11institutetext: Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69,
D-53121 Bonn, Germany 22institutetext: Instituto de Radioastronomía y
Astrofísica, Universidad Nacional Autónoma de México, 58090, Morelia,
Michoacán, Mexico 33institutetext: Max-Planck-Institut für Astronomie,
Königstuhl 17, D-69117 Heidelberg, Germany

# Very long baseline interferometry detection of nearby ($<100$ pc) young
stars

Pilot observations

Sergio A. Dzib 11 Laurent Loinard 22 Ralf Launhardt 33 Jazmín Ordóñez-Toro 22

To increase the number of sources with very long baseline interferometry
(VLBI) astrometry available for comparison with the Gaia results, we have
observed 31 young stars with recently reported radio emission. These stars are
all in the Gaia DR3 catalog and were suggested, on the basis of conventional
interferometry observations, to be nonthermal radio emitters and are therefore
good candidates for VLBI detections. The observations were carried out with
the Very Long Baseline Array (VLBA) at two epochs separated by a few days and
yielded ten detections (a roughly $30$% detection rate). Using the astrometric
Gaia results, we extrapolated the target positions to the epochs of our radio
observations and compared them with the position of the radio sources. For
seven objects, the optical and radio positions are coincident within five
times their combined position errors. Three targets, however, have position
discrepancies above eight times the position errors, indicating different
emitting sources at optical and radio wavelengths. In one case, the VLBA
emission is very likely associated with a known companion of the primary
target. In the two other cases, we associated the VLBA emission with
previously unknown companions, but further observations will be needed to
confirm this.

###### Key Words.:

astrometry — techniques: interferometry — radiation mechanisms: non-thermal —
stars:kinematics

##  1 Introduction

Stellar radio astrometry with very long baseline interferometry (VLBI) at
centimeter wavelengths achieves angular resolutions of a few milli-arcseconds
and an astrometric accuracy on the order of $0.1$ mas (e.g., Reid & Honma,
2014). For young stars, VLBI radio astrometry has primarily been used to
better constrain star-forming region distances, which were poorly known at the
time or were debated due to discrepant results from indirect methods (e.g.,
Loinard et al., 2005, 2007; Dzib et al., 2010; Ortiz-León et al., 2017b). In
young stars, radio emission detectable with VLBI is nonthermal and is produced
in stellar coronae (Güdel, 2002) that typically extend up to a few radii and
remain unresolved with VLBI observations at centimeter wavelengths. Thus, the
astrometry derived from these radio observations directly traces the stellar
motion. Recently, the Gaia mission (Gaia Collaboration et al., 2016, 2021,
2023) has been delivering astrometry of over a billion stars, including many
nearby young stars. Not surprisingly, Gaia results have confirmed the
astrometric results previously reported using the VLBI technique (e.g., Ortiz-
León et al., 2018; Galli et al., 2019).

Gaia and VLBI astrometry accuracies are comparable, so the two techniques are
complementary (e.g. Dzib et al., 2021). Comparing trigonometric parallaxes
measured with both can constrain the Gaia zero-point offset (Lindegren et al.,
2018; Lindegren, 2020; Riess et al., 2018; Zinn et al., 2019; Xu et al.,
2019). Also, high precision VLBI astrometry allows for the linking of the
optical stellar reference frames to the radio extragalactic reference frame
since VLBI stellar astrometry is relative to them (Lestrade et al., 1999).
Recently, Lindegren (2020) analyzed 41 stars in the Gaia catalog with existing
VLBI stellar astrometry. He noticed discrepancies between the VLBI and Gaia-
derived astrometry and rejected 15 stars for further analysis. These rejected
stars are binary and multiple systems that do not comply with the linear
motion prescription currently used for Gaia astrometry. From the remaining 26
stars, he estimated the rotation of the Gaia bright stars ($G\leq 13$ mag)
reference frame with respect to the distant quasar frame (the so-called Gaia
Celestial Reference Frame). Lindegren (2020) recognized that stellar VLBI
astrometry can provide independent astrometric results that could be used to
test results from current and future Gaia releases. He also noted that for
this purpose, it will be necessary to re-observe stars with VLBI and identify
new stars that can be detected with the VLBI technique. There are, however,
few stars with suspected nonthermal radio emission that have not already been
observed with the VLBI technique.

Looking to increase the number of stars with both VLBI and Gaia astrometry,
the main goal of the present work is to obtain VLBI observations of nearby
young stars with suspected nonthermal radio emission and compare their
positions with those reported by Gaia. Here we report VLBI observations of 31
nearby ($<100$ pc) young stars ($<100$ Myr) that were recently reported to be
radio emitters by Launhardt et al. (2022). These authors also suggested that
the radio emission is likely of nonthermal nature and could thus be detected
with VLBI observations. The stars have optical counterparts, and most of them
are bright stars in the Gaia catalog. In Table LABEL:tab:tar, these stars are
listed together with their astrometric solution from Gaia DR3 (Gaia
Collaboration et al., 2021, 2023) and other known properties.

Table 1: Target sources.

