Hydrogen Hot Ion Precipitation in the Martian Ionosphere #P13B-1317 Christopher D. Parkinson 1, Michael Liemohn 1, Xiaohua Fang 2 1 AOSS Dept., University of Michigan, Ann Arbor, MI; 2 LASP, University of Colorado, Boulder, CO I. ABSTRACT High energy H/H+ ion precipitation into Mars' upper atmosphere is modeled and discussed. Such particle transport has previously been modeled for Earth and we have extended this work for the Martian ionosphere using different cross sections for relevant Martian "background" species. Atmospheric effects of these precipitating hot ions in the Martian atmosphere are studied and reported on. Solar wind protons as well as pick-up ions from the planetary exosphere routinely enter and alter the upper atmosphere. Atmospheric effects of these precipitating hot ions in the Martian atmosphere will be described. In addition, the influence of the Martian remnant crustal fields on particle entry will be examined. A study of the ionization, excitation, and energy deposition is conducted. The result is a robust examination of the influence of energetic ion transport on the Mars upper atmosphere. II. MODEL DESCRIPTION  The 3-D Monte Carlo test particle precipitation model [Parkinson et al., 2008; Fang et al., 2004]: IV. RESULTS V. SUMMARY AND CONCLUSION  SWIM model atmosphere [Parkinson et al., 2008]: III. DISCUSSION (1) Higher energy particles have less energy deposition at high altitudes, but more energy deposition at lower altitudes. (2) The dip angle of the assumed magnetic field lines has only a small influence on the fluxes and energy deposition rates, with the low dip angles (more horizontal field lines) resulting in slightly more deposition at high altitudes and a bit less at low altitudes. (3) A nightside atmospheric profile results in substantially less deposition at high altitudes, but nearly the same deposition at low altitudes. (4) There are 4 distinct layers for the precipitating particles, as evidenced in the integrated flux plots. There is a high-altitude region of few collisions, a transition layer below this where the ion fluxes are depleted and the fast neutral fluxes increase, a quasi-equilibrium zone with relatively constant fluxes, and finally a low- altitude depletion region where the fluxes rapidly drop to zero. (5) The secondary high-altitude peak in some quantities is an artifact of the isotropic pitch angle distribution for the injected energetic ions. (6) The peak of ionization from precipitating solar wind ions is found to be typically around 120 km which is lower than the typical EUV ionization peak which is found above 130 km altitude. (7) Mars has a more complicated scenario for energetic particle precipitation than does the Earth because of the localized nature of the crustal field anomalies. Reference for the hydrogen hot ion precipitation model: Parkinson, C. D., M. W. Liemohn, X. Fang, Hydrogen Hot Ion Precipitation in the Martian Ionosphere, J. Geophys. Res., SWIM Special Issue, in press,  Discussion  Features to Notice Figure 1 a, b: SWIM model atmosphere noon and midnight.Figure 2: Ray diagram showing slant path geometries (dip angles) for a point source crustal field anomaly Our approach uses a three-dimensional Monte Carlo test particle trajectory calculation to simulate the motion of precipitating ions through the Mars upper atmosphere. It follows the guiding center motion, their conversion to and from an energetic neutral hydrogen atom, and their eventual energy deposition to the ionosphere and thermosphere. To handle the 3-D scattering of precipitating particles, the Monte Carlo model used for this work has been previously described for an Earth-specific model by Fang et al. [2004, 2005]. This code was written independently in 3-D, but has been extensively compared with the 1-D model of Solomon [2001]. The Monte Carlo model monitors the trajectories of incident energetic particles in a collision-by- collision manner down to an assigned low-energy cutoff limit. A variety of effects due to inelastic and elastic collisions are accumulated over the course of the particle traveling in a planetary upper atmosphere. Note that this is the initial application of this previously Earth-specific model (Fang et al., 2004) to the Mars space environment. The code launches particles from a single location on the upper boundary, and then calculates the 3-D spread in the atmosphere. Multiple point source results over a latitude-longitude grid and the results are summed to create a 3-D total precipitation result. This is what was done for the data-model comparisons presented in Fang et al. [2004] and for the auroral arc spreading simulations of Fang et al. [2005]. For the Mars precipitation calculations, the precipitation is calculated on a coarse grid in longitude and latitude over the day- and nightside thermosphere with sensitivity studies of key parameters performed, viz., dip angle, gyration radius, and injection energy and day/night atmospheric conditions. For our calculations, we have adopted a standard reference 1-D profile based on the Solar Wind Interaction with Mars (SWIM) Competition model atmosphere [S. Bougher, private communication] going from 100 km to approximately km. Key species include O, O 2, and N 2 in a background atmosphere of CO 2, as illustrated in Figures 1a and b, corresponding to Martian noon and midnight (solar and anti-solar), respectively. All species are more abundant with altitude for the noon model atmosphere, with the exception of O, which remains similar for both noon and midnight. Figure 4 a, b, c Figure 3 a, b Figure 5 a, b, cFigure 6 a, b, cFigure 7 a, b, c Figure 10 shows the integrated proton (p = H+) and hydrogen fluxes for Martian noon and standard reference case parameteric