JuanCastillo29/space-mission-intelligence-data huggingface dataset:
≲15 au (e.g. Shober et al 2024). Comets that pass sufficiently close to the Sun lose parts of their outer layers, exposing fresher, relatively unprocessed material beneath, as seen with comet 67P (Bockel´ee-Morvan et al, 2016). Additionally, a comet’s evolution is influenced by its rotational state. For instance, the high obliquity of comet 67P results in pronounced
7The preferred target is a dynamically new comet, or even an interstellar object, but if a suitable target is not found another long-period comet will be selected.

seasonal effects that result in uneven surface evolution (Bockel´ee-Morvan et al, 2016), with northern and southern hemispheres being preferentially illuminated during different parts of its orbit as a result of the orientation of its spin axis. On the other hand, the complex rotation of comet 103P/Hartley allows solar illumination over its entire surface at perihelion (A’Hearn et al, 2011; Guilbert-Lepoutre et al, 2023). A significant fraction of JFCs shows activity near aphelion (e.g., Kelley et al, 2013), indicating that regions deep within the nucleus may remain warm enough for volatile sublimation. Rosetta’s observations of 67P revealed that JFC activity is a complex interplay of near surface sublimation and the recondensation of volatiles migrating upward from deeper layers. This process can obscure the true bulk abundances of highly volatile species such as O2 (Luspay-Kuti et al, 2022).
At sufficiently low cometary activity levels, the Rosetta mission found that the solar wind can reach the surface where it sputters refractory atoms such as Na, K, Si, and Ca (Wurz et al, 2015). However, the interaction between the solar wind and dust in the coma has not been extensively studied (Hor´anyi and Deca, 2024), though it does have an observable effect on dust tails (Price et al, 2019). For airless bodies like asteroids, the Moon, and Mercury, exposure to the solar wind alters the properties of their surfaces through processes such as space weathering, implantation of solar wind ions, and the formation of molecules like hydrogen (H2) and water (H2O). This exposure can also result in the release of gases—including Na, K, S, and Ca—through processes such as electron-stimulated desorption and sputtering (Chapman, 1996; Nagao et al, 2011; Wurz et al, 2022). Cometary comae, along with their formation and interaction with the stellar wind, are discussed further in Vrignaud et al (2026).
For most exocomets, it is unknown whether their nuclei are volatile-rich, as is the case for typical Solar System comets, or volatile-poor. We therefore now turn our attention to the Solar System’s asteroids, as they may provide better analogues than Solar System comets for some bodies, for example the transiting exocomets of β Pic. The majority of asteroids in the Solar System today reside in the main asteroid belt, a region between the orbits of Mars and Jupiter. This belt contains millions of small rocky bodies orbiting the Sun, representing a dynamically evolving remnant of the planetesimals that formed within the early solar nebula. The total mass of the asteroid belt is estimated at approximately 1.8 × 10−9M⊙(∼6 × 10−4M⊕), inferred by analyzing planetary perturbations from 300 major asteroids and a ring of mass for the smaller asteroids, combined with statistical extrapolation of asteroid sizes (Krasinsky et al, 2002).