Star ID | R.A. (J2016.0) | Dec. (J2016.0) | $\pi$ | Dist. | $\mu_{\alpha}\cdot\cos{(\rm\delta)}$ | $\mu_{\delta}$ | spT | age | age | G | Dates for Max. | Obs. | Phase | Separation
---|---|---|---|---|---|---|---|---|---|---|---|---|---|---
|  (∘) |  (∘) | (mas) | (pc) |  (mas yr-1) |  (mas yr-1) |  | (Myr) | ref. | (mag) | Parallax | Group | Cal. |  (∘)
BD+17 232 | $024.4144646604$ | $+018.5922019429$ | $18.95\pm 00.03$ | $52.77\pm 00.09$ | $73.79\pm 00.04$ |  $-43.35\pm 00.03$. | K3VE | 10 | 1 | 10.18 | Jan 16 - Jul 18 | A | J0139+1753 | 0.9
HIP 12545 | $040.3582211403$ | $+005.9881913726$ | $22.58\pm 00.02$ | $44.30\pm 00.04$ | $79.07\pm 00.02$ |  $-56.70\pm 00.02$. | K6Ve | 10 | 2 | 9.82 | Feb 2 - Aug 4 | B | J0239+0416 | 1.8
HIP 12635 | $040.5877605186$ | $+038.6220662083$ | $20.09\pm 00.02$ | $49.78\pm 00.05$ | $82.11\pm 00.02$ |  $-107.55\pm 00.02$. | K2V | 149 | 3 | 9.92 | Feb 2 - Aug 5 | B | J0253+3835 | 2.1
_V 875 Per_ | $043.0736090107$ | $+036.2798896433$ | $04.07\pm 00.02$ | $245.92\pm 01.08$ | $52.87\pm 00.02$ | $-38.55\pm 00.02$ | K2IV | 63 | 4 | 10.45 | Feb 4 - Aug 7 | B | J0246+3536 | 1.4
TYC 3301-2585-1aa$a$Hereafter TYC 3301.  | $043.9317736370$ | $+047.7795639137$ | $19.97\pm 00.02$ | $50.07\pm 00.05$ | $85.50\pm 00.02$ | $-79.53\pm 00.02$ | K5Ve | 42 | 3 | 10.71 | Feb 5 - Aug 8 | B | J0303+4716 | 1.4
_HD 22213_ | $053.5684734519$ | $-012.0688463141$ | $19.46\pm 00.02$ | $51.38\pm 00.05$ | $74.46\pm 00.02$ | $-34.36\pm 00.01$ | G7V | 45 | 3 | 8.58 | Feb 15 - Aug 18 | C | J0336-1302 | 1.1
HD 23524bb$b$Star in Gaia DR3 without parallax and proper motions. Position is from Gaia DR3, and parallax and proper motions are taken from the Hipparcos catalog (van Leeuwen, 2007). | $057.0963206572$ | $+052.0375180643$ | $19.35\pm 01.62$ | $51.85\pm 04.50$ | $61.87\pm 01.98$ | $-70.99\pm 01.67$ | K1V | 42 | 3 | 8.62 | Feb 18 - Aug 21 | D | J0350+5138 | 0.5
HD 24681 | $058.8352088930$ | $-001.7296295356$ | $17.94\pm 00.02$ | $55.73\pm 00.06$ | $43.50\pm 00.02$ | $-91.47\pm 00.02$ | G8V | 149 | 3 | 8.79 | Feb 20 - Aug 23 | C | J0351-0301 | 1.7
_HD 285281_ | $060.1294810416$ | $+019.5890872952$ | $07.34\pm 00.02$ | $136.22\pm 00.43$ | $04.04\pm 00.03$ | $-12.36\pm 00.02$ | K1 | 1.5 | 4 | 9.97 | Feb 21 - Aug 24 | D | J0401+2110 | 1.6
HD 284135cc$c$Star in Gaia DR3 catalog without parallax and proper motions. Position is from Gaia DR3, proper motions are from Tycho-2 catalog (Høg et al., 2000) and distance information from Carpenter et al. (2009). | $061.4191041862$ | $+022.8032778542$ | $12.50\pm 01.00$ | $80.69\pm 06.55$ | $05.50\pm 01.00$ | $-14.70\pm 01.00$ | G3V: | 1.5 | 1 | 9.17 | Feb 22 - Aug 25 | D | J0412+2305 | 1.7
HD 31281 | $073.7900713168$ | $+018.4419023094$ | $08.25\pm 00.03$ | $121.29\pm 00.43$ | $-2.51\pm 00.04$ | $-16.10\pm 00.02$ | G1V: | 1.5 | 5 | 9.06 | Mar 6 - Sep 6 | E | J0449+1754 | 1.5
BD-08 995 | $074.7023574082$ | $-008.7277965955$ | $11.47\pm 00.02$ | $87.15\pm 00.12$ | $25.83\pm 00.01$ | $-21.61\pm 00.01$ | K0V | 42 | 3,6 | 10.00 | Mar 7 - Sep 7 | E | J0457-0849 | 0.4
HD 286264 | $075.2054343407$ | $+015.4499325911$ | $18.83\pm 00.02$ | $53.09\pm 00.04$ | $18.31\pm 00.02$ | $-58.72\pm 00.01$ | K2IV | 24 | 3 | 10.37 | Mar 7 - Sep 8 | E | J0459+1438 | 0.9
HD 293857dd$d$Star in Gaia DR3 catalog without parallax and proper motions. Position is from Gaia DR3, proper motions are from the Tycho-2 catalog, and distance information is from da Silva et al. (2009). The last column lists the angular separation between the target source and the phase calibrator. | $077.7903401655$ | $-004.1820005272$ | $12.82\pm 01.00$ | $78.45\pm 06.24$ | $18.70\pm 01.00$ | $-57.50\pm 01.00$ | G8V | 24 | 3 | 9.04 | Mar 10 - Sep 10 | E | J0505-0419 | 1.3
TYC 713-661-1 | $084.2085922983$ | $+013.6317744334$ | $17.61\pm 00.02$ | $56.77\pm 00.06$ | $04.90\pm 00.02$ | $-109.84\pm 00.01$ | K0V | 149 | 3 | 10.16 | Mar 16 - Sep 16 | F | J0539+1433 | 1.2
TYC 5925-1547-1 | $084.8465611076$ | $-019.5583471893$ | $14.30\pm 00.17$ | $69.95\pm 00.82$ | $21.30\pm 00.13$ | $-48.63\pm 00.16$ | K1V | 150 | 3,6 | 10.28 | Mar 16 - Sep 17 | F | J0536-2005 | 0.9
_AI Lep_ | $085.0864740191$ | $-019.6697559694$ | $13.86\pm 00.01$ | $72.16\pm 00.07$ | $19.32\pm 00.01$ | $-11.69\pm 00.01$ | G2V | 42 | 3 | 8.94 | Mar 17 - Sep 17 | F | J0536-2005 | 1.0
_HD 62237_ | $115.6106888262$ | $-016.2834934596$ | $08.02\pm 00.10$ | $124.76\pm 01.60$ | $-9.44\pm 00.09$ | $-13.21\pm 00.09$ | G5V | 42 | 3,6 | 9.39 | Apr 14 - Oct 16 | G | J0746-1555 | 1.0
_SAO 135659_ | $123.4623602751$ | $-007.6403837042$ | $19.97\pm 00.20$ | $50.07\pm 00.51$ | $-32.15\pm 00.23$ | $-32.31\pm 00.19$ | K0 | 42 | 3 | 9.01 | Apr 22 - Oct 23 | G | J0808-0751 | 1.4
WDS 09035+3750 B | $135.8623936283$ | $+037.8407266874$ | $32.75\pm 00.12$ | $30.54\pm 00.11$ | $-65.85\pm 00.11$ | $-163.11\pm 00.08$ | M3-6V | 42 | 3 | 6.91 | May 4 - Nov 5 | H | J0850+3747 | 2.6
_HD 82159_ | $142.6484056528$ | $+010.6016777097$ | $28.23\pm 00.04$ | $35.43\pm 00.05$ | $-200.06\pm 00.04$ | $-12.88\pm 00.03$ | G9V | 150 | 7 | 8.41 | May 11 - Nov 11 | H | J0935+0915 | 1.8
_HD 82558_ | $143.1054112092$ | $-011.1844843913$ | $54.74\pm 00.02$ | $18.27\pm 00.01$ | $-248.04\pm 00.02$ | $34.28\pm 00.02$ | K0V | 43 | 8 | 7.50 | May 12 - Nov 12 | H | J0933-1139 | 0.6
GJ 2079 | $153.5792772472$ | $+021.0741011408$ | $42.74\pm 01.00$ | $23.41\pm 00.55$ | $-119.55\pm 01.00$ | $-191.19\pm 01.00$ | M0.7V | 24 | 3 | 9.36 | May 23 - Nov 23 | H | J1016+2037 | 0.7
HD 135363 | $226.9819941355$ | $+076.2014908036$ | $33.71\pm 00.05$ | $29.66\pm 00.05$ | $-127.19\pm 00.08$ | $165.47\pm 00.07$ | G5V | 45 | 9 | 8.39 | Feb 8 - Aug 11 | I | J1448+7601 | 1.2
UCAC4 832-014013 | $226.9860077519$ | $+076.2338226259$ | $33.71\pm 00.03$ | $29.66\pm 00.02$ | $-127.64\pm 00.04$ | $163.10\pm 00.04$ | M4.5V | … | … | 12.40 | Feb 8 - Aug 11 | I | J1448+7601 | 1.2
_HD 199143_ | $313.9489042861$ | $-017.1144445844$ | $21.78\pm 00.03$ | $45.91\pm 00.07$ | $55.73\pm 00.03$ | $-58.67\pm 00.03$ | F8 | 24 | 3 | 7.32 | May 2 - Nov 3 | J | J2107-1708 | 2.7
HD 358623 | $314.0116622465$ | $-017.1818893172$ | $21.70\pm 00.02$ | $46.08\pm 00.05$ | $57.32\pm 00.02$ | $-62.20\pm 00.02$ | K6Ve | 24 | 3 | 9.90* | May 2 - Nov 3 | J | J2107-1708 | 2.7
SAO 50350 | $315.1960688588$ | $+045.5029486058$ | $19.31\pm 00.01$ | $51.78\pm 00.03$ | $38.84\pm 00.01$ | $08.50\pm 00.01$ | F8 | 200 | 8 | 8.36 | May 4 - Nov 4 | J | J2102+4702 | 1.6
_GJ 4199_ | $322.7577934607$ | $+023.3347377313$ | $41.29\pm 00.02$ | $24.22\pm 00.01$ | $134.65\pm 00.01$ | $-144.89\pm 00.01$ | K5Ve | 149 | 3 | 8.85 | May 11 - Nov 12 | K | J2125+2442 | 1.9
SAO 51891 | $335.0299106318$ | $+049.5033014554$ | $29.07\pm 00.02$ | $34.40\pm 00.02$ | $93.24\pm 00.02$ | $07.37\pm 00.02$ | K1V | 37 | 8 | 8.22 | May 24 - Nov 24 | K | J2201+5048 | 3.2
SAO 108142 | $341.1734897431$ | $+017.9047178334$ | $20.85\pm 00.07$ | $47.97\pm 00.17$ | $82.49\pm 00.08$ | $-82.11\pm 00.07$ | K0 | 149 | 3 | 9.17 | May 31 - Dec 1 | K | J2242+1741 | 0.5
111Stars with a detected radio source in this work are highlighted. The
positions, parallax, and proper motion were obtained from the Gaia DR3 catalog
(Gaia Collaboration et al., 2021, 2023), with some exceptions.

Age references: (1) Galicher et al. (2016), (2) Weise et al. (2010), (3) Bell
et al. (2015), (4) Carpenter et al. (2009), (5) Kenyon & Hartmann (1995), (6)
da Silva et al. (2009), (7) Desidera et al. (2015), (8) Stanford-Moore et al.
(2020), and (9) Montes et al. (2001).

##  2 Observations

The observations were obtained at 20 epochs in the first half of 2021 with the
Very Long Baseline Array (VLBA) using the C-band receivers ($\lambda=6.0$ cm;
$\nu=4.8$ GHz). This band was selected as a compromise between the angular
resolution, strength of radio fluxes, and the high sensitivity offered by this
band. We used the newly offered recording bit rate of 4 Gbps, which improves
the sensitivity of the array by a factor of $\sqrt{2}$ with respect to the
previous system for an equal observing time. For the present observations,
about 45 minutes were spent on each target source per observation, expecting a
noise level of 19 $\mu$Jy beam-1, while with the previous system was 27
$\mu$Jy beam-1.

###  2.1 Observation strategy

The radio emission from the stars in our sample is known to be rapidly
variable, and thus we scheduled two observing runs separated by a few days to
increase the chances of detection. The observations were scheduled close to
the maximum trigonometric parallax extension in right ascension so that they
could be combined with future observations to measure the trigonometric
parallaxes of the detected sources. The observing dates are listed
individually for all 31 radio stars in Table 2.

As the target sources are weak, phase-referencing was used (e.g., Lestrade et
al., 1990; Beasley & Conway, 1995). This means that the target observing scans
(2 min.) are bracketed with scans (1 min.) on a relatively bright quasar with
an angular separation of $\lesssim 3^{\circ}$. The phase corrections
determined for the quasar were transferred to the target so the target source
positions are referenced to the quasars. The use of this technique, including
antenna slewing time, increases the total telescope time by 60% (for instance,
for 30 minutes on source, we would need 48 minutes in total). Additionally, to
reduce the impact of errors in the troposphere models used by the correlator,
we observed so-called geodetic blocks (i.e., $\sim 15$-minute blocks at the
beginning and/or middle and/or end of the observation) dedicated to
observations of five to six bright quasars distributed over the entire sky
(Reid & Brunthaler, 2004).

