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Planetary and Space Science 57 (2009) 1834–1846

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Planetary and Space Science

0032-06

doi:10.1

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journal homepage: www.elsevier.com/locate/pss

Structure of Titan’s ionosphere: Model comparisons with Cassini data

I.P. Robertson a, T.E. Cravens a,�, J.H. Waite Jr.b, R.V. Yelle c, V. Vuitton d, A.J. Coates e, J.E. Wahlund f,K. Agren f, K. Mandt b, B. Magee b, M.S. Richard a, E. Fattig a

a Department of Physics and Astronomy, Malott Hall, University of Kansas, Lawrence, KS 66045, USAb Southwest Research Institute, P.O. Drawer 28510, San Antonio, TX 78284, USAc Lunar and Planetary Laboratory, University of Arizona, P.O. Box 210092, Tucson, AZ 85721, USAd Laboratoire de Plan�etologie de Grenoble, BP 53, 38041 Grenoble Cedex, Francee University College London, Mullard Space Sciences Laboratory, Holmbury St. Mary, Dorking, Surrey RH5 6NT, UKf Swedish Institute of Space Physics, Box 537, SE-751 21 Uppsala, Sweden

a r t i c l e i n f o

Article history:

Received 14 November 2008

Received in revised form

17 July 2009

Accepted 21 July 2009Available online 29 July 2009

Keywords:

Titan

Ionospheres

Atmospheres

Chemistry

Composition

Dayside

33/$ - see front matter & 2009 Elsevier Ltd. A

016/j.pss.2009.07.011

esponding author. Tel.: +1785 864 4739; fax:

ail address: cravens@ku.edu (T.E. Cravens).

a b s t r a c t

Solar extreme ultraviolet and X-ray radiation and energetic plasma from Saturn’s magnetosphere

interact with the upper atmosphere producing an ionosphere at Titan. The highly coupled ionosphere

and upper atmosphere system mediates the interaction between Titan and the external environment.

New insights into Titan’s ionosphere are being facilitated by data from several instruments onboard the

Cassini Orbiter, although the Ion and Neutral Mass Spectrometer (INMS) measurements will be

emphasized here. We present dayside ionosphere models and compare the results with both Radio and

Plasma Wave–Langmuir Probe (RPWS/LP) and INMS data, exploring the sensitivity of models to

ionospheric chemistry schemes and solar flux variations. Modeled electron densities for the dayside leg

of T18 and all of T17 (dayside) had reasonable agreement with the measured RPWS electron densities

and INMS total ion densities. Magnetospheric inputs make at best minor contributions to the

ionosphere for these flybys, at least for altitudes above about 1000 km. At lower (o1100 km) altitudes,

the total ion densities measured by the INMS are less than the electron densities measured by the

RPWS/LP which could be due to heavy (4100 daltons) ions, which the INMS is not able to detect.

Qualitatively, INMS spectra exhibit the same ion species and 12 amu family separations for the dayside

ionospheres of T17 and T18 as were seen in the mass spectra measured during T5 (nightside). However,

the relative abundance of high-mass (m450) ion species is about 10 times less for the dayside T17 and

T18 passes than it was for the polar nightside T5 flyby, which can perhaps be explained in several ways

including differences in neutral composition, less dissociative recombination on the nightside than on

the dayside (due to lower electron densities and affecting heavier ion species more than lighter ones),

and transport of longer-lived high-mass species from day-to-night.

& 2009 Elsevier Ltd. All rights reserved.

1. Introduction

Solar radiation and energetic particles from Saturn’s magneto-sphere interact with neutrals in Titan’s atmosphere, producing anionosphere. The electron density profile in Titan’s ionosphere wasmeasured by Voyager 1 in 1980 using the radio occultationtechnique (Bird et al., 1997). Prior to the Cassini mission a largenumber of models of the ionosphere and its composition werecreated (Keller et al., 1992; Banaskiewicz et al., 2000; Galand et al.,1999; Fox and Yelle, 1997; Wilson and Atreya, 2004; Cravens et al.,2004). The Langmuir probe (LP) part of the Radio and PlasmaWave Spectrometer (RPWS) instrument measured in-situ electrondensities and temperatures in Titan’s ionosphere during the Ta

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+1785 864 5262.

