JuanCastillo29/space-mission-intelligence-data huggingface dataset:
A handful of WDs are now known to possess material that passes in front of the star, detected as transits in broadband photometric lightcurves. Transits of WDs by planets can easily achieve 100% transit depth owing to the small stellar radius, and transits of small planets, asteroids and comets can be detected much more easily than towards MS stars (Agol, 2011). However, the small size of the WD makes the a priori geometric probability of a transit small. An early search for transits of WDs in WASP data was unsuccessful (Faedi et al, 2011). The first successful detection was with the Kepler space telescope by Vanderburg et al (2015), who detected asymmetric transits of WD1145+017. These transits, with their steep ingress and shallower egress, were attributed to a comet-like morphology arising from a parent body possibly surrounded by a coma and with a trailing tail. Several orbital periods were identified at ∼4.5 hr, suggesting multiple bodies close to the Roche radius, possibly indicative of a recent break-up event, and that the cometary tails may arise from tidal disruption, rather than sublimation (Veras et al, 2017). Multi-wavelength photometry of the transits has shown that they are independent of wavelength, posing challenges for our understanding of the dust being released by the parent body: at such small orbital radii, velocities are high and therefore small dust should abound from fragmentation, yet small grains would induce a significant colour dependence to the transits. Possible resolutions are that small grains sublimate extremely quickly (Xu et al, 2018a), or that the dust is sufficiently abundant to be optically thick at all wavelengths (Izquierdo et al, 2018). Recently this system has become more quiescent (Aungwerojwit et al, 2024), showing that these transiting phenomena are transitory on human-relevant timescales. This raises the prospect that revisiting previously-surveyed WDs with non-detections could
13An earlier discovery, a super-Jovian companion to WD0806-661 at 2500 au projected separation, is often regarded as a brown dwarf owing to its large mass and very large separation Luhman et al (2011); Rodriguez et al (2011).
This is in contrast not only to WD1145+017 (Alonso et al, 2016; Izquierdo et al, 2018; Xu et al, 2018a) but also to several other WDs which also exhibit grey
14The 65:1 commensurability is between the observed periods; this likely implies not a 65:1 mean motion resonance, but a 66:65.
15Previously, (Hallakoun et al, 2017) had detected bluing during transits of WD1145+017, attributed to gas absorption during transit rather than dust properties.

transits: WD J0328-1219 (Gary and Kaye, 2024), ZTF J1944+4557 (Guidry et al, 2025), SBSS 1232+563 (Hermes et al, 2025), and WD1054-226 (Farihi et al, 2022); especially in this latter case, the lower temperature of the dust means that the lack of observed small grains cannot be attributed to their sublimation (Korth et al, 2026a).
Discs orbiting WDs can comprise dust, gas, or both. The observed discs all lie within ∼1 R⊙, consistent with the outcome of tidal disruption and subsequent orbital decay. The descendants of main-sequence debris discs, on much wider orbits, are unfortunately mostly undetectable owing to the extremely low luminosities of WDs (Bonsor and Wyatt, 2010). An exception may be the debris disc around the very young WD at the centre of the Helix Nebula (Su et al, 2007); this object may indeed be showing cometary activity and the dust could be produced by outgassing, rather than collisions as in typical debris discs (Marshall et al, 2023). The dust discs orbiting WDs are inferred indirectly from excess IR emission compared to the stellar spectral energy distribution: they cannot be resolved, unlike many MS debris discs, owing to their extremely small radii. Over 1% of WDs exhibit such IR excess (Wilson et al, 2019; Madurga Favieres et al, 2024; Murillo-Ojeda et al, 2026). Where good-quality mid-IR spectra exist, mineralogy is possible: analysis of Spitzer IRS data showed that the G29-38 disc contains crystalline pyroxene, amorphous silicates, water ice, metal sulphides, and possible amorphous carbon (Reach et al, 2009). The improved sensitivity afforded by JWST is now enabling analyses of many more discs (Swan et al, 2024; Farihi et al, 2025). Many IR excesses are variable on timescales of months to years (e.g., Xu et al, 2018b; Swan et al, 2020, 2024), indicating significant changes in the quantity of dust in the systems (Fig 6C,G). The near-WD environment is clearly highly dynamic, with ongoing production of dust grains through fresh delivery and disruption of new parent bodies, collisions, and/or evaporative outbursts. Discs of circumstellar gas are detected spectroscopically, both in absorption (which requires an unusually favourable geometrical alignment to have the disc edge-on) or emission (which does not). There are fewer detections of gas discs than dust discs, but they do permit detailed measurements of the elemental abundances of the material (Steele et al, 2021). Unlike typical astrophysical gas discs, these are formed from vapourised rock, and are not dominated by H/He. Disc kinematics can be inferred from the Keplerian line profiles, and they sometimes show an orbital eccentricity and relativistic precession (Manser et al, 2016). The disc emission of SDSS J1228 shows variability on a 2-hour timescale, attributed to a disturbance in the disc by an embedded asteroid or minor planet (Manser et al, 2019), orbiting interior to the fluid Roche limit but held together by its internal strength.
In retrospect, we can identify the first detection of extrasolar planetary material by Humanity (Farihi et al, 2008): van Maanen (1917) identified a peculiarly faint highproper motion star of spectral class F0, which we now know to be the first discovery of a metal-polluted white dwarf.

