et al. 2014) RV.
The TESS light curve is not available but rotation period is available from
Messina et al. (2010). Considering the fully consistent indirect age
indicators, we adopt AB Dor membership and age. The components are not
resolved in Gaia but there is a significant $\Delta\mu$ (3.5 $\sigma$ between
Gaia DR2 and Tycho2).

TYC 7080-0147-1 = CD-35 2749 Star identified as a member of AB Dor MG by
Elliott et al. (2014). The star was flagged as a possible SB in Torres et al.
(2006), from the the RV difference in two epochs. The close companion detected
with SPHERE (projected separation 61 mas = 4.9 au) is likely the responsible
for the RV variability and the significant difference of the astrometric
parameters between Gaia DR1 and DR2 and other astrometric catalogues. The
kinematic analysis with the Gaia DR2 astrometric parameters yields a very low
membership probability in the AB Dor MG (3.7%), while when adopting Gaia DR1
the probability jumps to 36.8%. Li EW is slightly larger than the mean locus
of Pleiades and AB Dor MG members but well within the observed distributions.
The rotation period by Kiraga (2012) indicates that the star is part of the
sequence of fast rotators. TESS data are not available. Pending a full orbital
solution and assessment of the impact of the binarity on the Gaia astrometric
solution, we consider the system as probable member of AB Dor MG and we adopt
the group age.

TYC 7627-2190-1  The star is classified as a member of AB Dor MG by Torres et
al. (2008). However, the kinematic analysis with Gaia DR2 parameters yields 0%
membership probability.The trigonometric parallax indicates a distance
significantly larger than the photometric one adopted in Torres et al. (2008),
which was underestimated because of the unrecognized binarity. There are no
indications of significant RV variability from sparse observation (Elliott et
al. 2014, , Gaia DR2, RAVE DR5). Our data indicates that this is a close
system with a small luminosity difference between the two components. TESS
periodogram shows two very significant short period peaks consistent with both
the components being fast rotators. One of the periods is very similar to the
one measured by Messina et al. (2011). The Li EW from Torres et al. (2006) is
larger than the observed values for Pleiades and AB Dor MG, suggesting a
younger age, and similar to members of IC 2391/IC 2602. We then adopt an age
of 50 Myr, with limits 30 to 100 Myr.

UCAC2 06727592 = GSC 8544-1037 Flagged as member of AB Dor MG in Torres et al.
(2006) while the membership was rejected in Elliott et al. (2014). Our
kinematic analysis also yield null membership probability. The binary was
originally discovered by Elliott et al. (2014). Both components are also in
Gaia DR2 with $\Delta$ G=0.31. TESS photometry yields two periods of P=6.23 d
and P=7.86 d, that are in agreement with the expected rotation periods of
primary and secondary, respectively, if the system has an age close to that of
AB Dor MG. The Li EW by Torres et al. (2006) is slightly above the mean loci
of the Pleiades and AB Dor MG but compatible within the observed
distributions. Independently on the kinematic membership, we then adopt an age
of 120 Myr.

Figure 21: Results from the SPOC analysis of TESS data for TYC 9493-0838-1.
Black dots are the pre-whitened TESS photometry phased at a period of 18.558
d; blue dots are running median values. The red line is a trapezoidal transit
model. The green points are residuals.

TYC 9493-0838-1 = CD-84 80 Member of AB Dor MG. The membership is confirmed by
BANYAN analysis with the updated kinematic parameters. The star has constant
RV according to SACY (Elliott et al. 2014) and Brems et al. (2019), with good
agreement between these radial velocity determinations.

