On the properties of O + and O 2 + ions in a hybrid model and in Mars Express IMA/ASPERA-3 data: A case study

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1 Planetary and Space Science 56 (2008) On the properties of O and O 2 ions in a hybrid model and in Mars Express IMA/ASPERA-3 data: A case study E. Kallio a,, A. Fedorov b, E. Budnik b, S. Barabash c, R. Jarvinen a, P. Janhunen a a Finnish Meteorological Institute, Erik Palmenin aukio 1, Helsinki, Finland b CESR-CNRS, Toulouse, France c Swedish Institute of Space Physics, Kiruna, Sweden Received 12 January 2008; received in revised form 16 March 2008; accepted 26 March 2008 Available online 7 April 2008 Abstract We have performed a numerical simulation to analyze the energy spectra of escaping planetary O and O 2 ions at Mars. The simulated time energy spectrograms were generated along orbit no. 555 (June 27, 2004) of Mars Express when its Ion Mass Analyzer (IMA)/ASPERA-3 ion instrument detected escaping planetary ions. The simulated time energy spectrograms are in general agreement with the hypothesis that planetary O and O 2 ions far from Mars are accelerated by the convective electric field. The HYB-Mars hybrid model simulation also shows that O ions originating from the ionized hot oxygen corona result in a high-energy (E41 kev) O ion population that exists very close to Mars. In addition, the simulation also results in a low-energy (Eo0.1 kev) planetary ion population near the pericenter. In the analyzed orbit, IMA did not observe a clear high-energy planetary ion or a clear low-energy planetary ion population near Mars. One possible source for this discrepancy may be the Martian magnetic crustal anomalies because MEX passed over a strong crustal field region near the pericenter, but the hybrid model does not include the magnetic crustal anomalies. r 2008 Elsevier Ltd. All rights reserved. Keywords: Mars; Mars solar wind interaction; Ion escape; Ion acceleration; Solar wind; Numerical simulation 1. Introduction The ASPERA-3 instrument onboard Mars Express spacecraft has provided the largest data set so far on the properties of the fast escaping planetary ions at Mars by its Ion Mass Analyzer (IMA) sensor (see, for example, Fedorov et al., 2006). The IMA instrument can distinguish solar wind ions (H,He ) from heavy planetary ions (O,O 2 and heavier ions). The IMA data set is collected during various solar activity and solar wind conditions providing a possibility to study the response of the ion escape processes at various parameters and, consequently, to identify planetary ion acceleration processes. Another approach to studying the escape and acceleration of planetary ions is to compare the IMA observations with a numerical self-consistent model. In such a way, one has a possibility to study which Corresponding author. Tel.: address: Esa.Kallio@fmi.fi (E. Kallio). observed features can be reproduced by the employed model, which is always a simplified and limited description of reality. In other words, typical features seen in the data may be distinguished from untypical features that cannot be understood with the simulation model. Various global models have been used to study the plasma environment and the properties of planetary ions near Mars: test particle simulations, gas dynamic models, magnetohydrodynamic models (e.g. Ma et al., 2004; Harnett and Winglee, 2005) and quasi-neutral hybrid models, or hybrid models for briefly (e.g. Brecht and Ferrante, 1991; Shimazu, 1999). We extend the pilot study (Kallio et al., 2006a), where a hybrid model was compared with IMA observations made in June 27, 2004 (orbit no. 555). In this paper we, instead, generate, for the first time, simulated IMA type of time energy spectrograms for O and O 2 ions along the orbit of MEX by using a hybrid model. The paper is organized as follows. We first briefly describe the employed hybrid model, putting emphasis to /$ - see front matter r 2008 Elsevier Ltd. All rights reserved. doi: /j.pss

2 E. Kallio et al. / Planetary and Space Science 56 (2008) introducing the method that was used to generate IMA type time energy spectrograms. Then we present the macroscopic parameters along the orbit of MEX based on the hybrid model. After that we introduce the simulated O and O 2 time energy spectrograms and compare them with the measured ones. Finally, possible sources and implications for the found similarities and differences between the simulation and the data are discussed, especially the possible role of the Martian crustal magnetic anomalies. 