JuanCastillo29/space-mission-intelligence-data huggingface dataset:
The evolution of exocomets and their source populations
Alexander J. Mustill1*, Tim Pearce2†, Michele Bannister3†, Susanne Pfalzner4†, Dag Evensberget5,19, Dimitri Veras2,6,7, Rosita Kokotanekova8,9, Matthew Hopkins3,10, Dennis Bodewits11, Darryl Z. Seligman12,13, Isabel Rebollido14, Raphael Marschall15,16, Bin Yang17, Klaudia Jaworska1, Xabier P´erez Couto18
arXiv:2604.08190v1  [astro-ph.EP]  9 Apr 2026
1*Department of Physics, Lund University, Box 118, Lund, 22100, Sweden. 2Department of Physics, University of Warwick, Gibbet Hill Road, Coventry, CV4 7AL, United Kingdom. 3School of Physical and Chemical Sciences – Te Kura Mat¯u, University of Canterbury, Private Bag 4800, Christchurch, 8140, Aotearoa New Zealand. 4J¨ulich Supercomputing Centre, Research Centre J¨ulich, Wilhelm-Johnen-Strasse, J¨ulich, 52428, Germany. 5Leiden Observatory, Leiden University, PO Box 9513, Leiden, 2300 RA, The Netherlands. 6Centre for Exoplanets and Habitability, University of Warwick, Gibbet Hill Road, Coventry, CV4 7AL, United Kingdom. 7Centre for Space Domain Awareness, University of Warwick, Gibbet Hill Road, Coventry, CV4 7AL, United Kingdom. 8Institute of Astronomy and NAO, Bulgarian Academy of Sciences, 72 Tsarigradsko Shose Blvd., Sofia, 1784, Bulgaria. 9International Space Science Institute, Hallerstrasse 6, Bern, 3012, Switzerland. 10Department of Physics, University of Oxford, Keble Road, Oxford, OX1 3RH, United Kingdom. 11Department of Physics, Auburn University, Edmund C. Leach Science Center, Auburn, 36849, AL USA. 12Department of Physics and Astronomy, Michigan State University, East Lansing, MI 48824, USA. 13NSF Astronomy and Astrophysics Postdoctoral Fellow.

14European Space Astronomy Centre (ESAC), European Space Agency, Camino Bajo del Castillo s/n, Villanueva de la Ca˜nada, 28692, Madrid Spain. 15Laboratoire J.-L. Lagrange, Observatoire de la Cˆote d’Azur, CNRS, CS 34229, Nice Cedex 4, 06304, France. 16Physics Institute, Space Research and Planetary Sciences, University of Bern, Sidlerstrasse 5, Bern, CH-3012, Switzerland. 17Instituto de Estudios Astrof´ısicos, Facultad de Ingenier´ıa y Ciencias, Universidad Diego Portales, Av. Ej´ercito Libertador 441, Santiago, 8370191, Chile. 18Centre for Research in Information and Communication Technologies, Universidade da Coru˜na, Campus de Elvi˜na s/n, A Coru˜na, 15071, Galiza Spain. 19Centre for Planetary Habitability (PHAB), Department of Geosciences, University of Oslo, Sem Sælands vei 2A, Oslo, 0371, Norway.
*Corresponding author(s). E-mail(s): alexander.mustill@fysik.lu.se; Contributing authors: tim.pearce@warwick.ac.uk; michele.bannister@canterbury.ac.nz; s.pfalzner@fz-juelich.de;
dag.evensberget@geo.uio.no; dimitri.veras@aya.yale.edu; rkokotanekova@nao-rozhen.org; matthewj.hopkins@canterbury.ac.nz;
dennis@auburn.edu; dzs@msu.edu; isabel.rebollidovazquez@esa.int;
raphael.marschall@unibe.ch; bin.yang@mail.udp.cl; kl8058ja-s@student.lu.se; xabier.perez.couto@udc.gal; †These authors contributed equally to this work.
(Accepted 8 May 2026) Submitted to ApJ ABSTRACT Nearby short-period exoplanet systems may produce detectable stellar radio emission due to sub-Alfv´enic star-planet interaction (SPI), but there are no confirmed cases yet. We targeted five slowly-rotating M dwarfs with transiting terrestrial planets, observing at GHz frequencies throughout their sub-day orbital periods. We did not detect any bursty SPI-like emission, but detected two stars in quiescence: LHS 3844 (unpolarized) and LHS 1678 (circularly polarized). These detections imply persistent magnetic activity at Gyr ages, especially notable for LHS 1678 given its low photometric variability, and can serve as targets for radio transit experiments. Our SPI non-detections may be due to radio beaming geometry, a sub-GHz maximum emission frequency, or undetectable flux density. If the last case applies, then flux density upper limits constrain the exoplanet magnetosphere. GJ 367 b has the tightest constraints – no extended magnetosphere and an exoplanet field <0.8 G – although these results depend strongly Corresponding author: Jackie Villadsen j.villadsen@bucknell.edu 2 Villadsen et al.
terrestrial planet hosts with literature radio observations, and a hypothetical reference system with easily scalable basic parameters. Our 5 target systems were all discovered via the transit method and thus have radii available, whereas for RV-discovered comparison systems YZ Cet b, Prox Cen b, and GJ 3323 b we assume Earth density and edge-on orientation: Rp = (Mp sin i)1/3 in Earth units. The nearness to Earth and short orbits of our targets make them interesting follow-up targets even though they are too hot to be habitable, for instance LHS 3844 b has potential for JWST