The TESS data indicates a period of P=5.00 d, that agrees with expectation for
membership to the AB Dor moving group. Similar periods were measured by
Messina et al. (2010); Kiraga (2012). In addition, the SPOC analysis of the
TESS data detect a transit with SNR=12.2 and amplitude 2365 ppm (see Figure
21). In Gaia DR2 catalogue there is no other star bright enough ($G<16.24$) to
be responsible for this signal within 30 arcsec from the primary. The stellar
RV is stable within about a hundred m/s (Brems et al. 2019). The secondary has
dJ=2.7 mag, sep=585 mas, and PA=135 degree, so that we cannot exclude that the
transit occurs on the secondary. The in-transit/off transit offset is of about
$\sim$5 arcsec and PA=275 degree, that is roughly the expected accuracy for
such a small transit. If the transit occurred on the secondary (that should be
a star with an approximate spectral type of M2, and then a radius of about
0.434 R⊙ and a mass of 0.44 $M_{\odot}$) that should have a contrast of about
4 mag in visible light, it should have depth of about 94000 ppm. This requires
an object of about 0.14 R⊙. Given the estimated age of $149^{+31}_{-49}$ Myr,
this radius would correspond to a $68^{+10}_{-16}$ MJ BD or low-mass star
(Baraffe et al. 2015). On the other hand, if the transit is on the primary,
the radius would be of only 0.38 RJ, and then a planet with a mass of the
order of 5 $M_{\oplus}$ (Linder et al. 2019). This seems more probable than
the other hypothesis, and it is clearly still consistent with a lack of
detection from radial velocities.

HIP 36985 = BD-02 2198  Orbit and dynamical masses coupling imaging, HARPS RV,
and astrometry and comprehensive analysis of the stellar properties will be
presented in Biller et al., in preparation. The star is considered here only
for statistical purposes.

TYC 8911-2430-1 = CD-60 1850 = V838 Car Young object, classified as possible
member of Carina association by Elliott et al. (2016) but ruled out as a
member from our kinematic analysis. The lithium from Torres et al. (2006) is
slightly larger than the mean loci of Pleiades/AB Dor MG, and below Tuc-Hor
one. The rotation period from TESS is faster than the rotation sequence of
Tuc-Hor and Argus associations, suggesting a younger age. More likely, the
star is part of fast rotators sub-group. We adopt an age of 100 Myr with
limits 50/150 The primary is a solar type star (mass of about 0.94
$M_{\odot}$), while the secondary has a mass of 0.37 $M_{\odot}$.

Figure 22: Results from the SPOC analysis of TESS data for TYC 8911-2430-1.
Black dots are the pre-whitened TESS photometry phased at a period of 18.558
d; blue dots are running median values. The red line is a trapezoidal transit
model. The green points are residuals. The three panels are for the three
candidates, with periods of 218.3 (Planet 1, upper panel), 142.1 (planet 2,
middle panel) and 155.8 days (planet 3, lower panel).

Being close to south ecliptic pole, a very long TESS sequence is available for
this binary (contrast of $\Delta J=2.71$ mag, separation of 172 mas, that is a
projected separation of $\sim$17 au, PA=172 degree, and age of $\sim$50-60
Myr). Gaia DR2 detected only a star slightly fainter then the limit requested
to explain the candidate transits detected by TESS, at rather large separation
of 20.8 arcsec and PA=5.2 degree. The SPOC analysis revealed three possible
candidate planets (see Figure 22). In transit-off transit offsets measured by
SPOC analysis ($<4$ arcsec) are much smaller than this separation, are in a
completely different direction, and are within the errors. Given the small
separation of the binary, these offsets cannot be used to attribute the
transits to each of the two components. The SPOC analysis detected three good
candidates with periods of 218.3, 142.1 and 155.8 days, respectively, and an
additional one with very short period that however does not look of good
quality. It does not seem possible to attribute all these transits to the same
component. If they are due to faint companion around the primary, they may be
of planetary mass, while around the secondary they would be stellar. There
should be serious concern about dynamic stability of a system with two close
stellar companions. If these are all real transits, it seems reasonable to
attribute two of the candidates to the primary and the third to the secondary.
Still for dynamical stability, it seems then more likely that the primary has
two planets with periods of 218.3 and 142.1 days (close to the 3:2 resonance)
and the secondary has a stellar companion with period of 155.8 days. The
planets would have radii of about 1.3 and 0.7 Jupiter radii, and should then
be an external giant planet (mass of about 2 Jupiter mass) and a Neptunian
inner one (mass in the range 10-20 M⊕, using the models by Linder et al.
(2019). The companion of the secondary would have a radius of 0.27 R⊙, that
for the age of the system would imply a mass of $\sim 0.2$ $M_{\odot}$, using
the models by Baraffe et al. (2015). TYC 8911-2430-1 is then potentially very
interesting from a dynamical point of view because the binary looks quite
close - though of course this may be a projection effect. We cannot test this
because we have only a single SPHERE epoch.