2. HYB-Mars hybrid model and IMA instrument The properties of the quasi-neutral HYB-Mars model used in this paper have been discussed thoroughly in our previous works (see Kallio et al., 2006a, b). Here we present only the basic properties of the model and new features focusing on a new tool for generating simulated energy spectra. In the hybrid model, ions are treated as particles while electrons form a massless charge neutralizing fluid. The analyzed run contains four ion species, H,He,O and O 2. There are two sources of H ions (the solar wind and the hot hydrogen corona), one source of He ions (the solar wind), two sources of O ions (the hot oxygen corona and emission from the exobase) and one source of O 2 ions (emission from the exobase). The total H and O photoion production rates from the spherically symmetric neutral corona are and s 1, respectively. The total ion loss rate from Mars of O and O 2 ions emitted from the exobase are and s 1, respectively, by assuming that the ion emission flux from the exobase has the following solar zenith angle (SZA) dependence: Ion flux cos(sza)0.1. Here the cosine term was used only on the dayside and the ion emission flux on the nightside was therefore a constant. The size of the simulation box is 4.2R M ox, y, zo4.2r M (R M =3393 km, the radius of Mars). The coordinate system is Mars-centered Solar Orbital (MSO) where x-axis is pointing from Mars to Sun, z is perpendicular to the orbital plane of Mars pointing to the North and y completes the right-handed coordinate system. The simulation contains three grid sizes above Mars: the nominal grid size is 720 km (0.2R M ), the 360 km (0.1R M ) grid size is used at 2.1R M oro3.2r M (r=(x 2 y 2 z 2 )) 1/2 ) and the smallest grid size 180 km (0.05R M )atro2.1r M (see Kallio et al. (2006b), Fig. 1 for the plot of the used grid structure). The total number of ions in the simulation box is on average about 11.4 million, the average number of particles per cell is 30 and the time step, dt, is 0.02 s. The upstream plasma parameters are identical with our previous study for consistency: the density n(h ) ¼ 3cm 3 and the bulk velocity U(H ) ¼ [U x, U y, U z ] ¼ [ 450, 0, 0] km s 1 (for more plots about the plasma and magnetic parameters in the analyzed upstream case see Kallio et al., 2006a). The solar wind contains also 4% of He ions (n(he ) ¼ 0.04 n(h )) that have the same velocity as the solar wind H ions. The interplanetary magnetic field (IMF) is B sw ¼ [B x, B y, B z ] ¼ [cos(551), cos( 601) sin(551), sin( 601) sin(551)] 1.12 nt[0.64, 0.46, 0.79] nt. It is noteworthy that there is no magnetometer onboard MEX and that the direction of the IMF on the YZ plane is estimated from Mars Global Surveyor (MGS) magnetic field measurement (see Kallio et al. (2006a), for details of the used method and Brain et al. (2006), for the discussion about inaccuracies to determine the direction of the IMF from the MGS observations). The crustal magnetic field found by MGS (Acun a et al., 1998) is not included in the model. Inclusion of the crustal magnetic anomalies is not a technical or a physical problem and, indeed, the authors have made low-resolution test runs which include also the contribution of the crustal magnetic field. However, a low spatial resolution run which models accurately small-scale magnetic anomaly features and which includes strong magnetic field regions near Mars is computationally extremely expensive to perform. Such a run requires massively parallelized supercomputer environment and is a challenge for forthcoming simulations. The main physical improvements of the used HYB-Mars hybrid model compared with the previously used hybrid Mars models are the following. The model includes H O-HO and H H-HH charge exchange processes. The hybrid model does not include a selfconsistent ionosphere, but in the new hybrid model the electrodynamics near the exobase is improved by introducing a thermal plasma layer at the inner boundary that mimics thermal plasma in the upper ionosphere where ions move slowly a small distance during the simulated time interval. The thermal plasma layer is implemented by generating ions at the obstacle boundary whose positions are kept unchanged during the simulation. In the simulation the thermal plasma layer also reduces statistical fluctuations in the plasma density which could result in fluctuations to the other plasma parameters, for example, to the magnetic field. The population of O ions is now also recorded, making it possible to distinguish O ions having originated from the hot oxygen corona and O ions having originated from the exobase. The model includes also new tools for a detailed comparison between the simulation and the ion observations. In the hybrid simulation, ions are modeled as particles which are accelerated by the Lorentz force and no assumptions about the ion velocity distribution function or the correlation between different ion species are needed. The ion distributions can therefore be fully three-dimensional (3D) and different ion species can have different distributions and thereby also different bulk velocities and temperatures. Existence of individual ions makes it possible to implement a virtual detector in the simulation where the velocity and position of an ion are recorded when it crosses the detector surface. One example of a virtual detector is a plane detector where the ion is recorded when it moves through a plane. For example, x ¼ constant plane