This mass is significantly lower than the estimated primordial mass of the asteroid belt, which, based on interpolated mass densities needed to form the terrestrial planets and Jupiter’s core, was on the order of 1 M⊕(Weidenschilling, 1977; Morbidelli et al, 2009). This discrepancy implies that the primordial asteroid belt underwent substantial mass depletion and dynamical excitation, driven by gravitational perturbations and collisional processes. Morbidelli et al (2015) divide asteroid evolution into three major phases: (1) the early stage during the existence of the gaseous proto-planetary disc, when giant planets formed and migrated, (2) the period following gas dissipation, during which the asteroid belt experienced additional dynamical perturbations from planetary instabilities, and (3) the long-term evolution influenced by collisional fragmentation (Bottke et al, 2023) and resonant interactions with planetary orbits (Lecar and Franklin, 1997; Nagasawa et al, 2002). Repeated collisions between asteroids lead to fragmentation, generating smaller debris, reshaping surface properties through regolith redistribution and space weathering (Hapke, 2001; Brunetto et al, 2015), and modifying the size distribution of the asteroid population (Bottke et al, 2015). These collisions also contribute to the formation of asteroid families: groups of genetically related bodies that share similar orbital characteristics (Nesvorn´y et al, 2015; Novakovi´c et al, 2022, and references therein). In some cases, particularly energetic impacts can excavate subsurface material, providing unique information about an asteroid’s internal structure and composition. A notable example is (596) Scheila, which underwent a collision in December 2010, temporarily displaying comet-like activity (Larson, 2010; Ishiguro et al, 2011, and see Figure 4). Follow-up observations indicated that the impact exposed fresh, unweathered material from the subsurface, temporarily altering Scheila’s spectral characteristics (Bodewits et al, 2014; Hasegawa et al, 2022).
Interstellar objects (ISOs) form a bridge between Solar System small bodies and exocomets: they formed in a different planetary system, but have been observed in our own Solar System on hyperbolic orbits with respect to the Sun (Seligman and Moro-Mart´ın, 2022)8. While probably abundant in the interstellar medium, they would rarely pass through planetary systems. As such, they likely constitute only a small fraction, if any, of the observed exocomets. Nevertheless, they afford us the only opportunity we realistically have to study extrasolar small bodies up close.
The first such interstellar objects were detected relatively recently: 1I/‘Oumuamua in 2017, 2I/Borisov in 2019, and 3I/ATLAS in 2025 (Meech et al, 2017; Williams et al, 2017; Borisov et al, 2019; Jewitt and Luu, 2019; Denneau et al, 2025; Seligman et al, 2025). 1I/‘Oumuamua and 3I/ATLAS were discovered by surveys searching for Solar System objects (Pan-STARRS [Chambers et al 2016] and ATLAS [Tonry et al 2018]
8Here we refer to bodies whose orbital eccentricities are much greater than unity: a number of Solar System comets have been observed with slightly hyperbolic osculating orbits owing to measurement error, or the effects of planetary perturbations and non-gravitational forces around perihelion. A fraction of these may actually leave the Solar System.

respectively), while 2I/Borisov was discovered by an amateur astronomer (Borisov et al, 2019). In the near future, we can expect the Legacy Survey of Space and Time (LSST), at the Vera C. Rubin Observatory, to produce several dozen more detections (Moro-Mart´ın et al, 2009; Cook et al, 2016; Engelhardt et al, 2017; Ivezi´c et al, 2019; Hoover et al, 2022; Marˇceta and Seligman, 2023; Dorsey et al, 2025; Forbes et al, 2025), enabling statistical characterisation of their properties such as compositions and kinematics.
Like exocomets, ISOs can be used to study planetesimal formation in a range of environments (Raymond et al, 2018). Individual ISOs can be studied in far more depth than individual exocometary events, with detailed observations possible over a large range of wavelengths and heliocentric distances (’Oumuamua ISSI Team et al, 2019; Bolin et al, 2020). In the future, there is also the possibility of visiting ISOs (Abell et al, 2024) with missions such as Comet Interceptor (Jones et al, 2024; Yaginuma et al, 2025; Sanchez and Snodgrass, 2025).
Here we summarise the observations of the three hitherto-detected ISOs, and their implications. Detailed reviews of 1I/‘Oumuamua and 2I/Borisov (prior to the discovery of 3I/ATLAS) can be found in ’Oumuamua ISSI Team et al (2019); Seligman and Moro-Mart´ın (2022); Jewitt and Seligman (2023); Jewitt (2024); Fitzsimmons et al (2024). We give somewhat more details about 3I/ATLAS, as at present there is no review of this object. At the time of writing, 3I/ATLAS is still under observation.
2I/Borisov was detected on its inward approach, before it passed perihelion at 2.0 au. It passed just outside the orbit of Mars, and like 1I/‘Oumuamua it exhibited non-gravitational acceleration, most likely caused by outgassing as it warmed on approach to the Sun (Hui et al, 2020). In contrast to 1I/‘Oumuamua, the brightness and observational geometry conditions of 2I/Borisov allowed for the characterization of its activity and the direct detection of multiple volatiles, including CN, C2, OH, and CO (Fitzsimmons et al, 2019; Bodewits et al, 2020; Cordiner et al, 2020). 2I/Borisov was also sub-km in size (Jewitt et al, 2017; Jewitt and Luu, 2019; Hui et al, 2020). It underwent a large outburst at a heliocentric distance of 2.8 au after its perihelion passage, which ejected a small fraction (∼10−4) of its mass (Jewitt et al, 2020). Preoutburst HST observations of 2I/Borisov did not detect significant lightcurve variation (Bolin et al, 2020).
The composition of 2I/Borisov was broadly similar in some respects to Solar System comets (Opitom et al, 2021): its dust had a comparable redness, and its gas production rates showed a low C2 to CN ratio, marking it as carbon-depleted (Opitom et al, 2019; A’Hearn et al, 1995). It also had Ni and Fe abundances similar to Solar System comets (Guzik and Drahus, 2021; Manfroid et al, 2021). However, 2I/Borisov was chemically distinct in other ways. Its abundance of CO relative to water and HCN was exceptionally high (Cordiner et al, 2020; Bodewits et al, 2020), demonstrating that ISOs are not exactly analogues of Solar System comets. The anomalously high hypervolatile ratio could be due to 2I/Borisov forming at a large distance from its host star (Cordiner et al, 2020), or alternatively forming around an M dwarf (Bodewits et al, 2020).