HIP 37918 = HD 63581 The star is a member of Carina-Near MG (Zuckerman et al.
2006). It has an additional wide component at 23 arcsec, HIP 37923. The two
components are very similar each other (same spectral type, K0, $T_{\rm eff}$
difference of 60-100 K, Desidera et al. 2006b, $V$ magnitude difference 0.10
mag). The two stars have some differences in Li EW, $V\leavevmode\nobreak\
\sin{i}$ and possibly metallicity (Desidera et al. 2006b, a; Torres et al.
2006). Both components are metal enriched with respect to the Sun. The TESS
light curve yields a period of P=6.25 d, quite consistent with membership to
Carina-Near MG.

Our two SPHERE observations revealed a companion around the star, at a
projected separation of 60.5 mas (2.23 au) with a magnitude difference of
$\Delta J=3.8$ mag. The companion was also detected by Kammerer et al. (2019)
with NaCo. The presence of the companion can also be inferred by the high
$\Delta\mu$ signature and by some RV variability. In particular, the star is a
RV variable by Nordström et al. (2004).

While we have only three astrometric epochs (that is less than the seven
parameters describing the orbit), we can set interesting constraints on the
orbit of HIP 37918. In fact, the quite similar position found in Kammerer et
al. (2019) and in the second SHINE observation (but very different from the
intermediate first SHINE observation) suggests a period of a little more than
3 yr. If we couple this with the small total mass of the system ($\sim 1.24$
$M_{\odot}$), the semi-major axis cannot exceed by far $\sim 70$ mas and the
orbit should be fairly eccentric, with both these observations obtained rather
close to apoastron. This constrains the orientation of the semi-major axis and
the epoch of passage at periastron. On the other hand, the first SHINE
observation has a separation not far from the expected semi-major axis,
suggesting that also inclination is quite strongly constrained. This is
further constrained by the very small differences between the radial
velocities measured by Desidera et al. (2006b), GAIA DR2 and the catalogue
value (Gontcharov 2006) (data from individual epochs by Nordström et al. 2004
are not published), that implies that the orbit is seen almost face on
($i<3.6$ degree). With these consideration, we used the code Orbit by
Tokovinin (2016) to find the best astrometric orbit assuming that the
inclination at $i=2.4$ degree to fit a tentative orbit to the observed data.
This preliminary orbit has a period of $3.78\pm 0.09$ yr, a semi-major axis of
$a=78.4\pm 1.3$ mas, an epoch of T${}_{0}=2017.707\pm 0.018$, and an
eccentricity of $e=0.546\pm 0.015$. On the other hand, we cannot constrain
significantly $\Omega$ ($68\pm 136$ degree) and $\omega$ ($106\pm 136$
degree). With this orbit, the sum of the masses of the two components is
$M_{A}+M_{B}=1.38\pm 0.14$ $M_{\odot}$, in good agreement with the estimate
based on photometry ($M_{A}+M_{B}=1.24\pm 0.02$ $M_{\odot}$). In Section 4.5.2
we showed that an orbit with parameters very close to this one matches very
well the PMA measured by Kervella et al. (2019) from Hipparcos and Gaia
measurements. Further observations are needed to confirm these values.