4 Villadsen et al.

[Table: Table 1. Target Systems]
Exoplanet b
Radio Observations

SpT M∗ R∗
d Prot Lbol Rp Porb
a Ref Fν
Freq. Ref

| LTT 3780 LHS 3844 GJ 367 LHS 1678 GJ 1252 | M3.5 0.379 0.382 22.03 137 1.67×10−2 M4.5 0.151 0.189 14.878 128 2.72×10−3 M1 0.454 0.457 9.419 48 2.88×10−2 M2 0.345 0.329 19.865 64 1.45×10−2 M3 0.381 0.391 20.385 64 1.96×10−2 | 1.35 0.768 6.69 1 1.303 0.463 7.09 2 0.718 0.322 3.32 3 0.7 0.860 8.11 4 1.193 0.518 5.03 5 | <75 2-4 * <350 1.1-3.1 * <330 1.1-3.1 * <610 1.1-3.1 * <460 1.1-3.1 * |
| --- | --- | --- | --- |
| YZ Cet GJ 486 Prox Cen GJ 3323 | M4.5 0.137 0.163 3.717 63.5 2.25×10−3 M3.5 0.333 0.339 8.079 49.9 1.21×10−2 M5.5 0.12 0.141 1.301 89 1.55×10−3 M4.5 0.164 0.119 5.375 99 2.70×10−3 | 0.89 2.02 21.3 6,A,† 1.305 1.47 11.1 7 1.1 11.2 73.7 8,† 1.26 5.36 59.3 9,† | 470-1070 2-4 A <216 0.55-0.75 B 200-5000 1.1-3.1 C 86±10 4-8 D |
| Reference | M2.5 0.3 0.3 10 100 1.00×10−2 | 1 1 9.4 10 | <100 0-2 |

characterization (Whittaker et al. 2022). However, at a reduced semi-major axis of a/R∗∼3-8, our targets should be tidally locked in circularized orbits, with stellar irradiation sufficient to have evaporated any Earth-like atmospheres (e.g., GJ 367; Poppenhaeger et al. 2024). Tidal locking, circularization, and lack of atmosphere are supported by IR secondary eclipse observations of GJ 1252 (Crossfield et al. 2022) and LTT 3780 (Allen et al. 2025), and IR orbital phase curves for GJ 367 b (Zhang et al. 2024) and LHS 3844 b (Kreidberg et al. 2019; Lyu et al. 2024), with the latter’s atmosphere also undetected in ground-based transit spectroscopy (Diamond-Lowe et al. 2020). The planets likely have volatile-poor, solid surfaces (e.g., Kane et al. 2020), but alternative proposed interpretations revive the possibility of an atmosphere, including explaining phase curves with a thick atmosphere with night-side clouds (Powell et al. 2024) and a novel form of tectonics driven by the dayside/nightside temperature contrast (Meier et al. 2021). Overall, the preponderance of evidence suggests that our target planets are unlikely to have an ionosphere that can form an induced magnetosphere similar to Io’s. However, if they have a strong extended magnetosphere above the planet’s surface similar to Ganymede, then strong SPI interaction can occur without an ionosphere (Saur et al. 2013). All of our target stars are slowly rotating early-to-mid M dwarfs with weak stellar magnetic activity and low flare rates. Slow rotators minimize the risk of contamination by non-SPI flaring (Pope et al.
39.76
69.16
39.78
J2000 Declination
J2000 Declination
Jy / Beam
Jy / Beam
69.18
39.80
69.20
39.82
69.22
39.84
340.46 340.48 340.50 340.52 340.54
68.14 68.16 68.18 68.20 68.22
J2000 Right Ascension
J2000 Right Ascension