These hypotheses were further explored and compared to Solar System comets in Seligman et al (2022).
In terms of their formation, there may be little difference between ISOs and observable exocomets, other than that ISOs were ejected from their systems, whilst exocomets got driven inwards towards their host stars. The dynamical mechanisms that can form ISOs closely overlap with those that can form exocomets, and in many cases a single process may form both comets and ISOs at the same time (Figure 1a-d, f). The mechanisms that can generate ISOs are those which change a body’s energy (or equivalently, semimajor axis); these include stellar flybys (Pfalzner and Govind, 2021), and scattering by a planet (Brasser et al, 2006). Mechanisms that approximately conserve a body’s energy (semimajor axis), such as MMRs and secular interactions, are unable to directly generate ISOs; however, they can indirectly generate ISOs by exciting small bodies onto planet-crossing orbits, so that those bodies are then scattered and ejected by planets (Figure 1c-e). If a planet undergoes migration driven by its scattering of planetesimals, an outward-migrating planet would scatter planetesimals inwards where they could become exocomets (although this likely requires more interior planets to further reduce the periapsis distances, Bonsor et al 2014), whilst an inward-migrating planet would scatter planetesimals outward where they could become ISOs. If just one planet were involved, then this could mean that ISOs are more typically sampled from debris discs closer to their central stars (because the planet would need to be exterior to the disc and migrating inwards), whilst exocomets would be typically sampled from debris discs further away (with the planet interior and migrating outwards). This could lead to a detectable difference between ISO and exocomet populations, which may not be the case if other mechanisms dominate instead. However, multiple planets would complicate this picture; for example, an outward-migrating planet could scatter bodies inwards, where they are then ejected by an inner planet (as is thought to have happened during the migration of Neptune; Fernandez and Ip 1984; Malhotra 1993). This then has implications for the volatile content of the ISOs, as they can be heated during the time when their periapsis has been reduced around the host star, possibly to a greater extent than Oort Cloud comets (Gkotsinas et al, 2024).
One mechanism that can produce ISOs considerably more efficiently than exocomets is post-main-sequence mass loss of a star (see Section 10), which causes loosely-bound bodies in Oort Cloud analogues to become unbound (Veras et al, 2011; Hansen and Zuckerman, 2017; Moro-Mart´ın, 2019). Hence, some ISOs are probably produced via the post-main sequence ejection of a large portion of exo-Oort clouds. These ISOs could have fundamentally different properties and formation histories from ISOs that are ejected from an inner planetary system by other mechanisms, but they are predicted to form a minority of ISOs overall (Veras et al, 2020b; Levine et al, 2023).
b; Taylor and Seligman, 2025).
Stellar metallicities and abundances also depend on their formation location, as well as age (Chiappini et al, 2001; Nordstr¨om et al, 2004). Therefore, the chemical composition of interstellar objects and the protoplanetary discs in which they form (Cabral et al, 2023) should vary with both age and formation location within the Galaxy. In principle, ISOs with different Galactic velocities should trace the variable
10Alpha elements are those formed by successive additions of helium nuclei, such as oxygen, magnesium, etc.