TYC 7657-1711-1 = CD-43 3604 Member of Argus group (98.1% probability). It is
a very close binary. There are significant RV difference (7.6 km/s) between
Gaia and SACY (Elliott et al. 2014) RVs, but the star is indicated as constant
RV by SACY (Elliott et al. 2014). This is qualitatively in agreement with the
very small separation (40 mas = 3.4 au). There are two other sources within 5
arcsec in Gaia (delta mag 8-9 in G band). We retrieved them in our SHINE
images, though at low SNR. While the time elapsed between the SHINE
observation and Gaia epoch is short, both astrometry and position in the
colour magnitude diagram indicates that they are background objects. No data
available yet from TESS. Very short rotation period is derived by Messina et
al. (2011); Kiraga (2012).

TYC 7133-2511-1 = KWW 1637 Binary first identified by Janson et al. (2017).
The star is also flagged as SB2 by Kim et al. (2005), with an estimated RV
difference between the components of 40-45 km/s and by Torres et al. (2006).
The estimated magnitude difference in the $V$ band is of 0.5 mag that is
smaller than the observed $\Delta H$ (0.7 mag), suggesting a triple system.
The star was classified as member of Columba by Malo et al. (2014), with a
kinematic distance of 90 pc. Gaia DR2 yields a very small parallax
(1.42$\pm$0.38), which would imply a giant star. Neglecting it (that is,
assuming that Gaia DR2 astrometry is highly biased by binarity and that the
star is a true young object close to main sequence), BANYAN analysis yields
null membership probability for Columba and other known associations. This is
further supported by the extreme Li EW from Torres et al. (2006) (570 mÅ),
which would suggest a younger object. There are no data yet from TESS, but the
rotation period was measured by Kiraga (2012) on ASAS time series. Recently,
Yep & White (2020) proposed the star to be a member of a small association
near Cometary Globule CG 30 in the Gum Nebula, with an age of 0.5-1 Myr at a
distance of 358.1$\pm$2.2 pc. The differences of the observed RV and proper
motion of TYC 7133-2511-1 with respect to the group values are rather small
and compatible with the multiplicity of the object. The discrepancy of the
Gaia DR2 parallax with respect to the group distance becomes also smaller
although still formally significant. We consider this membership assignment as
fully compatible with all the observational constraints and we adopt it. The
star has a prominent IR excess longward of about 5.7 $\mu m$ from WISE,
Spitzer and AKARI.

TYC 6004-2114-1 = HD 67945 F0 star proposed as member of Argus association by
Torres et al. (2008) but rejected as a member by Zuckerman (2019) because of
the discrepant Gaia DR2 kinematics. Our kinematic analysis confirms the null
probability for being member when adopting Gaia DR2 values. However, a 99%
membership probability was derived when adopting Gaia DR1 values, which differ
by about 10 mas/yr in both coordinated of proper motion and 3.2 mas in
parallax (30%). The very close separation (72 mas = 6-8 au) of the newly
detected companion is the likely explanation of the discrepancy in the
astrometric parameters. The target was observed by Rameau et al. (2013) but
with no detection of the companion identified with SPHERE. The star was also
flagged as possible SB2 by Torres et al. (2006). A firm evaluation of the
membership and age will require astrometric solution including the presence of
the companion. Considering the difficulty in the dating of early F type star,
we tentatively adopt Argus membership, with upper limit as resulting from
isochrone fitting. No data available yet from TESS.

TYC 8577-1672-1 = CD-57 2315 Bona-fide member of the Argus association.
Constant RV from SACY (Elliott et al. 2014) and Gaia DR2. No data available
yet from TESS. The photometric period by Kiraga (2012) (52.53 d) is not
compatible with the available stellar properties.