3 1206 ARTICLE IN PRESS E. Kallio et al. / Planetary and Space Science 56 (2008) Fig. 1. A virtual tube detector used in the hybrid model. Panels (a) (c) represent the orbit of MEX in the analyzed orbit no. 555 (June 27, 2004) as a 3D plot viewed along the x-, y- and z-axis, respectively. The dots are examples of O 2 ions that were collected on a virtual tube detector and which were used to generate simulated energy spectra in this paper. In panel (d) the same points are presented in cylindrical coordinates. The text labels O 2 show the approximate position of the regions where high flux of escaping planetary ions was recorded by IMA/ASPERA-3/MEX instrument. See text for details. makes it possible to study the spatial distribution of escaping planetary ions at Mars (see Kallio et al., 2006b). In this paper, instead, a new type of virtual detector, a tube detector, is implemented into the simulation code, enabling the generation of an ion energy spectrum along the orbit of MEX. The position and velocity of an ion are recorded by the tube detector if an ion comes close enough to the orbit of MEX, in our paper closer than 0.05R M. The simulated energy spectra presented in the paper are generated from ions that are collected from t ¼ 400 to 450 s from the start of the simulation by the tube detector. Fig. 1 shows orbit no. 555 in June 27, 2004, and an example of O 2 ions that are collected by the tube detector. Orbit no. 555 was chosen because it is one of the first clear examples of the escaping O and O 2 ions seen in IMA/ASPERA-3 data and the first event that the authors started to interpret by the hybrid model introduced in this paper. It also provides a typical example of IMA observations in the wake region (Fedorov et al., 2006). The advantage to analyze orbit no. 555 in the present study is that the basic properties of the plasma and the magnetic field in the analyzed run are well documented and presented earlier (see Kallio et al., 2006a). The virtual detector data are used to make IMA-type time energy spectrograms in the time region 02:20 04:40 UT at orbit no The collected ions are binned to IMA-type energy intervals [E i, E (i1) ] where E i is 3 ev (10.08) i and i ¼ 0,1,y,95. If the energy of an ion is below the smallest

4 E. Kallio et al. / Planetary and Space Science 56 (2008) energy interval then the ion is included into the lowest energy interval. In this paper all collected ions are used in order to obtain good statistics in the simulated energy spectra. The used field of view (FOV) of the tube detector is therefore 4p, i.e., much larger than the FOV of the IMA which is about per detector. Here 51 is the FOV width in the polar dimension. The IMA instrument has 16 such detectors scanning 7401 along polar angle and, consequently, the total FOV is (for details of the IMA instrument, see Barabash et al., 2004). In the analyzed time period IMA was in the electrostatic elevation scan mode. Two vectors in Fig. 1a show the direction of the IMF and the convective electric field (E sw ¼ U sw B sw, where U sw is the velocity of the solar wind and B sw is the IMF) viewed along the x-axis. Fig. 1 illustrates that MEX was approaching Mars on the hemisphere where E sw is pointing away from the planet. MEX is therefore anticipated to approach Mars from the hemisphere, usually referred to as the E sw hemisphere, where the flux of accelerated planetary ions is high (for a detailed discussion of the asymmetry of the properties of the escaping ions with respect to the direction of the E sw see Kallio et al., 2006a). It was in that hemisphere where IMA detected an intense flux of escaping planetary ions at 03:15 03:25 UT (see Kallio et al., 2006a, Fig. 1 for details). In Fig. 1a d the position of the orbit where the high flux of planetary ions is detected are marked as text O 2. It is finally worth noting that in the analyzed orbit the pericenter was on the nightside near the terminator plane (x ¼ 0). 3. Results 3.1. Macroscopic parameters As the first step in this case study, we will look at macroscopic plasma parameters along orbit no. 555 based on the hybrid model. Figs. 2a and b give the density and the bulk velocity of H,He,O and O 2 ions. It can be seen that the solar wind density slightly decreases while the density of planetary ions substantially increases when MEX approaches the pericenter. The particle density of Fig. 2. The macroscopic plasma and field parameters derived along orbit no. 555 at 02:20 04:40 UT. The panels from top to bottom are: (a) particle density, (b) bulk velocity and (c) magnetic field from the hybrid model in MSO coordinates. The vertical transparent shaded region at 03:30 04:10 UT is the region which is analyzed later in Fig. 4a.

