University of Kansas
Image courtesy of NASA/JPL-Caltech.
A description of the photochemical model of the ionosphere used for this paper can be found in Keller et al. [1992, 1998], Gan et al. , and Cravens et al. . We adopt N2 and CH4 neutral density profiles (these are the major species) from INMS measurements made during the Ta encounter [Waite et al., 2005]. The measured exospheric temperature was 149 K +/- 3 K. The N2 and CH4 density profiles are similar to those measured by Voyager 1 in 1982 [Smith et al., 1982; Vervack et al., 2004; see Waite et al., 2005]. The abundance of minor neutral species in the upper atmosphere strongly influences the ion composition but only weakly affects the total ion (or electron) density. The minor neutral abundances were adopted from Toublanc et al.  (also see Keller et al. ).
The chemical scheme, described by Keller et al. , includes more than 51 ion species and more than 300 chemical reactions. The electron density is the sum of the ion densities. The set of chemical equations is numerically solved at each altitude. The dominant ion species near the ionospheric peak are HCNH+, C2H5+, and heavy hydrocarbon ions. The details of the ion composition do not significantly affect the electron density because the dissociative recombination rate coefficients for the relevant polyatomic ion species are about the same. We estimate an electron density uncertainty due to this effect of approximately 10 percent or less.
We adopted the altitude-dependent electron temperatures (Te) from the Gan et al.  model. Near 1200 km this model predicted Te approx.= 600 K. The Roboz and Nagy  model predicted Te approx.= 1000 K at this altitude. Both models showed very steep temperature gradients near this altitude. The Te measured by the RPWS experiment near 1200 km was approx.= 1300 K. This range of temperatures can affect the electron density by about 25 percent (see the discussion section).
The model includes ionization from both solar radiation and precipitating magnetospheric electrons. Cravens et al.  presented ion production rates versus altitude and solar zenith angle for the solar ionization case, and the ion production rates in the current paper are very similar to these (for N2+ and CH4+), although some minor differences exist due to the new INMS neutral atmosphere adopted and due to a somewhat different solar irradiance spectrum. At higher altitudes, significant ionization persists in the model for SZA up to 110 deg., due to the large neutral scale height relative to the radius of Titan (i.e., H/RT approx.= .02; for Venus this ratio is approx. 10-3).
Solar extreme ultraviolet (EUV) and soft x-ray fluxes are known to be quite variable [cf. Tobisca and Eparvier, 1998; Tobisca et al., 2000]. The solar activity level during late 2004 was generally low (F10.7 proxy indices less than approx.= 100), but for the encounter date F10.7 was approx.= 138. We adopt solar EUV fluxes from the SOLAR2000 irradiance model with F10.7 approx.= 85 [Tobisca et al., 2000]. Maurellis et al.  described the soft x-ray part of the spectrum we use. This solar spectrum is denoted "Solar 1." We also used a solar spectrum (denoted "Solar 2") uniformly enhanced from Solar 1 by the F10.7 ratio of 138/85 = 1.6. Note that solar fluxes observed at Earth are not necessarily appropriate for Saturn and that solar EUV and x-ray fluxes exhibit complex time variations including solar rotational modulations. We will demonstrate that a solar flux enhancement of a factor of 1.6 results in a 25 percent increase in ne.
The model also includes ion production from suprathermal electrons (photoelectrons from photoionization or electrons from Saturn's magnetosphere). These electrons move along magnetic field lines that have been pushed into the ionosphere by the incident magnetospheric flow [cf. Ledvina and Cravens, 1998; Ma et al, 2004]. Electron fluxes were calculated using the two-stream method [Nagy and Banks, 1970; Gan et al., 1992, 1993]. Auger electrons due to K-shell ionization were also included [Cravens et al., 2004].
Two magnetic field line configurations were adopted -- radial field lines or parabolic field lines, which approximately account for field-line draping around Titan (see the description by Gan et al., 1993). We found that magnetospheric electrons penetrate more deeply into the atmosphere by about one scale height (approx.= 70 km) for the radial case than for the parabolic case. Results will be shown only for the parabolic case. Electron fluxes were measured in the outer magnetosphere both by Voyager [cf. Schardt et al., 1984] and by the Cassini CAPS instrument [Young et al., 2004]. The distribution is approximately Maxwellian with a density of ne,mag approx.= 0.1 cm-3 and a temperature of Te,mag 100 eV, although there were variations in these values. For our calculations we adopted incident electrons with temperatures of 25 eV, 50 eV, 100 eV, and 200 eV. N2+ production rates are shown in Figure 1. Not surprisingly, electrons penetrate deeper into the atmosphere for higher temperatures.
|Figure 1. Production rate profiles for N2+ ions for incident magnetospheric electron fluxes and for a parabolic magnetic field line configuration. The magnetospheric electron populations considered have electron densities and temperatures of: (ne, Te) = (0.1 cm-3; 200 eV), (0.1 cm-3; 100 eV), (0.2 cm-3; 50 eV), (0.3 cm-3; 25 eV).|
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