TYC 8582-1705-1 = CD-52 2706 Young star, not previously classified as member
of known groups. The kinematic analysis yields a 65.4% membership probability
in Carina-Near MG. The TESS light curve yields a period of P=3.03 d, very
close to the previously reported by Kiraga (2012). The Li EW (Torres et al.
2006) and X-ray emission are close to the typical ones for Pleiades of similar
colors. The rotation period is instead faster, being intermediate between the
slow and fast rotators sequences in the Pleiades OC and similar to that of
Tuc-Hor members if belonging to the sequence of slow rotators. Therefore,
there is some tension between rotation (indicating an age of about 50 Myr of
the star is on slow sequence) and lithium, which is only marginally compatible
with such an age and indicative of an age close the Pleiades. The possibility
that the object is evolving from the fast rotators sequence might reconcile
these indicators. This latter possibility may also be compatible with Carina-
Near membership, although Li would be quite high in this case. Considering
these results we adopt an age of 120 Myr with limits 40 to 300 Myr (this
latter to allow the possible membership to Carina-Near). The companion found
with SPHERE is a new discovery. It is a very low mass star (0.18 $M_{\odot}$
for the nominal age) at 70 mas (4 au) projected separation. Some constraints
on the orbit are obtained using the MC method (see Sect. 4.5.3 and Table
LABEL:tab:orbits).

TYC 8944-1516-1 = CPD-62 1197 K star, identified as member of the Argus
association by Torres et al. (2008) and confirmed in our kinematic analysis.
No data are available from TESS but a fast rotation period (0.55 d) was
derived by Kiraga (2012) while Messina et al. (2011) found a period of 1.26 d.
The first one is compatible with the large vsini of the object (84 km/s Torres
et al. 2006). The rotation and the other age indicators are compatible with
Argus membership and age. The companions detected in this work is a new
discovery. Some constraints on the orbit are obtained using the MC method (see
Sect. 4.5.3 and Table LABEL:tab:orbits).

GSC 08584-01898 The close companion (separation 264 mas) was previously
detected by Shan et al. (2017). The star was flagged as member of Carina MG by
Moór et al. (2013). Gaia DR2 and Moór et al. (2013) show a marginally
significant RV difference ($16.53\pm 2.87$ and $22.6\pm 0.8$, respectively).
Independently on the adopted RV, BANYAN analysis classifies it as a field
object when adopting Gaia astrometric parameters. The Li EW from Moór et al.
(2013) (315 mÅ) clearly indicates a young age (between beta Pic MG to Tuc-
Hor/COL/Car MG; being consistent with both groups). The position on the CMD is
also above the main sequence and consistent with the sequence of Tuc-Hor
members. The young age is further supported by the X-ray emission detected by
ROSAT. No data available yet from TESS. There are two other sources within 5
arcsec in Gaia DR2, both background objects.

Since only a small fraction of the orbit has been covered since first
detection, we can only derive quite loose constraints on the orbit. Given the
current apparent separation, the semi-major axis is likely larger than 40 au
(0.31 arcsec); also it is probably smaller than 100 au (0.75 arcsec), unless
we consider the improbable circumstance of having observed the companion just
when being projected very close to the primary in a very wide orbit. If in
addition we assume the total mass of the binary as given by photometry, then a
family of best solutions is found with all quantities correlated with each
other. In particular, the period is in the range $270<P<1130$ yr, the
eccentricity in the range $0.47<e<0.64$, the inclination is $105<i<113$
degree, $104<\Omega<170$ degree, $-98<\omega<-53$, $1740<T_{0}<1800$, and
$21<V_{0}<25$ km/s. All these solutions provide a similar $\chi^{2}$ value
when using Orbit, so they are equally acceptable. We notice that the selected
period/semi-major axis range is only a range of most probable values, and we
cannot exclude longer periods.

TYC 4895-1137-1 Included in the target list as possible member of Tuc-Hor
(Schlieder, priv. comm.). However, the kinematic analysis with Gaia parameters
rejects the membership. Bright X-ray emission and a rotation period of 7.745
days have been reported by Kiraga (2012), while TESS data are not yet
available. We measured a Li EW of 92.8$\pm$3.4 mÅ on a FEROS spectrum
131313Prog. ID 090.A-9010(A), PI J. Schlieder. Both the lithium and rotation
period are close to the sequences of Pleiades members suggesting a similar or
slightly older age. The new companion is very bright ($\Delta J=0.49$ mag at
280 mas = 17.1 au projected separation).