5 1208 ARTICLE IN PRESS E. Kallio et al. / Planetary and Space Science 56 (2008) planetary ions is higher than the density of the solar wind ions at 03:40 04:10 UT. Recalling that an O ion is 16 times and an O 2 ion 32 times more massive than an H ion, the mass density is dominated by planetary ions for even a longer time interval. It is worth noting that there are no O 2 ions after 04:10 UT in contrast to O ions. This difference is associated with the hot oxygen corona which is a source of O ions forming a low-density O ion background. This difference between the molecular and atomic oxygen ions will be studied in more detail in the next section by analyzing the simulated time energy spectrograms. The velocity of all ion species decreases when MEX is approaching the pericenter, being only few tens of km s 1 between 03:40 and 04:00 UT. The high-velocity spikes seen in U(He ) may be just statistical fluctuations resulting from a small number of He ions in that region. Note that if MEX is near the pericenter where the grid size is 180 km and if its velocity along the coordinate axis is, say 4kms 1, it takes 45 s ( ¼ 180 km/4 km s 1 ) for MEX to move across a grid cell. The simulation time resolution becomes poorer far away from Mars because the grid size increases and the velocity of MEX decreases: If the grid size is 720 km and if its velocity along the coordinate axis is 2kms 1, the crossing of MEX across a grid cell along the coordinate axis takes 6 min. We note that the macroscopic parameters cannot be accurately determined from the IMA measurements between 03:20 and 04:20 UT, partly because the count rate becomes low and partly because IMA cannot measure all ion species in its full energy range at a single data mode. Moreover, the IMA instrument can cover only about 2p space angle in velocity space while the simulated macroscopic plasma parameters include contribution of ions from the full 4p velocity space Simulated O and O 2 time energy spectrograms In this paper the time energy spectrograms were derived along the orbit of MEX by using a hybrid model for the first time. The simulated O 2 and O time energy spectrograms can be seen in Fig. 3 in panels (a) and (b), respectively. In these panels, three different planetary ion populations are distinguished: (1) The E SW population, (2) the E low population and (3) the O corona population. The E SW population in the O 2 energy spectra forms a narrow energy range ion population when MEX is approaching Mars from 02:20 to 03:40 UT (Fig. 3a). The energy of the planetary ions from 03:40 to 02:20 UT increases with increasing altitude due to the convective electric field E( ¼ U e B, where U e is the electron bulk velocity and B is the magnetic field). This increase of energy is consistent with previous hybrid model simulations (e.g. Kallio et al., 2006b). In the E low population region, low-energy (Eo0.1 kev) O and O 2 ions can be found near the pericenter and near the terminator plane at x ¼ 0 (Fig. 3a and b). In the region labeled as the population, the high-energy (E41 kev) O ions O corona are recorded during the whole analyzed time period, especially near and after the pericenter (Fig. 3b). Note that such population does not exist in the simulated O 2 ion time energy spectrogram (Fig. 3a). The difference between O and O 2 ions is associated with the hot oxygen corona which is a source for O ions but no such source exists for O 2 ions. The O and O 2 ions in the E SW population region and in the E low population region are originating from the ionosphere. The ionospheric ions are emitted to the simulation box through the model exobase. This was confirmed by making energy spectra separately for O and O 2 ions: O ions emitted from the exobase form a time energy spectrogram which is similar to Fig. 2a ando ions originating from the oxygen corona result in the O corona population (figures not shown). The O ions from the neutral corona are spread over a much wider energy range than the E SW population ions because these ions can originate from various locations from the neutral oxygen corona and, consequently, could have been the subject of different acceleration by the convective electric field Measured time energy spectrograms The observed IMA time energy spectrogram is given in Fig. 3c and d. It can be seen that in the beginning MEX is in the solar wind and that it crosses the bow shock at 02:50 UT. Note that the time energy spectrogram in the solar wind in Fig. 3c contain two peaks because the spectrogram includes both H ions (E1 kev peak) and He ions (E2keV peak). Moreover, in