magnetosheath, but given the uncertainty associated to these values, we use here as a lower limit the fast magnetosonic speed in the solar wind, where, VA,i = BSW/√4πminSW is the Alfv´en speed, Vs,i = � γTi,SW/mi is the ion sound speed, and γ = 2 is the adiabatic index of the plasma. The acceleration of solar wind protons is due to magnetic reconnection in the magnetosphere. For a northward IMF, magnetic reconnection is 8 weakly driven at the lobes at high latitudes.

For a southward IMF, magnetic reconnection is more strongly driven both at the nose and in the tail of the magnetosphere. Magnetic reconnection converts part of the magnetic energy stored in the planetary magnetic field configuration, corresponding to roughly 10 keV (shown by the green vertical solid line in Fig. 7(a)-(b)), to kinetic energy of the particles. The efficiency of this conversion is on the order of 10% (Shay et al. 2014; Phan et al. 2014; Haggerty et al. 2015), meaning that roughly 1 keV of magnetic field energy is converted to kinetic energy of the particles. These two competing processes (bow shock slowdown and acceleration by magnetic reconnection) are responsible for the proton fluences versus energy at the surface reported in Tab. 1. Substantial proton fluences below 500 eV –mostly at high latitudes– are a consequence of the plasma slowdown downstream of the bow shock. While, proton fluences above 1.5 keV –mostly at high latitudes (under northward IMF) or at low latitudes (under southward IMF)– are a consequence of magnetic reconnection and proton heating in the magnetosheath. This result has implications for current estimates of ion sputtering at Mercury that consider ions with a fixed energy of 1 keV amu−1 (corresponding to the solar wind proton energy upstream of the bow shock). A broader discussion on the impact of proton energy distributions on sputtering yield can be found in Sec. 4.2. For electrons, the net effect of Mercury’s magnetic field is energization by magnetic reconnection. Solar wind electrons have a Maxwellian energy distribution with temperature Te,SW=21.5 eV, as shown by the orange curves in Fig. 7(c)-(d). For electrons, the flow energy component meV 2 SW/2=8.26 eV (in the solar wind) is negligible with respect to their thermal energy component, both outside and inside Mercury’s magnetosphere. At the surface, the electron energy distribution displays a high-energy tail extending up to ∼10 keV, as shown by the blue curves in Fig. 7(c)-(d). Magnetic reconnection energizes a substantial fraction of solar wind electrons up to the maximum magnetic energy available, shown by the green vertical solid line in Fig. 7(c)-(d). As a consequence, precipitating electrons are mostly found in the energy range 100-4000 eV, as summarized by the fluences in Tab. 1.
4. DISCUSSION AND IMPLICATIONS The study of solar wind particle precipitation on weakly magnetized bodies (such as Mercury) is more complex than that on unmagnetized bodies (such as the Moon). The presence of an intrinsic magnetic field – although weak – has a considerable effect on both where particles precipitate onto the surface and what energy these particles have. In the following, we discuss these two effects separately and their implications for plasma-surface interaction processes. Particular relevance is given to the comparison between weakly magnetized and unmagnetized bodies.
4.1. Magnetosphere as a filter: the effect of a magnetic field on the spatial distribution of particles at the surface We consider here the interaction between an homogeneous, constant flow of plasma (i.e., an ideal solar wind) and a spherical body (i.e., an ideal rocky body such as Mercury or the Moon). The body rotates with angular velocity ω and spin axis perpendicular to the plasma flow. If the body is unmagnetized, the flux of particles reaching the surface integrated over a time ∆T ≫2π/ω corresponds to the solar wind fluence multiplied by the cross-sectional area πR2 and divided by the total surface body area 4πR2:
nSWVSW
4 ∆T (1) . If the body is weakly magnetized, the fluence of particles at the surface is further reduced by a factor α > 1 as follows:
nSWVSW
4α ∆T (2) . The effective shielding parameter α accounts for the average reduction of plasma flux from the solar wind onto the surface due to the magnetic field. From the averaged fluences reported in Fig. 4 for protons and Fig. 6 for electrons, we derive the effective shielding parameters for solar wind protons and electrons:
αp ≈20, (3) αe ≈10. (4) Our results indicate that the magnetic field effectively shields about 90% of the incoming solar wind particles, reducing the fluence at the surface by one order of magnitude as compared to the unmagnetized case. The shielding values do not vary appreciably between our two runs using a purely northward or a purely southward IMF, suggesting that the 9 IMF direction weakly affects the total number of particles precipitating onto the surface.