[Table: Table 1 Summary of broad stellar wind characteristics across spectral types and evolutionary stages. The ‘Type’ column gives the spectral type and subtype range. The ‘Stage’ column describes the evolutionary phase of the star. The ‘Mass’ and ‘Temperature’ columns list the ranges of initial masses and effective temperatures for stars in each category. The ‘Wind type’ column describes the type of wind produced.
Spectral Stage Mass Temperature Wind type Type (M⊙) (K)
O0 – B2 Main Sequence ≳8.0 ≳20000 Hot star B3 – B5 Main Sequence ∼8.0 – 4.0 ∼20000 – 15000 Hot star (weak) B6 – B9 Main Sequence ∼4.0 – 3.0 ∼15000 – 10000 Negligible A0 – A9 Main Sequence ∼3.0 – 1.5 ∼10000 – 7500 Negligible F0 – F4 Main Sequence ∼1.5 – 1.1 ∼7500 – 6500 Chromospheric (weak) F5 – M9 Main Sequence ≲1.1 ≲7150 Chromospheric
K0 – M5 Evolved ≳8.0 ∼5000 – 3000 Giant star K0 – M5 Evolved ≲8.0 ∼5000 – 3000 Chromospheric M6 – M9 Evolved ≲8.0 ≲3000 Giant star
β Pictoris (A6) Main Sequence 1.8 ∼8000 Negligible Sun (G2) Main Sequence 1.0 5778 Chromospheric]

compositions of planetesimals formed as a function of stellar population, which could be probed statistically with Rubin/LSST.
Using ISOs to probe planetesimal formation does have several issues. One is that their size, shape and volatile content may have changed during the ejection process. This is because ISOs could have been tidally disrupted during ejection, for example if they passed close to a giant planet or star and were ejected (´Cuk, 2018; Rafikov, 2018; Zhang and Lin, 2020; Veras, 2020). Other issues include thermal processing in their birth system before or during ejection (Gkotsinas et al, 2024), and exposure to cosmic ray radiation during their sojourn through the Galaxy (e.g., Jackson and Desch, 2021).
branch evolution

Once a star leaves the Main Sequence, it undergoes significant changes: large increases in luminosity and radius on the red giant branch (RGB) and asymptotic giant branch (AGB), strong mass loss especially at the AGB tip, and then a declining luminosity as a WD. These changes to the star have significant effects on orbiting bodies: both their physical evolution and their orbital evolution, and thence the global orbital dynamics of the entire system. We discuss these in the closing sections of this work.

[Table: Table 2 Mass loss rates and terminal wind speeds for different types of stellar winds. The ‘Pressure at 1 au’ column gives the wind ram pressure
1]

2 ρv2, assuming a steady, uniform wind and that v∞has been reached at a distance of 1 au. In this case ˙M = 4πr2ρv and the pressure becomes p = 1
4πr2 . The following references are used: 1Kudritzki and Puls (2000); 2Cranmer and Winebarger (2019); 3Wood et al (2024); 4H¨ofner and Freytag (2022); 5Vidotto (2021); 6Bruhweiler et al (1991).
2 ˙ Mv
Wind type Mass loss rate ˙M Speed v∞ Pressure at 1 au (M⊙/yr) (km/s) (Pa)
10−2–10−1
Hot star (O0–B5)1 10−7–10−6
mass, expand in size to the orbit of our Solar System’s Earth or Mars, and increase
their luminosity by many orders of magnitude. This means that physical effects on
orbiting bodies during the giant branch phases are “supercharged” compared to main-
sequence systems because of the significant changes to the stellar mass, radius and
luminosity. The most extreme mass loss, radius and luminosity values are attained
during the second giant phase (the AGB), rather than the first (the RGB).
Considering collisions between a body and the star, the increased size of the giant
branch star drastically changes the minimum survivable orbital periapsis of a small
body from about a solar radius to about 1-5 au. The envelope boundary of the star
is likely to be a good representation of the maximum destruction distance. Small
bodies are sufficiently small that tidal orbital decay would not increase this distance,
whereas for giant planets the minimum radius for survival can be significantly beyond
the maximum stellar radius (Mustill and Villaver, 2012; Adams and Bloch, 2013;
Madappatt et al, 2016). However, for small bodies radiation effects will likely result
in destruction beyond the stellar radius, as discussed below.
RGB and especially AGB stars have strong winds (see Appendix A.4), owing to an
interplay of low surface gravity, stellar pulsations, dust formation, and high radiation
pressure (H¨ofner and Freytag, 2022). Mass loss from the host star will have both
direct physical effects on the bodies, and indirect gravitational effects on their orbits.
The gravitational effect arises as the orbit evolves in response to the star’s changing
gravitational potential. If the body is close enough to the star (within a few hundred
au of the parent star, so that the orbital timescale is short compared to the mass-loss
shape in the form of eccentricity, inclination and orientation in space (Hadjidemetriou,
1963; Omarov, 1964). However, for exo-Oort cloud bodies on wider orbits, the orbit