the analyzed orbit IMA was in a data mode which was designed to study planetary ions in a wide energy range, but low-mass solar wind ions cannot be detected at the same time below E200 ev. Fig. 3d shows the sum of the observed O and O 2 counts from the whole measurement period in orbit no An intense flux of escaping planetary ions was detected at about 03:15 03:25 UT and a detailed presentation of the IMA data at that time interval can be found elsewhere (Kallio et al., 2006a). Note that in IMA instrument some of the solar wind protons have access to O and O 2 mass channels which causes noise to IMA O and O 2 data. This possible interference complicates the interpretation of the origin of the low-energy (Eo100 ev) counts seen in Fig. 3d at 03:50 04:10 UT because IMA instrument did not measure H and He ions at that time interval at the same energy channels as the heavy ions. There are some clear similarities, but also certain obvious differences, between the simulated and the observed oxygen time energy spectrograms. First, in both the simulation and in the data the accelerated oxygen ions can be found in the E SW population region. However, in the simulation the oxygen ions can be seen already at 02:20 UT while in the data this ion population arises not until 03:15 UT. Second, in the simulation a clear oxygen population can be seen in the E low population region while in the data the low-energy ion population contains only a few counts at 04:05 UT. Third, a clear high-energy O

6 E. Kallio et al. / Planetary and Space Science 56 (2008) Fig. 3. The simulated (a) O 2 and (b) O time energy spectrograms, and the IMA time energy spectrograms of (c) H and He ions and (d) heavy ions (e.g. O and O 2 ) on orbit no Note that the simulated spectrograms are derived by using 4p field of view (FOV) while the FOV of the IMA instrument is Furthermore, the simulated O and O 2 spectrograms show normalized counts in log 10 scale ( ¼ log 10 (count s 1 )/ max(log 10 (count s 1 ))) while in IMA spectrograms the unit is log 10 (count s 1 ). See text for details of the three regions marked as E SW population, E low population and O corona population in panels (a) and (b). ion population can be identified in the simulation but only few counts are observed in the O corona population region Discussion about the differences between the simulation and the data The differences between the simulated and the observed time energy spectrograms can be associated with several issues. First, the differences may result from an incomplete comparison between the simulated and the observed energy spectra. In a more comprehensive comparison, one should generate simulated count rates by using IMA FOV and the attitude of the FOV along the orbit of MEX. One should also add (energy dependent) conversion factors in order to transform the counts recorder to the tube detector to the counts which IMA type of detector would have detected. The addition of these more advanced features is, however, beyond the scope of the current study. Another reason of the differences may be associated with the upstream parameters. For example, the direction and strength of the IMF are not fully known. The solar wind density and velocity are also not well determined either because the exact value of the plasma parameters derived from the IMA counts depends on the adopted algorithm.

7 1210 ARTICLE IN PRESS E. Kallio et al. / Planetary and Space Science 56 (2008) Fig. 4. Martian crustal magnetic field on orbit no. 555 in June 27, (a) The total magnitude of the crustal magnetic field based on a spherical harmonic crustal magnetic field model and (b) the altitude of the spacecraft. (c and d) The magnetic field lines connected to orbit no. 555 viewed from the center of the tail toward the nightside of Mars. In panel (c) the black circles show 40 positions of MEX at 03:30 04:10 UT from which the field line tracing was started. In panel (d) similar kind of magnetic field line tracing was made for 90 positions of MEX at 03:00 04:30 UT. The two identical surface color maps in panels (c) and (d) show the radial component B r ( ¼ B anomaly ( r/ r ), where r is the position vector from the center of Mars) of the crustal magnetic field at 400 km above the surface of Mars derived from the spherical harmonic crustal magnetic field model. The attitude of the magnetic anomalies corresponds to 27 June 2004 at about 03:50 UT. Note that MEX was below the altitude of 400 km near the pericenter and, therefore, some of the circles along the orbit near the pericenter are not visible in panels (c) and (d).