To reduce the overheads caused by the geodetic blocks, we organized the
targets into 11 groups of sources with similar right ascensions (and therefore
also similar dates of maximum parallax elongation). We named these groups
using alphabetic letters from A to K. We have one group of one target, three
groups of two targets, four groups of three targets, and three groups of four
targets (see Table 2). For the groups with one and two targets, we used two
geodetic blocks (at the beginning and the end of the session). For groups with
three or four targets, we used three geodetic blocks (at the beginning,
middle, and end of the sessions). A second advantage of the grouping strategy
is that we could schedule the observations such that the $(u,\,v)$ plane
coverage is maximized. This was achieved by intertwining several blocks of 24
minutes of the different sources with similar right ascension throughout the
duration of the session.

Table 2: Observed epochs.

| Epoch 1 | Epoch 2 | No. observed
---|---|---|---
Group | yyyy/mm/dd | yyyy/mm/dd | targets
A | 2021/07/15 | 2021/07/17 | 1
B | 2021/01/22 | 2021/01/24 | 4
C | 2021/02/10 | 2021/02/13 | 2
D | 2021/02/12 | 2021/02/14 | 3
E | 2021/03/09 | 2021/03/10 | 4
F | 2021/03/12 | 2021/03/13 | 3
G | 2021/04/14 | 2021/04/16 | 2
H | 2021/05/16 | 2021/05/23 | 4
I | 2021/01/29 | 2021/01/30 | 2
J | 2021/05/07 | 2021/05/10 | 3
K | 2021/05/21 | 2021/05/22 | 3
222Columns are (from left to right) the name of the group, the two epochs of
the observation in civil dates, and the number of targets included in the
corresponding group.

###  2.2 Calibration and imaging

The data were edited and calibrated using the software Astronomical Image
Processing System (AIPS; Greisen, 2003). After a detailed inspection, flawed
data were removed. The data were calibrated following the standard schemes
plus a calibration of the group delay (i.e., the rate of phase change with
frequency). The calibration of the group delay was used to further reduce
systematic errors caused by tropospheric zenith delays and ionospheric content
delays (Reid & Brunthaler, 2004). With this purpose, we additionally observed
ICRF quasars distributed over the sky over a wide range of frequencies (also
known as the geodetic blocks; Reid & Brunthaler, 2004). First, the geodetic
block data were calibrated using a standard procedure. Second, the clock
delays and tropospheric terms were derived using the AIPS task DELZN and
applied to the target data set using the CALIB. A detailed description of the
calibration strategy has been summarized by Loinard et al. (2007); Dzib et al.
(2010) and Ortiz-León et al. (2017b).

Figure 1: VLBA images of the radio source related to V 875 Per. The images are
centered in the position of the radio source as detected in the first epoch
(see Table LABEL:tab:am). Images, from left to right, correspond to epochs 1,
2, and the combination of both epochs. Contour levels are –3, 3, 6, 9, 12, and
15 times the noise level of the image as listed in Table 3. The yellow open
ellipse in the bottom-left corner represents the size of the synthesized beam
of the image as listed in Table 3. The predicted optical position in epochs 1
and 2 are shown as red and cyan crosses, respectively.

After calibration, the data were imaged in two steps. First, for the detection
experiment, the $(u,\,v)$ data of the two epochs for each target were combined
using the task DBCON in AIPS. This was done with the purpose of decreasing the
noise level in the images (by a factor of $\sqrt{2}$) and increasing the
chances of detections. Some of the target sources have very high proper motion
(up to $250$ mas yr-1, see Table LABEL:tab:tar), so before concatenating the
data of the two observed epochs, we corrected the position of the second epoch
to coincide with that of the first epoch. This procedure is described in
Appendix A. After the concatenation of the two epochs, we imaged large areas
around the expected positions of the target sources. These images are
$2^{\prime\prime}\times 2^{\prime\prime}$ with 4096 pixels, 500$\,\mu$as each,
along both directions, and they used a natural weighting scheme (Robust=5 in
AIPS).333The robust scheme weights the visibilities considering the length of
the baselines. A natural weighting scheme can be set with Robust=5 in AIPS and
gives more weight to the interferometer’s shortest baselines. A uniform scheme
weights all the baselines equally and is set with Robust=–5 in AIPS. The first
scheme resulted in images with better sensitivity, and the second scheme
resulted in the smallest synthesized beam but with higher noise.

In the second step, a new image was obtained for each source around the peak
emission on the first image. These new images are $0.1^{\prime\prime}\times
0.1^{\prime\prime}$ with 1024 pixels, 100$\,\mu$as each, and they also used
natural weighting. We consider a source detection when the peak value has a
signal-to-noise ratio (SNR) $>6.5$ in this last image (see discussion by
Forbrich et al., 2021). The final images were produced at the position of the
detected sources in each of the data sets of individual epochs using the
parameters of the previous image. Only in two cases was the detected emission
strong enough (with S/N$>20$) that we could produce final images using an
intermediate weighting scheme between natural and uniform (Robust=0 in AIPS;
Briggs, 1995) and obtain a better position measurement. The rms noise levels
in the final images are $\sim 20$ $\,\mu$Jy beam-1, similar to the VLA images
reported by Launhardt et al. (2022). The flux densities and positions of the
detected sources were measured using a two-dimensional Gaussian fitting
procedure (task JMFIT in AIPS). Statistical errors resulting from JMFIT are
given for each detection in columns (4) and (6) in Table LABEL:tab:am.
Additionally, systematic residual phase errors caused by the transfer of the
calibration from the phase calibrator to the target were expected. These
errors are on the order of $0.1$ mas per degree of the angular separation
(e.g., Reid et al., 2017). Given that the maximum angular separation between
our targets and their phase calibrators is 3$\hbox to0.0pt{.\hss}^{\circ}2$,
our systematic errors are expected to be $\leq 0.32\,$mas. These systematic
errors were calculated specifically for each detected target source. Overall,
we note that the position errors for weak sources ($S_{\nu,{\rm
peak}}<200\,\mu$Jy beam-1) are dominated by the statistical errors, and for
the strongest sources, they are dominated by the systematic residual phase
errors.

##  3 Results

Radio emission was detected for ten of our 31 targets. The properties of the
final images and the flux densities are listed in Table 3, and the positions
of the radio sources are given in Table LABEL:tab:am.

In the plots of Figure 1, we show an example of a detected radio source, and
the remaining images of other detected radio sources are shown in Figure 3. In
most cases, single sources are detected in the images. A possible exception is
source HD 199143, where two emission peaks separated by a few milli-arcseconds
are detected in the first epoch. Given the low signal-to-noise ratio of the
second peak, however, it is not possible to clearly establish that the second
peak is a real source. Sources HD 82558 and HD 82159 are single sources in
both epochs, but their positions change appreciably from one epoch to the next
due to their high proper motion. As mentioned earlier, the combined image
aligns both epochs to the position in the first image. The radio sources are
rapidly variable, and in several cases, their flux density changes by a factor
larger than two in a few days. This high variability was also noticed in the
VLA observation (Launhardt et al., 2022) and is typical of stellar sources
with magnetic activity (e.g., Güdel, 2002).

Table 3: Image properties and measured fluxes of detected radio sources.