(and Tb) close encounter in October 2004 (Wahlund et al., 2005).During this first flyby the Ion and Neutral Mass Spectrometer(INMS) measured densities of important neutral species in theupper atmosphere, including N2 (95%) and CH4 (5%) (Waite et al.,2005). But the INMS open source ion (osi) mode was not operatedduring this encounter and the ion composition was not measured.During the Ta flyby, the spacecraft entered the atmosphere andionosphere on the dayside, crossed the terminator at closestapproach (at an altitude of 1176 km and solar zenith angle of91.11), and continued onto the nightside (Waite et al., 2005).Comparisons between a theoretical model of the ionosphere andthe measured electron densities demonstrated that most of thedayside ionosphere is produced by photoionization by solarradiation (Cravens et al., 2005; Galand et al., 2006).

The first measurements of the ionospheric composition atTitan were made by INMS during the T5 Cassini flyby in April2005 (Cravens et al., 2006). This flyby took place entirely on the

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nightside and solar radiation could not be the source of at leastthe short-lived ion species (Cravens et al., 2009). Pre-Cassinimodels did investigate the role of impact ionization by energeticelectrons from Saturn’s outer magnetosphere transported alongmagnetic field lines as an important ionization source (cf., Atreya,1986; Gan et al., 1992; Cravens et al., 2009). Cravens et al. (2008)explored the role of impact ionization by energetic ion precipita-tion from Saturn’s magnetosphere, which might be especiallyimportant for the ionosphere detected remotely at lower altitudes(below 1000 km and down to 400 km) by the Cassini RadioScience (RSS) experiment (Kliore et al., 2008).

The composition of Titan’s ionosphere as observed by theCassini INMS was partly explained by the pre-Cassini chemicalmodels (see the references listed above). For example, it wascorrectly recognized that HCNH+, C2H5

+, and CH5+ are abundant;

however, a number of other ion species were seen in the measuredmass spectra (Cravens et al., 2006) but were not predicted tobe present with significant abundances (e.g., species at massnumbers 18 and 30). Vuitton et al. (2006, 2007) re-examinedTitan’s ion–neutral chemistry and introduced many new ionspecies (e.g., CH2NH2

+ at m ¼ 30), as well as new minor neutralspecies which have low abundances, yet have effects on the ionchemistry. Krasnopolsky (2009) created a photochemical modelof Titan’s atmosphere and ionosphere using a coupled ion andneutral chemistry, as did De La Haye et al. (2008) earlier, andobtained reasonable results when compared to Cassini data. Horstet al. (2008) created a photochemical model demonstrating thatO+ deposition is probably at the origin of the O-bearing speciesobserved on Titan. Agren et al. (2007) and Cravens et al. (2009)quantitatively analyzed the ionosphere using electron transportcodes and ion chemistry schemes for T5 conditions, demonstrat-ing that magnetospheric electrons are able to reproduce theobserved ionospheric structure on the nightside.

The current paper presents INMS data and some modelcomparisons for the ionosphere both on the dayside and on thepart of the nightside close to the terminator (i.e., solar zenith angles,SZA, up to E1051). INMS data from the T17 and T18 Titan flybys arepresented, as well as some comparison data from the RPWS/LPexperiment. The RPWS data were presented and described by Agrenet al. (2009). The T17 flyby took place on September 7, 2006, andwas entirely on the dayside (SZA from 311 to 711 for altitudes belowabout 1600 km) and the T18 flyby took place on September 23,2006, and was on both the day and nightside (SZA ranging from 751to 1041 for altitudes less than 1600 km, with a SZA at closestapproach of about 901). The Ta ionosphere is also revisited inthis paper. Both the ion composition and the vertical and hori-zontal structure of the ionosphere will be explored in the paper.Only ionization due to solar radiation is included in the modelcalculations carried out for this paper because we found that for theT17 and T18 ionospheres, magnetospheric particle precipitation wasnot required for the model to produce electron densities inagreement with measured values.

