Cometary nuclei and tails are affected by stellar winds. They are produced when material from the outermost layer of a star is subject to significant acceleration, overcoming the hold of the star’s gravity. Stellar winds are frequently characterised by their acceleration mechanism (Lamers and Cassinelli, 1999; Vidotto et al, 2023; Owocki, 2023; Vink, 2026).
Massive stars have radiation-driven winds, where the star’s luminosity is so great that the radiation pressure blows material outwards. This includes O0–B5 stars, with the winds of B3–B5 stars being somewhat weaker. Stars with stellar types B6–A9 do not exhibit significant winds (this includes the exocomet host β Pictoris, which is an A6V star). FGKM-dwarfs exhibit coronal winds, powered by the high temperatures (in the megakelvin range) in the coronae of these stars. The stellar types F0–F4 exhibit weak coronal winds due to their narrow convection zones. This category of winds includes the solar wind, as the Sun is a G2V star. Finally, stars that have aged beyond the main sequence have winds driven by dust and pulsations. This information is summarised in Table 1.

The composition of stellar winds generally mirrors the composition of the outermost stellar layer but is modified by factors including ionization processes, mixing and enrichment from the star’s interior, elemental fractionation (e.g., first ionisation potential effect, Laming, 2015), and dust formation (in cool stars).
Winds can be described in terms of density ρ, velocity v, magnetic field B and temperature T (e.g. Vidotto, 2021). The mass flux of the wind through a closed surface S containing the star ˙M is a function of the wind density ρ and velocity v, ˙M = �
ρv · dS. (4)
S
In a situation of spherical symmetry the flux simplifies to ˙M(r) = 4πr2ρv where r is the distance to the stellar centre. For a steady wind there can be no accumulation of matter in the wind so the flux must be independent of r. This imposes a constraint on the density and velocity of the wind ˙M = 4πr2ρ(r)v(r) = constant in r. (5) In general the winds accelerate away from the stellar surface and approach a terminal velocity v∞at large distances from the star. Different wind acceleration mechanisms are possible and yield different velocity and density profiles v(r) and ρ(r). As it expands, the wind emanating from a star generally achieves supersonic speeds. The velocity profile v(r) of evolved stars and high-mass stars is often described by a “β law” where v(r) ∝(1 −R⋆/r)β where R⋆is the stellar radius (see Sections A.1 and A.4). The velocity profile of the solar wind is described by the Parker (1958) wind model (see Section A.3) and a host of extensions and modifications. For the Parker (1958) wind model the wind speed at large distances grows as ln r −ln(ln r), and thus never reaches a terminal velocity.
Typical ranges of mass loss rates, wind speeds, and pressures around different stellar types are tabulated in Table 2. Additional information on the dependence of stellar winds on host star mass, including a compilation of empirical formulae for the mass loss rates, can be found in the Appendix. On the Main Sequence the main influence of the stellar wind is on the formation of the cometary tail from outgassed material, with some influence on the surface as discussed in Section 7. Once stars leave the Main Sequence, the strong winds significantly change the mass of the star itself, with important implications for all orbiting bodies, as discussed next.