|  |  | Synthesized beam |  |
---|---|---|---|---|---
|  | Noise | $\theta_{\rm maj}\times\theta_{\rm min};{\rm P.A.}$ | $S_{\nu,{\rm Int.}}$ | $S_{\nu,{\rm Peak.}}$
Star ID | Epoch |  ($\mu$Jy bm-1) |  (mas $\times$ mas; ∘) |  ($\mu$Jy) |  ($\mu$Jy bm-1)
V 875 Per | B1 | 26 |  3.1$\times$1.3; $-5.9$ | $1042\pm 66$ | $585\pm 26$
| B2 | 42 |  4.5$\times$1.7; $-9.0$ | $1225\pm 86$ | $885\pm 42$
| B1+B2 | 23 |  3.3$\times$1.4; $-6.5$ | $1156\pm 56$ | $667\pm 23$
HD 22213 | C1 | 24 |  3.6$\times$1.3; $+1.5$ | $338\pm 88$ | $142\pm 24$
| C2 | 26 |  3.5$\times$1.1; $+9.7$ | $710\pm 54$ | $552\pm 26$
| C1+C2 | 17 |  $3.9\times 1.4$; $+9.9$ | $450\pm 40$ | $297\pm 17$
HD 285281 | D1 | 17 |  4.0$\times$1.7; $-6.1$ | $270\pm 62$ | $96\pm 17$
| D2 | 24 |  3.9$\times$1.6; $-3.6$ | $152\pm 47$ | $141\pm 26$
| D1+D2 | 13 |  3.9$\times$1.6; $-4.9$ | $170\pm 33$ | $103\pm 13$
AI Lep | F1 | 26 |  4.7$\times$1.8; $+6.7$ | $541\pm 59$ | $387\pm 27$
| F2 | 15 |  5.0$\times$1.8; $+0.1$ | $105\pm 33$ | $69\pm 15$
| F1+F2 | 12 |  4.8$\times$1.8; $+1.9$ | $220\pm 36$ | $144\pm 16$
HD 62237 | G1 | 21 |  4.3$\times$1.6; $-1.8$ | $415\pm 65$ | $190\pm 21$
| G2 | 19 |  4.3$\times$1.6; $-0.2$ | $172\pm 35$ | $164\pm 19$
| G1+G2 | 15 |  4.3$\times$1.6; $-1.0$ | $280\pm 38$ | $162\pm 15$
SAO 135659 | G1 | 20 |  4.2$\times$1.6; $-2.7$ | $648\pm 41$ | $540\pm 21$
| G2 | 22 |  4.2$\times$1.7; $-0.4$ | $524\pm 41$ | $444\pm 22$
| G1+G2 | 16 |  4.2$\times$1.7; $-1.6$ | $586\pm 31$ | $489\pm 16$
HD 82159 | H1 | 66 |  3.9$\times$1.6; $+3.2$ | $980\pm 138$ | $734\pm 66$
| H2 | 47 |  4.0$\times$1.7; $+1.5$ | $470\pm 76$ | $503\pm 47$
| H1+H2 | 38 |  3.9$\times$1.6; $+2.2$ | $669\pm 74$ | $577\pm 39$
HD 82558 | H1 | 38 |  4.0$\times$1.6; $+0.8$ | $462\pm 77$ | $362\pm 38$
| H2 | 25 |  4.2$\times$1.6; $-1.2$ | $219\pm 40$ | $230\pm 25$
| H1+H2 | 22 |  4.1$\times$1.6; $-0.3$ | $365\pm 65$ | $260\pm 27$
HD 199143 | J1 | 30 |  5.1$\times$1.9; $-5.4$ | $280\pm 81$ | $155\pm 31$
| J2 | 20 |  5.7$\times$1.6; $-9.7$ | $162\pm 57$ | $81\pm 20$
| J1+J2 | 17 |  5.4$\times$1.7; $-8.2$ | $221\pm 64$ | $113\pm 17$
GJ 4199 | K1 | 19 |  4.0$\times$1.8; $-15.9$ | $350\pm 47$ | $210\pm 19$
| K2 | 24 |  4.0$\times$1.8; $-12.0$ | $273\pm 51$ | $192\pm 24$
| K1+K2 | 14 |  4.0$\times$1.8; $-13.8$ | $306\pm 33$ | $204\pm 15$
444Columns are (from left to right) the target star name, the epoch (for this
table the epoch ID is constructed using the group name and the epoch number
from Table 2), the image noise level, the synthesized beam size in units of
milli-arcseconds, and the integrated and peak flux densities.

Table 4: Astrometry of detected radio sources.

|  |  | $\sigma_{\alpha}$ |  | $\sigma_{\delta}$ | $\Delta\alpha^{*}_{\rm radio}$ | $\Delta\delta_{\rm radio}$ | $\Delta\theta_{\rm radio}$ | $\Delta\alpha^{*}$ | $\Delta\delta$ | $\Delta\theta$ | $\sigma_{\rm sys}$ | $\Delta\lambda$
---|---|---|---|---|---|---|---|---|---|---|---|---|---
Star ID | Epoch | R.A. |  ($\mu$s) | Dec |  ($\mu$as) |  (′′) |  (′′) |  (′′) | (mas) | (mas) | (mas) | (mas) | (au)
(1) | (2) | (3) | (4) | (5) | (6) | (7) | (8) | (9) | (10) | (11) | (12) | (13) | (14)
V 875 Per | B1 | $02^{\rm h}52^{\rm m}17\mbox{${}^{\rm s}\mskip-7.6mu.\,$}687958$ | 3.3 | $+36^{\circ}16^{\prime}47\mbox{${}^{\prime\prime}\mskip-7.6mu.\,$}407545$ | 60 | $0.06\pm 0.10$ | $0.02\pm 0.07$ | $0.06\pm 0.10$ | $0.19\pm 0.16$ | $-0.61\pm 0.12$ | $0.6\pm 0.1$ | 0.14 | $0.10\pm 0.06$
| B2 | $02^{\rm h}52^{\rm m}17\mbox{${}^{\rm s}\mskip-7.6mu.\,$}687982$ | 3.8 | $+36^{\circ}16^{\prime}47\mbox{${}^{\prime\prime}\mskip-7.6mu.\,$}407500$ | 85 |  |  |  |  |  |  |  |
| B1+B2 | $02^{\rm h}52^{\rm m}17\mbox{${}^{\rm s}\mskip-7.6mu.\,$}687968$ | 2.7 | $+36^{\circ}16^{\prime}47\mbox{${}^{\prime\prime}\mskip-7.6mu.\,$}407501$ | 49 |  |  |  |  |  |  |  |
HD 22213 | C1 | $03^{\rm h}34^{\rm m}16\mbox{${}^{\rm s}\mskip-7.6mu.\,$}458216$ | 22 | $-12^{\circ}04^{\prime}08\mbox{${}^{\prime\prime}\mskip-7.6mu.\,$}026964$ | 374 | $-0.09\pm 0.10$ | $-0.07\pm 0.14$ | $0.11\pm 0.12$ | $-1.0\pm 0.3$ | $+0.2\pm 0.4$ | $1.1\pm 0.3$ | 0.11 | $0.05\pm 0.02$
| C2 | $03^{\rm h}34^{\rm m}16\mbox{${}^{\rm s}\mskip-7.6mu.\,$}458292$ | 1.9 | $-12^{\circ}04^{\prime}08\mbox{${}^{\prime\prime}\mskip-7.6mu.\,$}026581$ | 79 |  |  |  |  |  |  |  |
| C1+C2 | $03^{\rm h}34^{\rm m}16\mbox{${}^{\rm s}\mskip-7.6mu.\,$}458287$ | 2.5 | $-12^{\circ}04^{\prime}08\mbox{${}^{\prime\prime}\mskip-7.6mu.\,$}026752$ | 98 |  |  |  |  |  |  |  |
HD 285281 | D1 | $04^{\rm h}00^{\rm m}31\mbox{${}^{\rm s}\mskip-7.6mu.\,$}076436$ | 14 | $+19^{\circ}35^{\prime}20\mbox{${}^{\prime\prime}\mskip-7.6mu.\,$}650376$ | 580 | $-0.19\pm 0.18$ | $0.04\pm 0.19$ | $0.19\pm 0.18$ | $+0.3\pm 0.2$ | $+0.8\pm 0.6$ | $0.9\pm 0.5$ | 0.16 | $0.12\pm 0.06$
| D2 | $04^{\rm h}00^{\rm m}31\mbox{${}^{\rm s}\mskip-7.6mu.\,$}076421$ | 8.6 | $+19^{\circ}35^{\prime}20\mbox{${}^{\prime\prime}\mskip-7.6mu.\,$}648346$ | 332 |  |  |  |  |  |  |  |
| D1+D2 | $04^{\rm h}00^{\rm m}31\mbox{${}^{\rm s}\mskip-7.6mu.\,$}076421$ | 7.2 | $+19^{\circ}35^{\prime}20\mbox{${}^{\prime\prime}\mskip-7.6mu.\,$}649016$ | 317 |  |  |  |  |  |  |  |
AI Lep | F1 | $05^{\rm h}40^{\rm m}20\mbox{${}^{\rm s}\mskip-7.6mu.\,$}759902$ | 4.8 | $-19^{\circ}40^{\prime}11\mbox{${}^{\prime\prime}\mskip-7.6mu.\,$}183095$ | 152 | $0.55\pm 0.37$ | $0.32\pm 0.42$ | $0.64\pm 0.39$ | $+0.2\pm 0.1$ | $+0.4\pm 0.2$ | $0.5\pm 0.1$ | 0.10 | $0.04\pm 0.01$
| F2 | $05^{\rm h}40^{\rm m}20\mbox{${}^{\rm s}\mskip-7.6mu.\,$}759861$ | 18 | $-19^{\circ}40^{\prime}11\mbox{${}^{\prime\prime}\mskip-7.6mu.\,$}182392$ | 465 |  |  |  |  |  |  |  |
| F1+F2 | $05^{\rm h}40^{\rm m}20\mbox{${}^{\rm s}\mskip-7.6mu.\,$}759892$ | 8.6 | $-19^{\circ}40^{\prime}11\mbox{${}^{\prime\prime}\mskip-7.6mu.\,$}182835$ | 238 |  |  |  |  |  |  |  |
HD 62237 | G1 | $07^{\rm h}42^{\rm m}26\mbox{${}^{\rm s}\mskip-7.6mu.\,$}562160$ | 8.4 | $-16^{\circ}17^{\prime}00\mbox{${}^{\prime\prime}\mskip-7.6mu.\,$}668793$ | 321 | $-0.11\pm 0.12$ | $0.11\pm 0.23$ | $0.16\pm 0.19$ | $+12.3\pm 0.5$ | $-23.6\pm 0.6$ | $26.6\pm 0.6$ | 0.10 | $3.32\pm 0.08$
| G2 | $07^{\rm h}42^{\rm m}26\mbox{${}^{\rm s}\mskip-7.6mu.\,$}562170$ | 5.8 | $-16^{\circ}17^{\prime}00\mbox{${}^{\prime\prime}\mskip-7.6mu.\,$}668809$ | 230 |  |  |  |  |  |  |  |
| G1+G2 | $07^{\rm h}42^{\rm m}26\mbox{${}^{\rm s}\mskip-7.6mu.\,$}562169$ | 6.0 | $-16^{\circ}17^{\prime}00\mbox{${}^{\prime\prime}\mskip-7.6mu.\,$}668891$ | 214 |  |  |  |  |  |  |  |
SAO 135659 | G1 | $08^{\rm h}13^{\rm m}50\mbox{${}^{\rm s}\mskip-7.6mu.\,$}953026$ | 1.9 | $-07^{\circ}38^{\prime}25\mbox{${}^{\prime\prime}\mskip-7.6mu.\,$}622119$ | 76 | $-0.04\pm 0.03$ | $0.04\pm 0.03$ | $0.06\pm 0.03$ | $-10.5\pm 1.2$ | $-72.6\pm 1.0$ | $73.3\pm 1.0$ | 0.14 | $3.67\pm 0.06$
| G2 | $08^{\rm h}13^{\rm m}50\mbox{${}^{\rm s}\mskip-7.6mu.\,$}953037$ | 2.3 | $-07^{\circ}38^{\prime}25\mbox{${}^{\prime\prime}\mskip-7.6mu.\,$}622112$ | 98 |  |  |  |  |  |  |  |
| G1+G2 | $08^{\rm h}13^{\rm m}50\mbox{${}^{\rm s}\mskip-7.6mu.\,$}953031$ | 1.7 | $-07^{\circ}38^{\prime}25\mbox{${}^{\prime\prime}\mskip-7.6mu.\,$}622091$ | 65 |  |  |  |  |  |  |  |
HD 82159 | H1 | $09^{\rm h}30^{\rm m}35\mbox{${}^{\rm s}\mskip-7.6mu.\,$}542606$ | 5.5 | $+10^{\circ}36^{\prime}05\mbox{${}^{\prime\prime}\mskip-7.6mu.\,$}979934$ | 151 | $-0.04\pm 0.03$ | $0.02\pm 0.03$ | $0.04\pm 0.03$ | $-1.24\pm 0.2$ | $+0.2\pm 0.2$ | $1.3\pm 0.2$ | 0.18 | $0.05\pm 0.01$
| H2 | $09^{\rm h}30^{\rm m}35\mbox{${}^{\rm s}\mskip-7.6mu.\,$}542401$ | 4.6 | $+10^{\circ}36^{\prime}05\mbox{${}^{\prime\prime}\mskip-7.6mu.\,$}979351$ | 145 |  |  |  |  |  |  |  |
| H1+H2 | … | … | … | … |  |  |  |  |  |  |  |
HD 82558 | H1 | $09^{\rm h}32^{\rm m}25\mbox{${}^{\rm s}\mskip-7.6mu.\,$}204609$ | 5.9 | $-11^{\circ}11^{\prime}03\mbox{${}^{\prime\prime}\mskip-7.6mu.\,$}942747$ | 185 | $-0.07\pm 0.29$ | $0.03\pm 0.45$ | $0.07\pm 0.31$ | $-0.6\pm 0.2$ | $+0.2\pm 0.2$ | $0.7\pm 0.2$ | 0.06 | $0.010\pm 0.003$
| H2 | $09^{\rm h}32^{\rm m}25\mbox{${}^{\rm s}\mskip-7.6mu.\,$}204348$ | 4.9 | $-11^{\circ}11^{\prime}03\mbox{${}^{\prime\prime}\mskip-7.6mu.\,$}940469$ | 183 |  |  |  |  |  |  |  |
| H1+H2 | … | … | … | … |  |  |  |  |  |  |  |
HD 199143 | J1 | $20^{\rm h}55^{\rm m}47\mbox{${}^{\rm s}\mskip-7.6mu.\,$}759095$ | 20 | $-17^{\circ}06^{\prime}52\mbox{${}^{\prime\prime}\mskip-7.6mu.\,$}304352$ | 441 | $0.47\pm 0.15$ | $0.48\pm 0.15$ | $0.67\pm 0.15$ | $-2.6\pm 0.3$ | $+3.6\pm 0.5$ | $4.4\pm 0.4$ | 0.27 | $0.20\pm 0.02$
| J2 | $20^{\rm h}55^{\rm m}47\mbox{${}^{\rm s}\mskip-7.6mu.\,$}759173$ | 24 | $-17^{\circ}06^{\prime}52\mbox{${}^{\prime\prime}\mskip-7.6mu.\,$}305210$ | 608 |  |  |  |  |  |  |  |
| J1+J2 | $20^{\rm h}55^{\rm m}47\mbox{${}^{\rm s}\mskip-7.6mu.\,$}759121$ | 15 | $-17^{\circ}06^{\prime}52\mbox{${}^{\prime\prime}\mskip-7.6mu.\,$}304406$ | 333 |  |  |  |  |  |  |  |
GJ 4199 | K1 | $21^{\rm h}31^{\rm m}01\mbox{${}^{\rm s}\mskip-7.6mu.\,$}925936$ | 9.2 | $+23^{\circ}20^{\prime}04\mbox{${}^{\prime\prime}\mskip-7.6mu.\,$}291167$ | 144 | $0.05\pm 0.05$ | $-0.06\pm 0.05$ | $0.08\pm 0.05$ | $0.6\pm 0.1$ | $+1.3\pm 0.2$ | $1.4\pm 0.2$ | 0.19 | $0.03\pm 0.01$
| K2 | $21^{\rm h}31^{\rm m}01\mbox{${}^{\rm s}\mskip-7.6mu.\,$}925974$ | 11 | $+23^{\circ}20^{\prime}04\mbox{${}^{\prime\prime}\mskip-7.6mu.\,$}291129$ | 190 |  |  |  |  |  |  |  |
| K1+K2 | $21^{\rm h}31^{\rm m}01\mbox{${}^{\rm s}\mskip-7.6mu.\,$}925953$ | 6.7 | $+23^{\circ}20^{\prime}04\mbox{${}^{\prime\prime}\mskip-7.6mu.\,$}291177$ | 112 |  |  |  |  |  |  |  |