Metals are not expected to be present in most WD atmospheres, since the strong surface gravity should cause material heavier than H/He to sink below the photosphere on astrophysically short timescales (ranging, depending on the atmospheric temperature and whether H or He dominates, from days to ∼1 Myr; see for example Wyatt et al 2014). This implies ongoing or recent accretion of non-stellar material, now thought to arise from the accretion of planets, moons, or small bodies onto the WD (Fig 6D). This is an extremely common phenomenon: depending on the subpopulation of WDs studied, up to 50% exhibit spectroscopically detected metal lines (e.g., Koester et al, 2014; Wilson et al, 2019; Ould Rouis et al, 2024). From the elemental abundances measured in WD atmospheres, we can determine the bulk compositions of the progenitor bodies, and compare them to bodies of our own Solar System (Fig 6H): readers are referred to the reviews by Jura and Young (2014) and Xu et al (2024) for the geochemical conclusions. Information obtained in this manner is complementary to our study of the atmospheres of exoplanets via transmission or emission spectroscopy (which tells us about the surface but needs to be coupled to interior models to learn about the interior), and to knowledge of the bulk density from a planet’s mass and radius (which provides weak constraints on the bulk composition with many degeneracies). Disadvantages are that this can only be performed for destroyed bodies after the star has become a WD, that we may not know the primordial nature of the parent body (volatile-rich, volatile-poor, etc.), that the material has undergone thermal processing on the RGB/AGB altering its composition to some degree (see Sec 10), and that not all components of the destroyed parent body may be accreted at the same time and rate (Brouwers et al, 2023). Currently over 1700 polluted WDs have been identified (Williams et al, 2024). This number is expected to grow rapidly in the near future, as the combination of largescale spectroscopic surveys such as SDSS, J-PAS, LAMOST and Gaia, combined with machine learning techniques, allows for efficient screening of vast numbers of spectra to identify candidate polluted WDs and efficiently target high-resolution follow-up spectroscopy for confirmation and to increase the number of species detected (Kao et al, 2024; P´erez-Couto et al, 2024). Future surveys such as 4MOST (de Jong, 2019) will soon provide over 105 high-quality WD spectra for large-scale studies of WD abundances (Chiappini et al, 2019). While we typically infer the accretion of material from planets or small bodies indirectly via the presence of metals in the WD photosphere, direct evidence of accretion has been found at G29-38 in the form of X-rays emitted by cooling of the accreted material (Cunningham et al, 2022).
The above phenomena are thought to reflect stages in a process of dynamical excitation of the eccentricities of the parent bodies. The initial excitation may arise from the mechanisms discussed in Section 3, such as scattering, forcing from mean motion or secular resonances, or the Lidov–Kozai effect. With sufficiently extreme eccentricity excitation, a body’s periapsis is forced down to very small values, after which it may experience tidal disruption, circularisation, pulverisation, sublimation, and ultimately accretion onto the WD. A schematic of these processes is shown in Figure 6.

Secular resonance
sweeping
Parent body belt
[Figure: Fig. 7 Schematic overview of how changes wrought by stellar evolution effect the delivery of small bodies to white dwarfs. We show an imaginary system initially of three planets with two source belts, one destabilised by a mean-motion resonance, and one destabilised by a secular resonance. As a result of stellar mass loss, the MMR broadens, while previously stable bodies are moved by radiation forces into the unstable region. The secular resonance moves as a result both of the engulfment of one planet by the giant star, and of the differential motion of the remaining planets as tidal forces compete with the effects of mass loss.]