11This is sometimes known as “adiabatic” mass loss. 12As measured by the semimajor axis. For isotropic mass loss, the orbital periapsis distance always increases (Veras et al, 2011).
the location of the body along the orbit (Veras et al, 2011; Veras and Wyatt, 2012). Significantly anistropic stellar mass loss can give rise to more complex changes in the orbit (Parriott and Alcock, 1998; Veras et al, 2013; Stone et al, 2015; Dosopoulou and Kalogera, 2016a,b; Akiba et al, 2024). The stellar mass loss takes place through a strong wind (Appendix A.4), but the physical and chemical consequences of giant branch stellar winds impinging on the surface of a body have not yet been significantly explored. Veras et al (2015b) presented a general formalism for the orbital effects of drag from giant branch stellar winds, finding it to be significant, though detailed treatments including inhomogeneities in the wind have not been performed. Radiative effects are, however, much better studied; see Vrignaud et al (2026) for a discussion. Stellar luminosities on the giant branch attain values of thousands of solar luminosities, and this has extreme effects on heating, sublimation and devolatilisation of bodies (Jura and Xu, 2010; Jura et al, 2012; Malamud and Perets, 2016, 2017a,b; Katz, 2018; Harrison et al, 2021; Li et al, 2024). The result of these processes could fundamentally change the shape, chemistry and survival prospects for these bodies. Furthermore, such extreme radiation will change the spin and orbit evolution of these bodies through a supercharged version of the YORP (Veras et al, 2014b; Veras and Scheeres, 2020) and Yarkovsky (Veras et al, 2015b, 2019) effects (see Section 6), leading to both widespread destruction due to rapid and irreversible rotational spin-up, and orbital dispersal of the survivors and fragments. The Yarkovsky effect dominates over the Poynting–Robertson (PR) effect, an additional radiative force (Burns et al, 1979; Gustafson, 1994), for bodies orbiting giant branch stars (Veras et al, 2015b), though the PR effect may help bring small (sub-m) fragments down to the surface of a young white dwarf (Veras et al, 2022).
Direct evidence for these processes occurring is lacking, although the central star of the Helix Nebula shows a strong IR excess that may be due to a large population of disintegrating small bodies (Marshall et al, 2023), while the central star of the planetary nebula WeSb 1 has recently been observed to undergo occultations attributed to dust clouds (Budaj et al, 2025). On the other hand, the white dwarfs that emerge after the AGB phase show significant evidence for cometary phenomena arising from some form of small bodies—whether volatile-rich or volatile-poor—which likely stems ultimately from the aforementioned changes to their orbits. This is discussed in depth in the next Section.
to white dwarfs

As a result of the changes discussed in Section 10, a young white dwarf (WD) is expected to have no planets or small bodies within at least a few au, owing to the combination of stellar mass loss, tidal orbital decay, and strong radiation effects. However, a high fraction of WDs do in fact possess some material in very close proximity to the star, within roughly 1 R⊙. The evidence for this comes in several forms: photometric transits of material and/or parent bodies across the stellar disc; dusty circumstellar material inferred from an IR excess; circumstellar gas discs inferred from emission or absorption lines; and metals in the stellar photosphere deposited by infalling material,
and seen in absorption (Figure 6). These phenomena are thought to arise as a result of bodies on wide orbits being destabilised around the WD following stellar mass loss. Unfortunately, the source regions are difficult to probe observationally (Fig 6A,E). Some giant planet candidates have been imaged with JWST (Mullally et al, 2024), but subsequent epochs have revealed that they were false positives (Mullally et al, 2026)13. WDs typically descend from progenitors somewhat more massive than the Sun, as most solar-mass stars have not yet had time to evolve off the Main Sequence (Koester et al, 2014). Their progenitors are therefore often A-type stars, which are known to host debris discs very commonly (around a quarter or a third show detectable IR excess, Su et al, 2006; Thureau et al, 2014; Bannister et al, 2025), but little is known about the planets at several to 10s of au that can drive the dynamics. This has given theorists great freedom in speculating on the exact dynamical processes driving the delivery. For further reading on this topic, see the recent reviews by Veras et al (2024); Xu et al (2024); Malamud (2026); Mustill (2026).