555Columns are (from left to right) the target name, epoch ID as in Table 3,
measured positions in the VLBA images and their statistical errors derived
from JMFIT, and the estimated offset between the radio source detected with
the VLA and the radio source detected with the VLBA ($\Delta\alpha^{*}_{\rm
radio}$, $\Delta\delta_{\rm radio}$, and $\Delta\theta_{\rm radio}$). The
offsets $\Delta\alpha^{*}$, $\Delta\delta$, and $\Delta\theta$ are the angular
separations between the radio source (VLBA) and the Gaia DR3 predicted
position at the observed epoch in right ascension, declination, and total,
respectively. The column $\sigma_{\rm sys}$ is the position systematic error
expected from the phase difference of the target and its phase calibrator,
given their angular separation. Finally, column $\Delta\lambda$, which is
derived from $\Delta\theta$, is the linear separation in astronomical units at
the corresponding distance (see Table LABEL:tab:tar).

##  4 Discussion

###  4.1 Comparison with VLA results

The radio sources associated with our target stars are listed in Table
LABEL:tab:tar and were first reported by Launhardt et al. (2022). These
authors measured the variability of the radio sources and pointed out that
they are nonthermal emitter candidates, motivating the follow-up VLBA
observations presented here. We confirm the nonthermal nature of the radio
sources detected with the VLBA, as they require brightness temperatures
$>10^{6}\,$ K to be detected with interferometric baselines of $\sim$8 000 km
(see chapter 9 of Thompson et al., 2017). These brightness temperatures are a
strong indication of the nonthermal radio emission associated with young stars
(Feigelson & Montmerle, 1999).

The angular resolution of VLA images presented by Launhardt et al. (2022) is
$\sim 2^{\prime\prime}$, which is insufficient to resolve potential tight
binaries. As an example, in the case of EC 95 in Serpens, VLA observations
showed the presence of a single radio source (Eiroa et al., 2005). Follow-up
VLBA observations showed that the radio source is a hierarchical triple
system, with three nonthermal variable radio emitters unresolved at the VLA
angular resolution. The first VLBA observations of EC 95 showed a close binary
with an angular separation of about 15 mas (Dzib et al., 2010), and later a
third radio source was detected at an angular distance of about 140 mas from
the close binary (Ortiz-León et al., 2017a). Because of their variability, not
all components were detected in all observed epochs. Thus, whether the radio
sources detected in this work are identical to the radio sources reported by
Launhardt et al. (2022) needs to be analyzed. Indeed, the analysis by
Launhardt et al. (2022) showed that, in some cases, the radio emission can
have contributions from both components in known binary systems or be related
to an unknown companion.

For our analysis, we compared the positions derived from the VLBA and the VLA.
First, we assumed that the radio sources share the trigonometric parallax and
proper motions of the associated stars reported by Gaia (see Table
LABEL:tab:tar) to extrapolate the positions reported from the VLA observations
to the epoch of the corresponding first VLBA observation. Then, we computed
the offset between the position determined by the VLBA observations and the
extrapolated position from VLA observations ($\Delta\alpha_{\rm radio}$,
$\Delta\delta_{\rm radio}$ and $\Delta\theta_{\rm radio}$). The position
offsets of these radio observations are listed in columns (7) to (9) of Table
LABEL:tab:am and are shown graphically in Figures 2 and 2. Given the
differences in angular resolution, the dominant error in the position offsets
is the VLA position error, which is two orders of magnitude larger than those
from the VLBA.