Delivery of bodies to the region close enough to the WD for the above phenomena to occur is possible for bodies originating on wide orbits similar to the Solar System’s scattered disc or Oort Cloud objects (Alcock et al, 1986) These orbits are affected by near-impulsive anisotropic mass loss as the star transitions to a WD, known as WD “kicks”, and then continue to evolve under the Galactic tide and stellar flybys (Veras et al, 2014a; Stone et al, 2015; Veras et al, 2014c; O’Connor et al, 2023; Pham and Rein, 2024) and in some cases internal interactions within the planetesimal population (Akiba et al, 2024). While (exo-)Oort Cloud dynamics are identical for bodies orbiting a MS star or a WD, the combination of reduced stellar mass and increased orbital radius means that a body experiences a stronger effect from the Galactic tide following AGB mass loss. The presence of a planet or star between the Oort Cloud and the WD can reduce delivery of Oort Cloud material to the WD, as they can suppress eccentricity fluctuations induced by the Galactic tide and eject bodies that do get too close, preventing them from crossing the planet’s orbit (O’Connor et al, 2023;
The following are some of the major open questions in the field of exocomet evolution. Future studies answering these would help us understand these objects and the processes occurring in their planetary systems.
• What processes drive comets down to small periapsides around MS stars? Are the main driving bodies planets, stellar binary companions, or the Galaxy as a whole, and what are the exact dynamical pathway(s)?
• What process drives comets down to small periapsides around WDs? Is this the same as for MS stars, or do different mechanisms dominate? In both of these cases, we are currently hindered by our lack of knowledge of the architectures of outer regions of planetary systems.
• Are certain planetary system architectures required to produce low-periapsis exocomets, or are a range of architectures and mechanisms possible?
• How can we observationally probe planets in the outer regions of planetary systems, where comets probably originate?
• What is the relative importance of stellar flybys and the tidal field for exo-Oort Cloud depletion? How does this change depending on a star’s migration history within the Galaxy?
• What surprises can we expect from the upcoming census of ISOs from The Vera C. Rubin Observatory’s Legacy Survey of Space and Time? What will be the balance between more traditional objects such as 2I/Borisov and 3I/ATLAS, and more exotic objects such as 1I/’Oumuamua?
• How unevolved are the nuclei of dynamically new comets? What insights can we expect from Comet Interceptor?
• Is the physical evolution of exo-cometary nuclei the same as that for Solar System comets? How can we probe this when the individual bodies themselves are too small to detect?
• To what extent do asteroids on cometary orbits, low-activity comets, mainbelt comets, and related populations represent distinct reservoirs versus different evolutionary states of a common feedstock?
• What is the primordial nature of the bodies responsible for WD pollution? Were they originally volatile-poor, or have they become so by devolatilisation when the star was a giant?
We gratefully acknowledge support by the International Space Science Insitute, ISSI, Bern, for supporting and hosting the workshop on “Exocomets: Bridging our Understanding of Minor Bodies in Solar and Exoplanetary Systems”, during which this work was initiated in July 2024.

We thank two anonymous reviewers for their constructive comments which improved the paper.
M.T.B. appreciates support by the Rutherford Discovery Fellowships from New Zealand Government funding, administered by the Royal Society Te Ap¯arangi. M.J.H. acknowledges support from the Science and Technology Facilities Council through grant ST/W507726/1. K.J. acknowledges support by grants from eSSENCE (grant number eSSENCE@LU 9:3), the Swedish National Research Council (project number 2023-05307), The Crafoord foundation and the Royal Physiographic Society of Lund, through The Fund of the Walter Gyllenberg Foundation. R.K. would like to acknowledge the support from “L’Oreal UNESCO For Women in Science” National program for Bulgaria. This work benefited from discussions within ISSI International Team project 504 “The Life Cycle of Comets” funded by the International Space Science Institute (ISSI) in Bern. R.K. acknowledges partial support by project KP-06-D002/3 “CLIC – Cometary Life Cycle,” carried out under the PROMYS (Promotion of Young Scientists) component of the Swiss–Bulgarian Research Programme and supported by the Bulgarian National Science Fund. A.J.M. acknowledges support from the Swedish Research Council (Project Grant 2022-04043) and the Swedish National Space Agency (Career Grant 2023-00146). T.D.P. is supported by a UKRI Stephen Hawking Fellowship and a Warwick Prize Fellowship, the latter made possible by a generous philanthropic donation. X.P.C. acknowledges financial support from the Spanish National Programme for the Promotion of Talent and its Employability grant PRE2022-104959 cofunded by the European Social Fund. D.Z.S. is supported by an NSF Astronomy and Astrophysics Postdoctoral Fellowship under award AST-2303553. This research award is partially funded by a generous gift of Charles Simonyi to the NSF Division of Astronomical Sciences. The award is made in recognition of significant contributions to Rubin Observatory’s Legacy Survey of Space and Time.
High-mass (O0–B5) type stars