[Figure: Figure 1. (left) ATCA 1.1-3.1 GHz Stokes I image of LHS 3844 during its first epoch. The star is detected in Stokes I at 63 ± 12 µJy, where the uncertainty is estimated from the image RMS in empty regions near the star’s location. The star is undetected in Stokes V. (right) Stokes V image of LHS 1678 during its three epochs combined. The star is detected in Stokes V at -70±9 uJy, and undetected in Stokes I (Appendix A).]
on orbital and magnetic geometry (Kavanagh & Vedantham 2023), an emission cone thickness of
the targets’ short orbital periods.
Throughout the remainder of the paper, we use 3σ upper limits on the flux density based on the
standard deviation of the Stokes V time series with 5-minute time bins (shown in Appendix B).
The 3σ upper limits are summarized in Table 1; for multi-epoch observations, the table presents the
standard deviation of all epochs combined, except for GJ 367 where each epoch covers almost the
full orbital period so we report the most sensitive epoch’s limits. The CABB sensitivity calculator
predicts σ =60 µJy sensitivity in 5 minutes an ATCA 1.1-3.1 GHz with typical weather, 6 antennas,
and RFI removed. Our measured Stokes V sensitivities range from 90-200 µJy, due to a missing
antenna in some epochs plus poor weather, heavy RFI, and/or sidelobes of bright off-axis background
sources leaking into Stokes V.
Two targets were detected in the image plane in images of the full duration of one or more epochs.
LHS 3844 was detected in Stokes I with a flux density of 63 ± 12 µJy in the first of two epochs
(Figure 1) and 56 ± 17 µJy in the second epoch (Appendix A), although there is more visible sidelobe
contamination around the source in the second epoch. LHS 1678 is detected in Stokes V in an image
combining all 3 epochs (Figure 1), at a flux density of -70±9 µJy. While this is a 7σ detection and
there are no known background sources at this location, it has a somewhat non-point-like shape that
may be due to time-variable poor phase stability or polarization calibration, so we consider this a
tentative detection. Imaging single epochs did not improve the signal-to-noise ratio of the LHS 1678
images. The Stokes I flux density at the star’s location (no clear source) is 63 ± 40 µJy, estimated
from the intensity value at the center pixel in the image and the image RMS near that location. The
nearby Stokes I image RMS is large due to a nearby extended source with unclear boundaries, which
is clearly offset from the Stokes V source but whose edges may overlap this location.
almost entirely dipolar (Lehmann et al. 2024) [hereafter L24]. While none of our 5 targets have ZB or ZDI measurements, some of the comparison sample do, but we use the rotation-based magnetic field predictions throughout this work for simplicity.
6.2. Wind base conditions: Mass loss rate None of our targets have wind measurements. To estimate their mass loss rates, we employ the power-law fit of Wood et al. (2021) for a population of main-sequence stars including about half SPI Radio Search 9 [Table:

Table 3. Predicted stellar wind parameters]

| LTT 3780 LHS 3844 GJ 367 LHS 1678 GJ 1252 | 125 7.5 557 78 325 19 437 61 310 19 | 2.14 4.3×1025 * 0.27 0.039 0.99 5.6×1025 * 0.97 0.035 0.93 5.0×1026 1 1.34 0.28 0.91 3.8×1026 * 1.81 0.20 1.01 3.9×1026 * 1.41 0.21 |
| --- | --- | --- |
| YZ Cet GJ 486 Prox Cen GJ 3323 | 1800 252 562 78 1580 221 1940 271 | 0.47 6.4×1026 2 7.95 0.21 0.68 6.3×1025 3 0.43 0.050 0.62 1.7×1027 4 19.4 0.42 0.80 1.9×1027 5 30.1 0.43 |
| Reference | 301 42 | 1.23 1.1×1026 * 0.82 0.073 |
| --- | --- | --- | --- | --- |
|  | GJ GJ 125 3844 S 167 ence |  | et 3 |  |
|  |  | LHS LH Refer | YZ C J 332 Cen |  |
|  |  | 780 | 86 G Prox ury |  |
|  | 3 LTT Wind Model |  | GJ 4 Merc |  |
|  | Uncertainty |  |  |  |

|  | v A |  | v orb v wind |  |
| --- | --- | --- | --- | --- |
|  |  |  | v A |  |
|  | nd |  |  |  |
|  | v w i | rb o v |  |  |

3)
Wind Density (g cm
Speed (km/s)
1.0

|  | If B wind 0 |  |  | .5 G, Earth has |  |
| --- | --- | --- | --- | --- | --- |
|  |  |  | If B wind 0 | .5 G, Earth has |  |
|  | no extend |  |  | ed magnetosphere |  |
|  |  |  |  | B wind = B |  |
|  |  |  |  |  |  |
|  | Wind Model Uncertainty |  |  |  |  |

0.8
Alfven Mach Number MA
Stellar Wind B Field (G)
0.6 0.4
0.2
0.0
Distance from Star (R * )
Distance from Star (R * )

[Figure: Figure 2. Predicted wind parameters for our target systems (the 5 closest to the star) and comparison systems. We show the radial dependence of wind parameters for a hypothetical reference system defined in Table 1 (gray lines). (Top left) Stellar wind density. Later-type M dwarfs have greater predicted mass loss flux and thus higher wind density relative to the reference line. The density of the solar wind at Mercury is shown for comparison. The error bar at left shows the scale of uncertainty for all systems due to inferring wind properties from rotation-activity relations (Section 6.7). (Top right) Wind speed in stellar frame (w), planet orbital speed (o), and Alfv´en speed (A). (Bottom left) Stellar wind magnetic field, modeled as an open radial field. Later-type M dwarfs have stronger predicted surface fields compared to the early-type reference system. The dashed line shows a typical Earth field of 0.5 G; if the wind field exceeds this value, an Earth twin will not have an extended magnetosphere. (Bottom right) Alfv´en Mach number of the wind relative to the planet. All of our targets can have sub-Alfv´enic interactions (MA < 1) between the star and its innermost planet.]
planet orbital speed (o), and Alfv´en speed (A). (Bottom left) Stellar wind magnetic field, modeled as an open radial field. Later-type M dwarfs have stronger predicted surface fields compared to the early-type reference system. The dashed line shows a typical Earth field of 0.5 G; if the wind field exceeds this value, an Earth twin will not have an extended magnetosphere. (Bottom right) Alfv´en Mach number of the wind relative to the planet. All of our targets can have sub-Alfv´enic interactions (MA < 1) between the star and its innermost planet.]

12 Villadsen et al.
6.5. Wind magnitude at planet: Velocity Solar/stellar winds accelerate significantly at the 3-10R∗orbital distances of our target planets. We model this acceleration using a Parker wind (Parker 1958), which has an analytical solution for wind speed (Cranmer 2004) using the Lambert W function, available in scipy (Virtanen et al. 2020). This solution, derived from conservation of mass and momentum, assumes a steady-state isothermal wind with spherical symmetry (i.e., no angular expansion of open magnetic field lines), a better assumption for the equatorial slow solar wind than the polar fast wind. The 4 equations needed to implement this solution are reproduced here for ease of use:
1) Isothermal sound speed assuming the wind consists of fully ionized hydrogen (atomic mass mH, mean molecular weight µ = 0.5) at temperature T, where T6 is in MK:
� kT
= � 128 km s−1��
T6
cs =
µmH
2) Parker critical point rc where the wind exceeds the sound speed (in the rightmost expression, M∗ and R∗are in solar units):
rc = GM∗
= (5.8 R∗) M∗
2c2
R∗T6
s
3-4) (Eq. 15-16 in Cranmer 2004) Solution for wind speed u(r) using Lambert W0 and W−1 functions:
�−4
� −1 �
exp � 4 � 1 −rc
D(r) = �r
rc
r
 