The offsets between the VLA and VLBA positions are less than three times the
errors for nine out of the ten detected sources. In these cases, the VLA and
VLBA radio sources presumably correspond to the same emitting source. In the
case of HD 199143, the separation is larger than three times the errors. This
indicates that the radio source detected with the VLBA corresponds to a
different object than the one reported from VLA observations. Launhardt et al.
(2022) argued that the radio source detected in the VLA observations is
positionally consistent with the expected position of the fainter companion,
HD 199143 B (not in the Gaia DR3 catalog), located $0\hbox
to0.0pt{.\hss}^{\prime\prime}84$ northwest of the main stellar component
(Hagelberg et al., 2020, see also Fig. 2). In the next section, we show that
the radio source detected in the VLBA images corresponds to the primary
stellar component in this system or to a very nearby companion.

Figure 2: Positions of the radio and optical sources related to HD 199143. The
red star indicates the position of the radio source in the first detected
epoch by the VLBA observations. The Cyan dashed line indicates the optical
trajectory of the primary component of the HD 199143 stellar system from
2016.0 to the epoch of detection of the radio source with the VLBA. The
optical position of the star at the VLBA observed epoch is indicated with the
cyan square. The blue square indicates the position of the secondary, HD
199143 B, whose absolute coordinates are unknown. Coordinates relative to the
primary from the WDS catalog were used to estimate the absolute position. Gaia
DR3 and VLBA position errors are smaller than the symbol sizes. The black line
is the trajectory followed by the radio source detected with the VLA
(Launhardt et al., 2022), assuming the derived astrometry at optical
wavelengths from the VLA observed epoch to the VLBA observed epoch. The black
square indicates the extrapolated position of the VLA radio source at the
epoch of the VLBA detection. The black cross size indicates the VLA positional
error.

###  4.2 Comparison with optical astrometry

Table 5: Previous companion information.

| Companion | Angular | $\Delta\theta$ |  PMaa | RUWE |  SORSb | SORS
---|---|---|---|---|---|---|---
Star ID | Name | separation | (mas) | (SNR) |  |  | This work
V 875 Per |  WDS 02523+3617 Bc | $5\hbox to0.0pt{.\hss}^{\prime\prime}4$ | $0.6\pm 0.2$ | … | 1.2 | Main | Main
HD 22213 | WDS 03343–1204 B | $1\hbox to0.0pt{.\hss}^{\prime\prime}7$ | $1.1\pm 0.3$ | … | 1.1 | Main | Main
HD 285281 | WDS 04005+1935 B | $0\hbox to0.0pt{.\hss}^{\prime\prime}8$ | $0.9\pm 0.5$ | … | 1.5 | Main | Main
AI Lep | WDS 05403–1940 B | $8\hbox to0.0pt{.\hss}^{\prime\prime}4$ | $0.5\pm 0.2$ | … | 1.0 | Main | Main
HD 62237 | … | … | $26.6\pm 0.6$ | … | 7.3 | Main | New companion
SAO 135659 | WDS 08138–0738 B | $0\hbox to0.0pt{.\hss}^{\prime\prime}1$ | $73.3\pm 1.0$ | … | 13.8 | Main/Companion | Companion
HD 82159 | WDS 09306+1036 B | $13\hbox to0.0pt{.\hss}^{\prime\prime}8$ | $1.3\pm 0.3$ | 1.74 | 1.6 | Main | Main
HD 82558 | … | … | $0.7\pm 0.2$ | 0.79 | 1.1 | Main | Main
HD 199143 | HD 199143 B | $0\hbox to0.0pt{.\hss}^{\prime\prime}84$ | $4.4\pm 0.5$ | 56.68 | 1.0 | Companion | New Companion
GJ 4199 | WDS 21310+2320 B | $9\hbox to0.0pt{.\hss}^{\prime\prime}2$ | $1.4\pm 0.3$ | 0.99 | 1.0 | Main | Main

666Angular separations of previously known companions are from the Washington
Visual Double Star Catalog (Mason et al., 2001, 2022), except for HD 199143,
which is from (Hagelberg et al., 2020). The quoted error value in
$\Delta\theta$ was estimated by adding in quadrature the statistic and
systematic errors listed in Table LABEL:tab:am.
Sources discussed further in the text are highlighted.
$${\rma}$$$${\rma}$$footnotetext: As reported by Kervella et al. (2019) or
Kervella et al. (2022). $${\rmb}$$$${\rmb}$$footnotetext: Suggested origin of
the radio source (SORS) as reported by Launhardt et al. (2022).
$${\rmc}$$$${\rmc}$$footnotetext: No physical companion (Mason et al., 2022).

To examine the relationship between the radio sources detected with the VLBA
and the targeted stars, we have used the astrometric results from optical
wavelengths (i.e., Gaia) to determine the stellar positions at the epoch of
the VLBA observations. Then, we determined the angular offsets between the
radio source and the optical counterpart ($\Delta={\rm Pos}_{\rm radio}-{\rm
Pos}_{\rm Opt.}$). The initial parameter uncertainties (optical: position,
proper motion, and parallax; radio: position) were propagated to estimate the
angular separation error by using a Monte Carlo approach. In columns 10, 11,
and 12 of Table LABEL:tab:am, we list the angular separation in right
ascension, $\Delta\alpha*=\Delta\alpha\cdot\cos{({\rm Dec.})}$ and in
declination, ($\Delta\delta$), and the total angular separation
($\Delta\theta$) between the VLBA and optical sources. Expected systematic
errors in the source position caused by residual phase calibration errors are
listed in column 11 and were added quadratically to the errors provided by
JMFIT.

Seven of the ten radio sources are consistent with the expected position from
Gaia results within five times the errors. We argue that, in these cases, the
radio emission is tracing the stellar corona of the stars seen at optical
wavelengths. Stellar coronae of young stars could be of sizes up to 10 R∗ (see
the review by Güdel, 2002). For a conservative value for the stellar radii, we
assumed 1 R∗= 1.15 R⊙, as the observed stars are of spectral types F8 or later
and about to enter the main sequence. At the distances of our targets, this
implies angular sizes of less than $2.9$ mas. On the other hand, the measured
positions of the other three radio sources (HD 62237, SAO 135659, and HD
199143) are not consistent with the results from Gaia, and the offset cannot
be attributed to the stellar coronae structure, as they would require
structures greater than $40$ R∗.

To look for clues on the origin of the radio-optical position discrepancies in
these three sources, we first turned to the Gaia results. The Gaia catalog
includes a measure of the quality of the astrometric fits called the
renormalized unit weight error (RUWE; listed Table LABEL:tab:bin for our
targets). The RUWE index is expected to be close to 1.0 for stars whose motion
is well described by a uniform linear motion, while a large RUWE value
($>1.4$) indicates a significant deviation from a linear uniform motion and is
suggestive of the presence of an unseen companion. Interestingly, two of the
three sources where we found discrepancies between the VLBA and optical
positions (HD 62237 and SAO 135659) have very large RUWE values (7.3 and 13.8,
respectively). A second relevant piece of information is provided by the
proper motion anomaly (PMa) introduced by Kervella et al. (2019) and Kervella
et al. (2022) and listed in the fifth column of Table LABEL:tab:bin. The PMa
quantifies discrepancies between the proper motions measured by the Hipparcos
and Gaia astrometric missions. A large PMa indicates significantly different
proper motions in the Hipparcos and Gaia catalogs and suggests the presence of
a low-mass companion with an orbital period larger than the Gaia observation
window (668 days for Gaia DR3). The third source where we found a discrepancy
between the VLBA and optical positions (HD 199143) has a very large PMa
(56.7).

In summary, the three sources where the VLBA and optical positions do not
coincide have astrometric evidence of multiplicity. In contrast, the sources
where the VLBA and optical positions do coincide have RUWE values close to
unity and either small PMa values or no PMa at all.

Given the evidence presented above that the sources with discrepant VLBA and
optical positions are multiple, our next step was to look for known
companions. We searched the Washington Visual Double Star Catalog (WDS; Mason
et al., 2001, 2022) for companions to our target stars.777For completeness, we
searched the WDS catalog for companions to all of our targets, not just for HD
62237, SAO 135659, and HD 199143. We list the stellar companions and their
angular separation in Table LABEL:tab:bin. In cases where there is more than
one known companion, we include only the companion with the smallest angular
separation to the targeted star. For HD 62237 and HD 199143, the angular
separation of the visual binaries is much larger than the VLBA-optical
astrometric discrepancy $\Delta\theta$. In the case of SAO 135659, however,
the values are comparable. We come back to this point momentarily. We also
investigated if our targets had been reported as spectroscopic binaries. Four
of our target sources (HD 22213, SAO 135659, GJ 4199, and HD 199143) were
included in the spectroscopic analysis by Zúñiga-Fernández et al. (2021). Only
HD 22213 was confirmed as a spectroscopic binary. In particular, none of the
stars with discrepant optical and VLBA positions are a known spectroscopic
binary.

Armed with the information gathered for each system, we are now in a position
to discuss the three targets where the VLBA and optical positions do not
coincide. Particularly, we discuss if the origin of the position discrepancy
can be explained by known companions.