8 E. Kallio et al. / Planetary and Space Science 56 (2008) In addition, there might be temporal variations in the solar wind during the analyzed 2 h and 20 min time period which can disturb the asymmetries that arise in the hybrid model. The third reason is that the hybrid model does not include a self-consistent ionosphere. It is also important to recall that the orbit analyzed in this paper is especially problematic because the pericenter is at the nightside ionosphere. Finally, the differences may be associated with the Martian crustal magnetic anomalies which are not taken into account in the hybrid model. That can be seen in Fig. 4a which gives the total magnetic field associated with the crustal magnetic anomalies derived from a 901 spherical harmonic model (Cain et al., 2003; see Kallio et al. (2008); Frahm et al. (2008) for details of the crustal magnetic anomaly model). Fig. 4a shows how the total magnetic field associated with the crustal anomalies is anticipated to increase noticeably when MEX approaches the pericenter at 03:50 UT (Fig. 4b). MEX was therefore in a strong crustal magnetic field region near the pericenter. Two shaded regions are added in Fig. 4a to show two time intervals which are analyzed in more detail in Fig. 4c and d: the pericenter region (P) about 720 min from the pericenter at 03:30 04:10 UT and the region of high crustal magnetic field (H) when the magnitude of the crustal magnetic field is about half of the magnitude of the IMF. The lines in Fig. 4c and d are magnetic field lines associated with the crustal magnetic field which are connected to the circles on orbit no Two identical color maps in Fig. 4c and d give the magnetic anomaly map at the altitude of 400 km. The positive (negative) B r value means that magnetic field is pointing away (toward) from the surface. The value of B r is derived from the spherical harmonic crustal anomaly field model. Note that the magnetic field line tracing was performed by using a constant UT time of 03:50 UT where MEX was close to the pericenter. The rotation of Mars is not taken into account in Fig. 4c and d because only at a constant UT time both the magnetic field lines and the magnetic anomaly map can be represented simultaneously making it possible to study how the morphology of the magnetic field lines is linked with the strength of the magnetic anomaly field near the planet. Note that in Fig. 4c the shown positions are only about 720 min from the pericenter crossing and, therefore, the calculated field lines are also anticipated to illustrate well the situation when Mars is rotating around its axis. In Fig. 4d, on the contrary, the magnetic field lines show how 90 positions along orbit no. 555 at 03:00 04:30 UT are connected to Mars. These points correspond the positions about 50 min before, and about 40 min after the pericenter pass. In this case the rotation of Mars cannot be omitted and, therefore, in Fig. 4d the field lines far from the pericenter only give an illustration of the real situation. Nevertheless, Fig. 4d illustrates how the morphology of the magnetic field lines near the pericenter are associated with the morphology of the magnetic field lines further away from Mars. As can be seen in Fig. 4c and d, MEX pass near the pericenter high crustal magnetic anomaly region on the nightside. In large regions near the pericenter in the nightside, the magnitude of already one magnetic field component (B r ) exceeds noticeably the total magnetic field associated with the magnetic field in the hybrid model (c.f. Fig. 2c). In fact, the relative contribution of the crustal magnetic field to the total magnetic field at a given latitude longitude point is anticipated to increase when the point is in the nightside because in the hybrid model the induced magnetic field decreases with increasing SZA angle (see, for example, Kallio et al., 2006a, Fig. 7). It can also be seen in Fig. 4d that the footpoints of the magnetic field lines that are connected to the points along orbit no. 555 are focused to a relatively small magnetic cusp like region at a height of 400 km where the crustal field is high. It is an illustration of the situation in which the crustal magnetic field produces mini magnetospheres, or magnetic loops through which MEX may flow. These magnetic loops can be so strong that the region in the nightside near the magnetic anomalies cannot be connected to the IMF, as would be the case without the crustal magnetic field (see Frahm et al., 2008). 4. Discussion This paper presents the first case study in which the properties of the simulated O and O 2 time energy spectrograms were derived along the orbit of MEX by a hybrid model. The simulated spectrograms were then compared with the spectrograms measured by IMA/ ASPERA-3 instrument. The study is a continuation of the preliminary IMA data hybrid model comparisons where the comparison was based on the simulated macroscopic plasma parameters (Kallio et al., 2006a) and planetary ions collected on the x ¼ constant planes (Kallio et al., 2006b). In this paper the possibility to perform a detailed case study was based on a technique where a virtual tube detector was implemented in the hybrid model which counts ions coming close to the orbit of MEX. Only an observed vs. simulated time energy spectrogram comparison provides the maximum advantage what a kinetic model can provide to help to interpret the features seen in the ion data. In the future, more detailed data-model comparisons are called for a more comprehensive comparison between the simulation and the observations (cf. also discussion in Section 3.4). In those comparisons one should look through all MEX orbits and try to find a suitable orbit, or orbits, by keeping in mind several issues associated with its choice criteria. First, the correct FOV and its attitude: The correct narrow FOV should be used for a detector, and the attitude and the geometrical factor of IMA detectors should be taken into account. It should be noted that the attitude of IMA FOV does not remain constant in the MSO coordinate system from orbit to orbit.