TYC 8182-1315-1 Star identified as member of AB Dor by Torres et al. (2008).
However, BANYAN analysis returns a very small membership probability (1%) for
this group. The Li EW is compatible with Pleiades and AB Dor age (slighty
below the mean value). TESS photometry indicates a period of 3.126 d, somewhat
short but not incompatible with membership to AB Dor. When coupled with the
observed $V\leavevmode\nobreak\ \sin{i}$ (6 km/s, Torres et al. 2006) this
would indicate a fairly pole-on observations. Alternatively, the true rotation
period may be two times the observed one because of the configuration of the
active regions. X-ray emission and $\log{R_{HK}}$ (measured by us on FEROS
spectrum 141414Prog. ID 083.A-9003(A), PI R. Launhardt as in Desidera et al.
2015, $\log{R_{HK}}$=-4.44) are slightly below the mean locus of Pleiades and
AB Dor MG but well within the observed distributions. Considering these
ambiguities, we adopt an age of 150 Myr with limits 100 - 200 Myr.

HIP 49767 = HD 88201 Field star. The TESS photometry indicates a period of
P=5.00 d, close to the rotational sequences of the Hyades and Praesepe. The
Lithium EW we measured on FEROS spectrum 151515Prog. ID 083.A-9003(A), PI R.
Launhardt (77.0$\pm$1.7 mÅ) is just smaller than that of the Hyades members of
similar colors, while $\log{R_{HK}}$ and X-ray emission are slightly above
Hyades loci. The companion detected by SPHERE is then a very low mass star
(0.11 $M_{\odot}$). There are good perspective for dynamical mass
determination, as Hipparcos detected an astrometric acceleration and the star
is flagged as RV variable by Nordström et al. (2004) (see also the large PMA
between Hipparcos and Gaia DR2 measured by Kervella et al. 2019).

TYC 7191-0707-1 HD 89326 = CI Ant Visual companion previously known (Tycho,
WDS). The very short period (0.32d) photometric modulation seen in ASAS time
series was interpreted as a contact eclipsing binary by Bernhard (2011);
Kiraga (2012), indicating a triple system. There are no data yet from TESS.
Kinematic analysis indicates a field object. Lithium EW (Torres et al. 2006)
is intermediate between Hyades and Pleiades. none Existing data allows to put
significant constraints on the orbit of TYC 7191-0707-1. We can use the
positions measured at three epochs over about 90 years. The intermediate
measure (from the Tycho binaries catalogue) gives the widest separation. In
addition, the three measures have very similar PA with variations of about 10%
in separation. These facts essentially fix $\Omega$ and constrain $a$
(depending on $e$) and indicate that the orbit is seen at high inclination. If
we now assume the masses given by the photometry and $i=82$ degree, we can
find the best solution that has $a=1.4\pm 0.3$ arcsec (that is $a=175\pm 35$
au), $P=1600\pm 500$ yr, $T_{0}=1920\pm 380$ yr, $e=0.25\pm 0.10$,
$\Omega=123.1\pm 0.3$ degree, and $\omega=160\pm 20$ degree.

HIP 54477 = HD 96819 = 10 Crt = TWA 43 = HR 4334 This star was originally
proposed as a possible member of the TWA association by Mamajek (2005). It was
classified as a confirmed member and labelled as TWA 43 by Gagné et al.
(2017). The updated kinematic analysis including Gaia DR2 astrometry fully
supports this classification (membership probability 99.5%). It was observed
in several direct imaging surveys (De Rosa et al. 2014; Meshkat et al. 2015;
Nielsen et al. 2019) without mention of the close companion identified with
SPHERE. The RV monitoring by Lagrange et al. (2009) is not expected to have
been sensitive to the new companion due to the short time baseline and the
large RV errors due the high $V\leavevmode\nobreak\ \sin{i}$ of the star. The
star has a $\Delta\mu$ in right ascension significant at $>5\sigma$ between
Gaia DR2 and Hipparcos and $>4\sigma$ between Gaia DR2 and Tycho2 (see also
Kervella et al. 2019).