� −W0[−D(r)] r ≤rc
cs
u =
� −W−1[−D(r)] r ≥rc
cs

Our targets’ critical radius is ∼3R∗compared to ∼6R⊙for the solar wind. This difference arises because our targets’ hot coronae have higher thermal pressure gradients that cause rapid acceleration.
6.6. Wind magnitude at planet: Density
We determine wind density using mass continuity on open field lines:
ρ(r) = ˙M fopen(r) 4πr2u (7) where fopen(r) is the fraction of a sphere with radius r surrounding the star that is pierced by open field lines. At the orbital distances of our targets, we assume fopen(r) ≈1, which may slightly underestimate density near GJ 367 b at 3.4R∗. The Parker solution assumes the wind is purely radial, enabling an analytic solution, whereas a numerically-integrated Weber-Davis wind (Weber & Davis 1967) allows for azimuthal wind velocity caused by stellar rotation. We compared a radial Parker wind model for YZ Cet to the Weber-Davis SPI Radio Search 13 solution of Pineda & Villadsen (2023)
Villadsen et al.

| 52 844 e |  |  |
| --- | --- | --- |
| 7 GJ 12 LHS 3 eferenc |  | Cet 3 |
|  | GJ 36 80 1678 R | 486 YZ GJ 332 ox Cen |
| 37 LHS LTT Wind Model Uncertainty |  | GJ Pr |
|  |  |  |

|  |  |  |  |
| --- | --- | --- | --- |
|  |  |  |  |
| Wind Model Uncertainty |  |  |  |
| Predicted for for R |  |  | AW = R p |
|  | Predicted for for R |  | AW = R p |
| Predicted Flux - R Observed | Predicted Flux - R Observed |  | eference |

SPI Power for RAW = Rp (erg s 1)
SPI Flux F  ( Jy GHz)

| Wind Model Uncertainty | e : 1 0 0 Jy |
| --- | --- |
| (impacts "Observed") | R e fe re n c |

|  | 0 0 Jy |
| --- | --- |
| Wind Model Uncertainty | R e fe re n ce : 1 |

Alfven Wing Radius (Rp)
Planet Field Bp (G)
Predicted for for Bp = 0.5 G
Reference RAW UL
Bwind = B
Observed
Distance from Star (R * )
Distance from Star (R * )
(Top left) Predicted SPI power for an Alfv´en wing the width of the planet. For a given system,
the predicted power declines sharply with orbital distance, but the later spectral types of our comparison
sample somewhat offset their greater orbital distances. The error bar at left shows the scale of uncertainty
for all systems due to inferring wind properties from rotation-activity relations (Section 8.9). (Top right)
Observed radio upper limits or detections for our targets and comparison sample, compared to predicted
is an estimate of spectrally-integrated flux. For comparison systems YZ cet and Prox Cen, the error bars
denote the range of observed coherent burst flux densities in different epochs, whereas for GJ 3323 the error
bars show the uncertainty on a single-epoch detection. (Bottom left) The Alfv´en wing radius needed to
match the observational constraints, compared to its predicted size for an Earth-like planet field of 0.5 G.
We also show how the sensitivity to Alfv´en wing size varies with orbital distance, for a radio sensitivity
of 100 µJy on the reference system. Closer-in orbits offer sensitivity to smaller Alfv´en wings, and thus to
weaker exoplanet magnetic fields. The error bar at left shows the wind model uncertainty for all systems
that enters when converting flux density to Alfv´en wing radius, which is not included in the 3σ upper limits.
(Bottom right) Exoplanet magnetic dipole field strength consistent with observed flux densities or upper

[Figure: Figure 3.]
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