values. Figures 11 and 12 show the integrated proton (p = H+) and hydrogen fluxes for Martian noon over the range of injection energies and dip angles considered, respectively. Figure 13 shows a similar plot, only for a noon/midnight comparison. We see from the view of particle precipitation that the atmosphere can be divided into 4 regions, viz., (1) >190 km (collisionless): downward H+ fluxes are larger than downward H fluxes; (2) km (transition layer): downward H+ and H fluxes are comparable; (3) km (quasi-equilibrium zone): the ratio of H+ and H fluxes stays almost invariant. (4) <110 km (depletion zone): the precipitating particle fluxes drop rapidly to zero. Our code additionally allows us to determine trajectories of the particles along the magnetic field lines through our model atmospheres. Mars has a more complicated scenario for energetic particle precipitation than does the Earth because of the localized nature of the crustal field anomalies. Multiple dip angles feed into an isolated magnetic anomaly, and several nearby anomalies further complicate the magnetic structure. Our results for different dip angles (30-90°) were not significantly different and so the precipitating particles can most likely enter the thermosphere from a broad locus of points around cusp features. This is different for Earth because the Earth has such a huge internal dipole field that its cusp region is highly localized with nearly vertical fieldlines. What we have seen generally is that higher energy precipitating particles penetrate deeper into the atmosphere than lower energy precipitating particles, which is expected, and the rates due to ionization are much larger than those for excitation. Excitation and ionization rates are much lower in the upper atmosphere for midnight since atmospheric density is much lower (cf. Figure 1 a and b, the SWIM noon and midnight model atmospheres). As noted previously, this study presents the initial results from the modification of our precipitation code for Earth for use with the Mars scenario. This work will be extended further to include O+/O precipitation with the view to keeping track of secondary hot ions/neutrals, which is important for sputtering in the Martian upper atmosphere. A further development is the removal of the guiding center approximation, to more accurately resolve the influence of crustal fields on large gyroradius particles. This modification will also allow the inclusion of a nonuniform magnetic field. In addition, the model can be applied to the Venus case with relative ease, because the major upper atmospheric constituents are the same. Figure 2 shows the geometry of our slant path calculations through the Martian atmosphere. Each ray represents a different dip angle and slant path to a point source crustal magnetic field anomaly from a point at the top of the atmosphere. Figures 3 a and b show the ionization and excitation rates, respectively, as a function of altitude for our standard reference parameters, viz., dip angle of 90° (vertical field lines), gyration radius of 3 km, injection energy of 1 keV, and the noon model atmosphere. Our results for Mars are normalized at the top to a tenth of that for the Earth, viz., 0.1 erg/cm 2 and 10 5 test particles. This is analogous to typical solar wind ions entering the atmosphere directly above a typical strong crustal field source region. Figures 4 a, b and c and Figures 5 a, b and c illustrate how the ionization and excitation rates respectively vary with altitude for the injection energies 0.5 keV, 1 keV, and 4 keV (corresponding to ~ 300, 400 and 800 km/s for the noon and midnight model atmospheres) for the species CO 2, O, and N 2, respectively. We see that lower energy particles do not penetrate as deeply into the atmosphere and have much more of their ionization and excitation happening in the upper part of the atmosphere as compared to higher energy injection particles. Figures 6 a, b and c and Figures 7 a, b and c respectively show how the ionization and excitation rates vari with altitude as a function of changing dip angle. Above the peak value, the excitation/ionization rate is somewhat lower at a given altitude for increasing dip angle, asymptotically approaching the value at 90 degrees, and the opposite below the peak rate at ~ km depending on energy. This is true for either the noon or midnight cases. This means that a higher dip angle implies path length through atmosphere is shorter and precipitating particles can penetrate and deposit more of their energy deeper in the atmosphere. An interesting feature that is particularly noticeable in the oxygen ionization rate profiles is a double peak, with a main maximum near 120 km altitude and a second relative maximum near 200 km altitude. This feature was explained by Fang et al. [2004] to be the result of the topside incident ion injection having an isotropic pitch angle distribution. Figures 8 and 9 show how the ionization and excitation rates vary with altitude for both the noon and midnight model atmosphere cases, respectively. For each scenario, the rates fall several orders of magnitude over the range of altitude considered. Excitation and ionization rates are much lower in the upper atmosphere for midnight since atmospheric density is much lower (cf. Figures 1 a and b). Clearly, hot hydrogen ions are precipitating to approximately to 100 km where the excitation and ionization rates are falling rapidly to zero and to altitude levels lower than EUV photons can penetrate (~130 km). Peak rate values are occurring at approximately 120 km. Hence, they (and their associated phenomena) should be more readily detectable below 130 km where EUV photons no longer dominate. Figures 8, 9, and 10Figures 11, 12, and 13