HIP 55334 = HD 98660 F2 star, probable member of LCC (membership probability
of 66.6%). The star is moderately close on the sky (2.78 deg) and with similar
kinematic parameter to the planet host HD 95086 (see discussion in Desidera et
al. 2021). It may be slightly older than the bulk of LCC, as also resulting
from isochrone fitting. Considering the somewhat ambiguous membership
probability, we adopt LCC age with an upper limit from the isochrone fitting
(pre-MS fit). The companion at 133 mas is a new discovery. A background object
at 3.3” was identified by Janson et al. (2013). Kervella et al. (2019)
measured a large PMA between Hipparcos and Gaia DR2. The motion and mass ratio
of the companion we found makes it fully compatible with being responsible for
the acceleration observed by Kervella et al. (2019).

We tried to constrain the orbit of this system using our data. Unfortunately,
only a limited portion of the orbit is available. The method we used, first
derivation of best possible orbit and then application of the PMA data on the
family of best orbits, is not adequate for this system, mainly because the
fraction of the orbit covered is too limited. The only solid conclusion that
we may draw with this approach is that the inclination at which we see the
orbit is large (in the range 75-86 degree) and that PMA data do not agree well
with highly eccentric orbit.

HIP 56128 = CD-33 7779 = TYC 7223-227-1 The companion detected with SPHERE at
211 mas (7 au) is likely the responsible for the observed astrometric features
(astrometric acceleration in Hipparcos, large Gaia/HIP/Tycho $\Delta\mu$,
large RUWE in Gaia). Three RV measurements (RAVE, Gaia, and our own using
FEROS 161616Prog. ID 083.A-9003(A), PI R. Launhardt) are consistent within
errors. The object is not associated with known moving groups. The TESS
photometry gives a period of 11.11 days, close to the rotational sequences for
the Hyades and Praesepe open clusters. The much shorter period (P=1.03 d)
reported by Oelkers et al. (2018) is likely an alias. The star has no lithium,
indicating an age older than 300 Myr, and was not detected by ROSAT. The
$\log{R_{HK}}$ (rederived from S index using B-V from Nascimbeni et al. 2016,
as the B-V from Hipparcos adopted in Gray et al. 2006 is highly discrepant
with respect to other photometric measurements and the spectral
classification) is -4.71, lower than the locus of the Hyades. We adopt an age
of 600 Myr with limits of 300 and 2000 Myr.

HIP 56963 = HD 101523 A3V star, classified as LCC member by de Zeeuw et al.
(1999) and Rizzuto et al. (2011). The kinematic analysis without including the
RV (as the RV in SIMBAD is the astrometric one from Madsen et al. 2002) yield
51.5% membership probability for LCC, 0.8% for UCL and 47.7% for a field
object. Slightly larger membership probability in LCC were found for Hipparcos
astrometry (57.2%). The star shows significant $\Delta\mu$ between various
catalogs, most likely due to the companion at 0.2” and supporting its physical
association. The companion detected in our study remained undetected by
Kouwenhoven et al. (2005). Considering the somewhat ambiguous membership, we
adopt the LCC age, but with the upper limit derived by isochrone fitting as
for an isolate field object.

HIP 58859 = HD 104839 B9 star confirmed member of LCC. New close binary
(separation 87 mas) from our observations.

TWA 24 = CD-58 4411 TWA 24 was included among the TWA members by Song et al.
(2003). The membership was however rejected by Mamajek (2005) and other
studies, because of the distance larger than 100 pc, putting the star in the
LCC association, in the background of TWA. The Gaia trigonometric parallax
confirms this assessment and our kinematic analysis yields a membership
probability to LCC of 99.9% probability. Both components of the TWA 24 system
are in Gaia DR2, confirming the physical association. The TESS periodogram
yields two significant periods of 2.198 and 5.41 d, respectively; they are
consistent with membership to the LCC association if they are interpreted as
due to the rotation of the primary and secondary, respectively. Messina et al.
(2010) provided a shorter period of 0.68 days. The 2.198d period from TESS is
compatible with the observed $V\leavevmode\nobreak\ \sin{i}$ of the primary
(17.1 km/s), the 5.41 is too long, while 0.68 would imply an unlikely
orientation very close to pole-on.