Firstly, SAO 135659 is known to be a binary star with stellar components of
similar masses (0.77 and 0.64 M⊙ for the primary and secondary, respectively).
The angular separation of the system is $\sim 0\hbox
to0.0pt{.\hss}^{\prime\prime}1$ almost in the north-south direction at the
epoch of the observations reported by Elliott et al. (2015). At a distance of
50 pc and given the total mass, we can estimate the period of the system to be
roughly 9 years. Coincidentally, this is similar to the separation between the
observations realized by Elliott et al. (2015) in May 2012 and our VLBA
observations in April 2021. Consequently, we expect the relative offset
between the two stars to also be on order of $\sim 0\hbox
to0.0pt{.\hss}^{\prime\prime}1$ in the north-south direction at the time of
our VLBA observations. In our observations, the position of the VLBA source
was found to be located about $\sim 0\hbox to0.0pt{.\hss}^{\prime\prime}07$
north of the expected Gaia position (Fig. 2). The similarity between the
relative position of the optical and VLBA sources and the expected relative
positions of the two stars in SAO 135659 strongly suggests that the VLBA
source is associated with the secondary (WDS 08138-0738). The small difference
in separation ($\sim 0\hbox to0.0pt{.\hss}^{\prime\prime}1$ vs. $\sim 0\hbox
to0.0pt{.\hss}^{\prime\prime}07$) could easily be caused by the somewhat
different orbit phase seen during our observations compared to those of
Elliott et al. (2015) and by the fact that Gaia presumably traces the position
of the photocenter of the system rather than the exact position of the
primary.

The source HD 199143 exhibits a large PMa. The detailed analysis from
Launhardt et al. (2022) shows that this can be explained by the known
companion HD 199143 B located at $\sim 0\hbox to0.0pt{.\hss}^{\prime\prime}84$
(see Fig. 12 from Launhardt et al., 2022). The VLA source reported by
Launhardt et al. (2022) coincides with that companion (see also Fig.2). The
VLBA source we detect here, however, is located very near the primary of the
system (Fig.2). The small but significant offset between the VLBA position and
the expected position of the primary (4.4 $\pm$ 0.5 mas) could be due to the
presence of another hitherto undetected companion located very close to the
primary. Given the small offset, however, further observations will be needed.
To finish on this source, we note that we searched for a second VLBA radio
source in the expected position of the companion but found none. The flux of
the companion star was below 110 $\mu$Jy beam-1(6.5 times the noise level in
our image) during the VLBA observations. The radio source detection by
Launhardt et al. (2022) and the one reported here is interesting but
unsurprising, and it points to the highly variable nature of the radiation
mechanisms in magnetically active stars (e.g., Güdel, 2002).

Regarding HD 62237, it has no known companion, but the large Gaia RUWE value
strongly suggests that such a companion exists. We argue that the VLBA source,
offset by about 27 mas from the expected position of the primary, traces that
companion. We note that, given the modest resolution of the VLA observations
reported by Launhardt et al. (2022), the VLA source could coincide with either
the primary or the secondary putatively traced by the VLBA (Fig.2). This
companion should be less luminous than the primary star, as a clear offset
between the Gaia and VLBA positions indicates that the system photocenter is
coincident or close to the primary star. However, the mass of the companion
must be significant to cause the linear fitting by Gaia DR3 to not properly
describe the motion of the system (RUWE value is 7.3). The spectral type of HD
62237 is G5V, indicating a star with a mass of 0.93 M⊙ (Zombeck, 1990).
Assuming that the companion has a mass between 0 and 0.5 M⊙, the mass of the
system ranges from 0.93 to 1.43 M⊙. We also assume that the optical-radio
offset of $26.6\pm 0.6$ mas ($3.32\pm 0.08$ au at a distance of 124.76 pc) is
the semi-major axis of the orbit, and we roughly estimate an orbital period of
the system in the range from 5.1 to 6.3 years.

##  5 Conclusions

We have presented a series of VLBA observations of a sample of 31 nearby young
stars previously reported as radio sources, obtaining ten detections (for a
detection rate on the order of 30%). The radio sources are compact at our
angular resolution of a few milli-arcseconds and are highly variable; both
properties confirm they are nonthermal radio emitters. All detected radio
sources are associated with stars with good astrometry from Gaia DR3. Using
these results, we extrapolated the position of the optical stars to the epoch
of the radio observations to compare the radio and optical positions. For
seven of the ten detected radio sources, the optical and radio positions are
consistent. Three radio sources, however, have position discrepancies above
eight times the errors. All three of these sources show independent evidence
of multiplicity, and we attribute the position discrepancies to the presence
of companions.

Follow-up observations of the detected radio sources will be of high interest
for comparison with Gaia astrometry. Such observations could provide some
clues on the Gaia parallax zero point and on the relative orientation of the
Gaia and ICRF reference frames by increasing the current numbers of stars with
Gaia- and VLBI-derived astrometry.

###### Acknowledgements.

We thank the anonymous referee for his/her careful review and insightful
comments. We thank Eduardo Ros for reading the paper and providing suggestions
that improve the readability of the paper. S.A.D. acknowledge the M2FINDERS
project from the European Research Council (ERC) under the European Union’s
Horizon 2020 research and innovation programme (grant No 101018682). L.L.
acknowledges the support of UNAM-DGAPA PAPIIT grants IN108324 and IN112820 and
CONACYT-CF grant 263356. J.O.-T. acknowledges a stipend from CONAHCYT, México.
The National Radio Astronomy Observatory is a facility of the National Science
Foundation operated under cooperative agreement by Associated Universities,
Inc. This work has made use of data from the European Space Agency (ESA)
mission Gaia (https://www.cosmos.esa.int/gaia), processed by the Gaia Data
Processing and Analysis Consortium (DPAC,
https://www.cosmos.esa.int/web/gaia/dpac/consortium). Funding for the DPAC has
been provided by national institutions, in particular the institutions
participating in the Gaia Multilateral Agreement.