9 1212 ARTICLE IN PRESS E. Kallio et al. / Planetary and Space Science 56 (2008) Second, one should find an orbit when IMA observed oxygen ions but when the strength of the crustal anomaly was so low that the plasma data in the analyzed region may be regarded to be crustal magnetic anomaly free. In such analysis one should take into account the fact that the crustal anomalies may have been affected to the properties of the measured plasma although the strength of the crustal magnetic field at the analyzed measurement point might be small. Third, the update of IMA energy tables in 2007 makes it possible for the IMA to observe lower energy ions than during the observation period In that respect, the new IMA data set provides new possibilities to make a more accurate low-energy ion comparison than presented in Fig. 2 because IMA is more sensitive to observe low-energy oxygen ions. However, on the other hand, the loss of MGS on November 2006 prevents its magnetic field observations to be used to determine the direction of the IMF clock angle in 2007 and afterward. Fourth, ASPERA-3 instrument includes also an electron detector (ELS) the data of which were not used in this paper. In the future studies, the electron data should also be taken into account because it provides information about physical processes in the analyzed orbit which can also impact the properties of ions, such as information about the magnetic connection between the IMF and the crustal magnetic anomalies and the strength of the wave activity. The electron data are also commonly used to derive estimation for the total plasma density which includes the contribution of the low-energy ions which may be difficult to obtain by ion instrument measurements. The ELS data are, indeed, available from orbit no. 555 and in the data there are clear high-energy (E4100 ev) electron flux enhancements near Mars (data not shown). It has been suggested that such electron flux enhancements may be associated with the Martian crustal magnetic anomalies (Soobiah et al., 2006). All the aforementioned orbital choice criterions and aspects, as well as other general orbit choice criteria (IMA data mode, temporal variations in the solar wind, etc.) should be taken into account in the future in more comprehensive studies. It is also worth to note that the self-consistent analysis of the crustal magnetic field is a very challenging computational task for a hybrid model and it will take a while before good spatial resolution (grid size of few tens of kilometers) results may become available. At the moment the only computationally feasible method to bring new insight to this issue is a non-selfconsistent approach where the magnetic field derived from the hybrid model (which does not contain the crustal magnetic anomaly magnetic field) is superimposed by vector addition to the crustal magnetic model (which does not include the solar wind magnetic field, see Kallio et al. (2008) and Frahm et al. (2008) for details of the superimposition method). Moreover, the spatial resolution of 180 km near the exobase implies that the ionization source terms and ion acceleration are not taken in the account in the hybrid model very accurately close to the planet. This may also be one factor which results in the lack of a very good agreement between observations and simulations. In the future, the role of the spatial resolution should be studied by using a smaller grid size near the planet. 5. Summary The present study suggests the following. First, the hybrid model produces time energy spectrograms where the planetary O and O 2 ions form a narrow energy ion beam population whose energy increases with increasing altitude from Mars (the E SW population region in Fig. 3a and b). This planetary ion population has also been observed by IMA/ASPERA-3 instrument (see Fig. 3d). Second, the simulation resulted in a high-energy (E41 kev) O ion population near Mars. These ions are pick-up O ions originating from the hot oxygen corona. Similar kind of high-energy (E41 kev) O 2 ion population does not exist because there is no similar hot molecular oxygen ion source. The properties of the high-energy O ions need further investigation but it is interesting to note there are cases where IMA instrument has observed highenergy planetary ions very close to Mars on the dayside, just a few hundred kilometers from its surface (Lundin et al., 2004). Third, the simulation resulted in low-energy planetary oxygen ion population near the pericenter. No such clear low-energy oxygen ion population was observed by MEX on its orbit no There are many possible sources for this difference such as an instrument response or a simplified ionosphere used in the model. Some of the differences may also be associated with the Martian crustal magnetic anomalies because MEX was in the high crustal magnetic anomaly region near the pericenter. Acknowledgments The authors thank Prof. Joseph Cain for the program and for the constants that were used to derive the crustal magnetic field used in Fig. 4a, c and d from the 901 spherical harmonic model. 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