Available NaCo archive imaging already allowed us to classify TWA24B as a
physical companion before Gaia. The individual spectra gathered by Mentuch et
al. (2008), with very strong Lithium line in both components, represent an
additional confirmation of the physical link of the system. The RV difference
between the components in Gaia DR2 (8.0$\pm$3.9 km/s) is only marginally
significant due to the large errors. Archive HARPS RVs 171717Prog. ID
074.C-0037(A), 076.C-0010(A), 077.C-0012(A), 079.C-0046(A), PI E. Gunther
allow us to exclude the presence of close stellar companions around the
primary. An additional comoving object, likely a physical companion, is
identified in Gaia Dr2 at wider separation (12.8”). It results a very low mass
star from G band photometry. To our knowledge, it was not previously noticed.

Figure 23: Upper panel: Results from the SPOC analysis of TESS data for
TWA-24. Black dots are the pre-whitened TESS photometry phased at a period of
17.403 d; blue dots are running median values. The red line is a trapezoidal
transit model. The green points are residuals.

SPOC analysis of TESS data detected two possible transit candidates. However
in the following we will only consider that with period P=17.4 d and depth of
5362 ppm (see Figure 23), because an exam of the light curve of the other one
(P=10.672) shows that this has not the shape expected for a transiting object
unless it is very marginal. In Gaia DR2 catalogue there are not other stars
bright enough to be responsible for this transit but the two components of the
binary. While the separation of the binary is not negligible, the in
transit/out transit offset is very small and in a direction completely
different from the secondary. We tentatively assume that the transit is on the
primary. Given the stellar radius of the star (R=1.36 R⊙ using COND isochrones
by Baraffe et al. 2015), the radius of the transiting object is 0.94 RJ, that
given the very young age of the system is due to a planetary object with a
mass of $\sim$20-25 M⊕ (using models by Linder et al. 2019), that is a
Neptune-like one. Further analysis and data would be highly welcomed.

HIP 59603 = HD 106218  The kinematic analysis yields a probability of 99.9%
for LCC, adopting Chen et al. (2011) RV. There is a marginal RV difference
between Gaia DR2 and Chen et al. (2011). The newly detected companion is at
very close separation (87.5 mas).

HIP 60577 = HD 108035 Member of LCC, as also confirmed by our kinematic
analysis including Gaia DR2 data (membership probability 99.2%). TESS
periodogram gives both short and long periods; if we interpret the long period
(P=2.151 d) as due to rotation of the secondary, this is consistent with
membership to LCC. It is flagged as binary in WDS because of the presence of a
companion detected in observation performed between 1929 and 1944. The
companion was instead not detected in Hipparcos, while it was detected by
Gaia, with $\Delta$G =3.59 and astrometric parameters compatible with a
comoving object. The secondary is also detected in SPHERE images and resolved
for the first time into a tight equal mass pair (separation=108 mas, PA=203.4
degree, $\Delta K=0$) making the system triple. Since Gaia did not resolve the
secondary, the position should refer to the photocenter of the system in the
Gaia band pass. The system shows prominent X-ray emission (source 1RXS
J122500.3-501121), likely originating from the late type components.

HIP 61087 F6 star, confirmed member of LCC (probability 98.8%). The companion
detected with SPHERE is a new discovery, it was not detected in previous
observations by Janson et al. (2013). The RV of the star measured by Chen et
al. (2011) differ by 4.6$\pm$0.8 km/s with respect to Gaia and by 5.0$\pm$1.8
km/s with respect to Nordström et al. (2004). Two RV epochs separated by 3307