High-mass stars with spectral types O0–B5 have stellar winds driven by the star’s radiation pressure. These stars emit copious amounts of ultraviolet light, which exerts pressure on the outermost stellar layer. The radiation pressure primarily acts on metal ions through absorption and scattering in spectral lines, leading to the acceleration of the wind (Castor et al, 1975; Sander, 2023). For these hot stars, the wind velocity profile v(r) is often (Lamers and Cassinelli, 1999; Krtiˇcka and Kub´at, 2011) described by a β law �β
v(r) = v0 + (v∞−v0) � 1 −R⋆
. (A1)
r The v0 value sets the wind velocity at an inner boundary but plays little role at large distances from the star; it is sometimes set to v0 = 0 (Krtiˇcka and Kub´at, 2011).

The value β determines the acceleration at different stellar distances. A value of β = 1
2 corresponds to an inverse square accelerating force F ∝r−2 (Chandrasekhar, 1934). Values of β = 0.7 to 1 are adopted for O star winds (Puls et al, 1996), whilst higher values are possible for evolved stars (Crowther et al, 2006) such as β ∼2 for BA-supergiants (Puls et al, 2008). More sophisticated models may vary β (e.g. for Wolf-Rayet stars, Gr¨afener and Hamann, 2005). The terminal velocity v∞is often linked to the escape velocity at the stellar surface (adjusted to include the effects of radiation pressure) vesc. The review by Kudritzki and Puls (2000) uses  
2.65 Teff > 21 000K 1.4 10 000K < Teff < 21 000K 1.0 Teff < 10 000K (A2)
v∞= C(Teff)vesc, C(Teff) =
 Hot star winds can reach speeds of several thousand kilometers per second and produce significant mass loss values.
Observationally-determined formulae governing the winds of hot stars can be found in Vink et al (2001) for O and B stars, and Nugis and Lamers (2000) for Wolf–Rayet stars.
Cool stars (FGKM)

Low-mass stars form stellar coronae, an outer layer of high-temperature matter in the megakelvin range (Parnell and De Moortel, 2012). The stellar winds of low-mass stars are driven by the high thermal pressure arising from the high temperature of their coronae. Stellar surface (i.e. photospheric) magnetic fields, and their continuation into the magnetised corona and wind regions, play a crucial role in accelerating and structuring the winds. Models of cool star winds are based on (magneto-)hydrodynamic models and the conservation of mass and angular momentum, the ideal gas law, and the polytropic equation of state. Under these conditions the one-dimensional velocity profile is independent of the wind mass loss rate ˙M, and given by Parker (1958, 1960) �2c2
� , (A3)
dv dr = v v2 −c2s
s r2
r −GM
where cs is the sound speed for a polytropic gas. For a more detailed description of cool star winds, see Vidotto (2021). For an a constant-temperature (isothermal) wind, cs is also constant, and then (A3) is the well-known isothermal Parker (1958) wind model. In this isothermal case the wind profile can expressed analytically via the the Lambert W function (see Cranmer, 2004).