## References

  * Beasley & Conway (1995) Beasley, A. J. & Conway, J. E. 1995, in Astronomical Society of the Pacific Conference Series, Vol. 82, Very Long Baseline Interferometry and the VLBA, ed. J. A. Zensus, P. J. Diamond, & P. J. Napier, 327
  * Bell et al. (2015) Bell, C. P. M., Mamajek, E. E., & Naylor, T. 2015, MNRAS, 454, 593
  * Briggs (1995) Briggs, D. S. 1995, PhD thesis, New Mexico Institute of Mining and Technology, United States
  * Carpenter et al. (2009) Carpenter, J. M., Bouwman, J., Mamajek, E. E., et al. 2009, ApJS, 181, 197
  * da Silva et al. (2009) da Silva, L., Torres, C. A. O., de La Reza, R., et al. 2009, A&A, 508, 833
  * Desidera et al. (2015) Desidera, S., Covino, E., Messina, S., et al. 2015, A&A, 573, A126
  * Dzib et al. (2010) Dzib, S., Loinard, L., Mioduszewski, A. J., et al. 2010, ApJ, 718, 610
  * Dzib et al. (2021) Dzib, S. A., Forbrich, J., Reid, M. J., & Menten, K. M. 2021, ApJ, 906, 24
  * Eiroa et al. (2005) Eiroa, C., Torrelles, J. M., Curiel, S., & Djupvik, A. A. 2005, AJ, 130, 643
  * Elliott et al. (2015) Elliott, P., Huélamo, N., Bouy, H., et al. 2015, A&A, 580, A88
  * Feigelson & Montmerle (1999) Feigelson, E. D. & Montmerle, T. 1999, ARA&A, 37, 363
  * Forbrich et al. (2021) Forbrich, J., Dzib, S. A., Reid, M. J., & Menten, K. M. 2021, ApJ, 906, 23
  * Gaia Collaboration et al. (2021) Gaia Collaboration, Brown, A. G. A., Vallenari, A., et al. 2021, A&A, 649, A1
  * Gaia Collaboration et al. (2016) Gaia Collaboration, Prusti, T., de Bruijne, J. H. J., et al. 2016, A&A, 595, A1
  * Gaia Collaboration et al. (2023) Gaia Collaboration, Vallenari, A., Brown, A. G. A., et al. 2023, A&A, 674, A1
  * Galicher et al. (2016) Galicher, R., Marois, C., Macintosh, B., et al. 2016, A&A, 594, A63
  * Galli et al. (2019) Galli, P. A. B., Loinard, L., Bouy, H., et al. 2019, A&A, 630, A137
  * Greisen (2003) Greisen, E. W. 2003, in Astrophysics and Space Science Library, Vol. 285, Information Handling in Astronomy - Historical Vistas, ed. A. Heck, 109
  * Güdel (2002) Güdel, M. 2002, ARA&A, 40, 217
  * Hagelberg et al. (2020) Hagelberg, J., Engler, N., Fontanive, C., et al. 2020, A&A, 643, A98
  * Høg et al. (2000) Høg, E., Fabricius, C., Makarov, V. V., et al. 2000, A&A, 355, L27
  * Kenyon & Hartmann (1995) Kenyon, S. J. & Hartmann, L. 1995, ApJS, 101, 117
  * Kervella et al. (2019) Kervella, P., Arenou, F., Mignard, F., & Thévenin, F. 2019, A&A, 623, A72
  * Kervella et al. (2022) Kervella, P., Arenou, F., & Thévenin, F. 2022, A&A, 657, A7
  * Launhardt et al. (2022) Launhardt, R., Loinard, L., Dzib, S. A., et al. 2022, ApJ, 931, 43
  * Lestrade et al. (1999) Lestrade, J. F., Preston, R. A., Jones, D. L., et al. 1999, A&A, 344, 1014
  * Lestrade et al. (1990) Lestrade, J. F., Rogers, A. E. E., Whitney, A. R., et al. 1990, AJ, 99, 1663
  * Lindegren (2020) Lindegren, L. 2020, A&A, 633, A1
  * Lindegren et al. (2018) Lindegren, L., Hernández, J., Bombrun, A., et al. 2018, A&A, 616, A2
  * Loinard et al. (2005) Loinard, L., Mioduszewski, A. J., Rodríguez, L. F., et al. 2005, ApJ, 619, L179
  * Loinard et al. (2007) Loinard, L., Torres, R. M., Mioduszewski, A. J., et al. 2007, ApJ, 671, 546
  * Mason et al. (2001) Mason, B. D., Wycoff, G. L., Hartkopf, W. I., Douglass, G. G., & Worley, C. E. 2001, AJ, 122, 3466
  * Mason et al. (2022) Mason, B. D., Wycoff, G. L., Hartkopf, W. I., Douglass, G. G., & Worley, C. E. 2022, VizieR Online Data Catalog, B/wds
  * Montes et al. (2001) Montes, D., López-Santiago, J., Gálvez, M. C., et al. 2001, MNRAS, 328, 45
  * Ortiz-León et al. (2017a) Ortiz-León, G. N., Dzib, S. A., Kounkel, M. A., et al. 2017a, ApJ, 834, 143
  * Ortiz-León et al. (2018) Ortiz-León, G. N., Loinard, L., Dzib, S. A., et al. 2018, ApJ, 869, L33
  * Ortiz-León et al. (2017b) Ortiz-León, G. N., Loinard, L., Kounkel, M. A., et al. 2017b, ApJ, 834, 141
  * Reid & Brunthaler (2004) Reid, M. J. & Brunthaler, A. 2004, ApJ, 616, 872
  * Reid et al. (2017) Reid, M. J., Brunthaler, A., Menten, K. M., et al. 2017, AJ, 154, 63
  * Reid & Honma (2014) Reid, M. J. & Honma, M. 2014, ARA&A, 52, 339
  * Riess et al. (2018) Riess, A. G., Casertano, S., Yuan, W., et al. 2018, ApJ, 861, 126
  * Stanford-Moore et al. (2020) Stanford-Moore, S. A., Nielsen, E. L., De Rosa, R. J., Macintosh, B., & Czekala, I. 2020, ApJ, 898, 27
  * Thompson et al. (2017) Thompson, A. R., Moran, J. M., & Swenson, George W., J. 2017, Interferometry and Synthesis in Radio Astronomy, 3rd Edition
  * van Leeuwen (2007) van Leeuwen, F. 2007, A&A, 474, 653
  * Weise et al. (2010) Weise, P., Launhardt, R., Setiawan, J., & Henning, T. 2010, A&A, 517, A88
  * Xu et al. (2019) Xu, S., Zhang, B., Reid, M. J., Zheng, X., & Wang, G. 2019, ApJ, 875, 114
  * Zinn et al. (2019) Zinn, J. C., Pinsonneault, M. H., Huber, D., & Stello, D. 2019, ApJ, 878, 136
  * Zombeck (1990) Zombeck, M. V. 1990, Handbook of space astronomy and astrophysics
  * Zúñiga-Fernández et al. (2021) Zúñiga-Fernández, S., Bayo, A., Elliott, P., et al. 2021, A&A, 645, A30

##  Appendix A Images of detected radio sources

We detected ten of our 31 targets. The obtained images from our VLBA
observations are presented in this appendix. The plots in Figure 3 are the
intensity maps of nine detected radio sources; the remaining source is
presented in the main text. The figures include the expected positions of the
stars at the epochs of the VLBA observations as calculated from Gaia DR3 (Gaia
Collaboration et al. 2021, 2023).

To produce the images combining the data from both epochs, we corrected the
position of the second epoch, assuming that the optical astrometric parameters
could be applied to the radio observations. In AIPS, this was done with the
task CLCOR using the operation code ANTC. The position correction for each
target was performed by specifying the offsets in right ascension and
declination in arcsecond units (clcorprm(5,6) =$\alpha_{\rm offset}$,
$\delta_{\rm offset}$). The values of these offsets are listed in Table
LABEL:tab:off.

We assumed the upper limits for the nondetected sources to be 6.5 times the
noise levels of images where the two observed epochs of each source were
combined. These noise level values are also listed in Table LABEL:tab:off.

Table 6: Offsets applied to combine data of the two observed epochs and the
noise level of the obtained image.

| Offsets (mas) | Noise
---|---|---
Star ID | $\alpha$ | $\delta$ |  $\mu$Jy bm-1
BD17232 | $-0.45$ | $+0.13$ | 25
HIP 12545 | $-0.30$ | $+0.22$ | 13
HIP 12635 | $-0.41$ | $+0.89$ | 14
V 875 Per | $-0.32$ | $+0.27$ | 23
TYC 3301 | $-0.37$ | $+0.79$ | 15
HD 22213 | $-0.55$ | $-0.23$ | 17
HD 23524 | $-0.44$ | $+0.76$ | 19
HD 24681 | $-0.21$ | $+0.43$ | 12
HD 285281 | $+0.02$ | $+0.07$ | 13
HD 284135 | $+0.05$ | $+0.11$ | 14
HD 31281 | $-0.00$ | $+0.03$ | 21
BD 08995 | $-0.08$ | $-0.05$ | 16
HD 286264 | $-0.06$ | $+0.12$ | 16
HD 293857 | $-0.05$ | $+0.06$ | 15
TYC 7136611 | $+0.01$ | $+0.25$ | 10
TYC 59251547 | $-0.04$ | $-0.03$ | 12
AI Lep | $-0.04$ | $-0.13$ | 12
HD 62237 | $+0.05$ | $-0.10$ | 15
SAO 135659 | $+0.26$ | $-0.15$ | 16
WDS 09035+3750B | $+0.42$ | $+4.74$ | 19
HD 82159 | $+3.42$ | $+0.14$ | 38
HD 82558 | $+4.01$ | $-3.24$ | 22
GJ 2079 | $+2.75$ | $+4.44$ | 16
HD 135363 | $+1.02$ | $-1.03$ | 20
UCAC 4832 | $+1.02$ | $-1.03$ | 14
HD 199143 | $-0.37$ | $+0.51$ | 17
HD 358623 | $-0.38$ | $+0.54$ | 16
SAO 50350 | $-0.35$ | $-0.93$ | 14
GJ 4199 | $-0.28$ | $-0.00$ | 14
SAO 51891 | $-0.43$ | $-0.45$ | 18
SAO 108142 | $-0.29$ | $+0.05$ | 13
888Offsets are as expected assuming the parallax and proper motions obtained
from the optical results. Detected sources are highlighted with gray-shaded
rows. The noise value was measured in the image of two combined epochs, that
is, after applying the correction of the offsets.

Figure 3: VLBA images of detected radio sources. The images are centered in
the position of the radio source as detected in the first epoch (see Table
LABEL:tab:am). The name of the stellar source related to the radio sources is
indicated in the top-left corner. The images, from left to right, correspond
to the epochs 1, 2, and the combination of both epochs. Contour levels are -3,
3, 6, 9, 12, and 15 times the noise level of the image as listed in Table 3.
The yellow open ellipse in the bottom-left corner represents the size of the
synthesized beam of the image as listed in Table 3. The predicted optical
position in epochs 1 and 2 are shown as red and cyan crosses, respectively.
For HD 62237 and SAO 135659, these positions fall outside the region shown.

Figure 1: Continued. Figure 1: Continued.

##  Appendix B Figures comparing VLBA, VLA, and Gaia astrometric results

Figure 2 shows the measured VLBA position of detected radio sources. In these
plots, for comparison, we also show the position of the VLA radio sources
detected by Launhardt et al. (2022), where we have extrapolated its position
to the epoch of the VLBA observation using the results of optical astrometry.
The expected optical position at the epoch of VLBA observations is also shown.
Finally, in two cases, SAO 135659 and HD199143, the relative position to the
primary of a companion was obtained from the Washington Visual Double Star
catalog (Mason et al. 2001, 2022) and are also shown in these plots.

|  |
---|---|---
|  |
|  |
Figure 2: Positions of radio and optical sources related to our target
sources. The name of source is indicated at the top of each plot. Red stars
indicate the position of the radio source in the first detected epoch by the
VLBA observations. The black line is the trajectory followed by the radio
source detected with the VLA (Launhardt et al. 2022), assuming the derived
astrometry at optical wavelengths (see Table LABEL:tab:am) from the VLA
observed epoch to the VLBA observed epoch. The black square indicates the
extrapolated position of the VLA radio source at the epoch of the VLBA
detection. The black cross size indicates the VLA positional error. The cyan
line indicates the trajectory of the optical results from 2016.0 to the epoch
of detection of the radio source with the VLBA; the position of the optical
source at this epoch is indicated with the cyan square. Gaia DR3 and VLBA
position errors are smaller than the symbol sizes. Companion relative
positions are shown as blue squares.