For the non-isothermal case there is no closed-form solution to Equation (A3); instead the equation must be solved numerically. Polytropic wind models, such as the Parker model and its extensions, are applied to an ‘already heated’ corona, i.e. these models are separated from the problem of coronal heating. For the Sun and other low mass stars, coronal-heating mechanisms are not well understood. An extended model that includes coronal heating is described by Gombosi et al (2018). In modelling based on the solar wind, velocities of coronal winds generally range from 300–800 km s−1 with younger, more rapidly rotating stars having hotter coronae and faster winds. The mass loss rates can range from 10−14 −10−12 M⊙yr−1. Coronal winds are composed of a mix of ions, electrons, and neutral atoms, primarily hydrogen and helium.
Waldron (1985) and de Jager et al (1988) have derived empirical relations for mass loss rates across the HR diagram. These have been combined with the aforementioned rates for hot stars by Glebbeek et al (2009), and subsequently in the well-known Modules for Experimentation in Stellar Astrophysics (MESA) code (Paxton et al, 2011, 2013, 2015)16. It should be noted that this approach is not based on a physical model of the wind acceleration but rather on empirical relations.
Stars that have aged off the main sequence

Low to intermediate-mass M ≲8M⊙stars evolve into asymptotic giant branch (AGB) stars, while massive stars M ≳10M⊙evolve into red supergiant stars (RSG). The winds of evolved stars are driven by the combination of stellar pulsations and radiation pressure on dust grains, which form in these stars’ cool, extended atmospheres. Stellar pulsations produce shock waves that lift material away from the stars’ surface. Dust grains, which form in the cool outer regions of the atmospheres, are pushed outwards by radiation pressure, dragging gaseous matter with them (see the reviews by H¨ofner and Olofsson 2018; Decin 2021). For evolved stars the wind velocities are low, ranging from 10–50 km s−1, but mass loss rates are high (10−7 −10−4 M⊙yr−1). The winds of these evolved stars are rich in molecules and dust grains with primarily carbon- and oxygen-rich compounds. The wind velocity profile for evolved stars is also often assumed to be that of a β-law as for massive stars (Equation A1). The mass loss rate is different from massive stars, however. For red supergiant stars, Reimers (1975) gave a semi-empirical law that describes the mass loss rate of a star as a function of its luminosity and radius:
˙M = η � L �
�� R �
�� M �
L⊙
R⊙
M⊙ �−1
(A4)
16In MESA this is referred to as the ‘Dutch’ wind mass loss prescription.

where η = 4 × 10−13 M⊙yr; this has become known as the ‘Reimers relation’. A different version from Schr¨oder and Cuntz (2005) is
�3.5 �
�
� (A5)
Teff
g⊙
˙M = η � L �
�� R �
�� M �
L⊙
R⊙
M⊙ �−1
1 +
4000 K
4300g
where η = 8 × 10−14 M⊙yr
Comets, asteroids, and other small bodies are thought to be remnants of the original planetesimal population of the Solar System. As such, their physical, chemical, and isotopic properties hold crucial details on how and where they formed and how they evolved. Yet, placing precise constraints on the formation region of these bodies has been challenging. Data from spacecraft missions have a particularly high potential of addressing the question of the origin of the visited bodies. ESA’s Rosetta mission to comet 67P/Churyumov-Gerasimenko returned data from the comet for two years on its journey around the Sun. This extensive data set has revolutionised our view on comets and still holds unsolved problems. Here, we aim to determine comet 67P’s bulk elemental composition from Rosetta data, including its refractory-to-ice ratio. We use these results to constrain the temperature in the protoplanetary disk where comets formed and, using a disk model, the formation location. We use the Rosetta/ROSINA (Rosetta Orbiter Spectrometer for Ion and Neutral Analysis) measurement of the volatile/ice composition and the Rosetta/COSIMA (COmetary Secondary Ion Mass Analyzer) measurements of the refractory composition of comet 67P. These measurements are combined using a Monte Carlo method. The refractory-to-ice ratio is a free parameter that is constrained a posteriori. Using only the composition, we constrain the refractory-to-ice ratio to 0.5 < χ < 1.7, and derive the bulk elemental abundances for 67P of H, C, N, O, Na, Mg, Al, S, K, Ar, Ca, Cr, Mn, Fe, Kr, and Xe. We find the noble gas xenon in near solar elemental abundance in comet 67P. Krypton is slightly depleted, while argon is heavily depleted. Comet 67P is enriched in all three noble gases by up to 2.5 orders of magnitude compared to CI chondrites. We show this is consistent with a formation region between 25 and 35 au in a protoplanetary disk region with temperatures between 30 and 40 K and with the trapping of dust for a long time in rings of the protoplanetary disk.