Fig. 1. Pre-flyby modeled neutral densities for C2H2. Because all density curves

show similar shapes at altitudes above about 900 km, the INMS observed neutral

densities were fitted to the Toublanc curve for use in our ionospheric model (Yung,

1987; Yelle, 1991; Vervack et al., 2004; Toublanc et al., 1995a, b).

2. Description of the model

The theoretical model of Titan’s ionosphere that we use tointerpret the Cassini data has evolved over many years and wasdescribed by Keller et al. (1992, 1994), Gan et al. (1992), Kelleret al. (1998), and Cravens et al. (2004, 2005, 2009). The modelincludes three main components: (1) a code that accounts for(primary) ion production (and photoelectron production) dueto the absorption of solar radiation, (2) an electron transport codethat determines the secondary ionization rate due to impactionization by photoelectrons, and (3) an ion chemistry code thattakes the total ion production rates for ‘‘primary’’ species and

determines ion densities versus altitude and solar zenith angle.For the current paper, a photochemical code is used for the thirdcomponent and ion transport effects are neglected. Depending onthe ion species this limits the reliability of the model results foraltitudes exceeding about 1500 km where plasma transport isexpected to be important (Cravens et al., 1998; Ma et al., 2006,2007, 2009; Cui et al., 2009a). Our model only predicts iondensities and requires as inputs the densities of neutral species.The neutral atmosphere model(s) used for this paper are mainlyderived from INMS neutral density measurements. Negative ionchemistry is not included. We will now briefly describe eachmodel component as applied to the T17 and T18 cases.

2.1. Titan’s neutral upper atmosphere for T17 and T18

The INMS in its closed source neutral (csn) mode has measureddensities of many neutral species on a number of Cassini flybys.For the energy deposition and main ion production the two majorspecies (N2 and CH4) are most important. The INMS-measured N2

and CH4 for Ta, T17, and T18 were used as model inputs asappropriate. The chemical model also requires densities of minorneutral species and where possible we used INMS data. However,our model requires densities over a large altitude range and weneeded to extend the INMS density profiles by fitting them to thedensity curve of the same species from a theoretical model(Toublanc et al., 1995a,b). Note that in the current paper ouremphasis is not on the minor neutrals themselves but on theireffects on the ionosphere, and a more general photochemicalmodel including neutral chemistry would need to take intoaccount recent developments (e.g., Lavvas et al., 2008; Vuittonet al., 2008; Krasnopolsky 2009). We attached the neutraldensities to an altitude grid, starting at 725 km, increasingby 10 km intervals and ending at 3005 km, although our mainaltitude range of interest is from about 960 to 1700 km. Forexample, Fig. 1 shows the neutral density profiles for C2H2. Froman altitude of 900 km up, the shape of the neutral density curve issimilar for each model. Consequently, the shape of the Toublanccurve was used to extend INMS observed neutral densities beyondthe observation altitudes.

For the photoionization and electron transport codes we used19 ‘‘minor’’ neutral species, in addition to molecular nitrogen andmethane. The minor species densities from the INMS aresupplemented in two ways: (1) using densities deduced fromthe ion chemistry for the T5 flyby by Vuitton et al. (2006, 2007),

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and (2) from the Toublanc et al. photochemical model. Weorganized the minor neutral densities in two different ways, asdescribed below. This approach follows to a large extent how wemodeled the nightside ionosphere for the T5 flyby (Cravens et al.,2009). Note also that densities can vary from egress to ingress,but because the derivation of minor neutrals from INMS datais thought to be more reliable during ingress, we only usedmeasured ingress neutral densities. Part of this reliability is due tostickiness to the ante-chamber wall during the egress leg of theflyby and also due to chemical reactions taking place during thattime (Cui, 2008; Cui et al., 2009b). Our two different cases ofneutral densities are as follows.

1.

Fig. 2. Neutral densities used in the model for the T18 flyby. Major neutral species

(N2 and CH4) are from INMS closed source measurements. The densities of minor

neutral species for the model case B1 are discussed in the text. For clarity, the box

around the name of the density is the same style as the line.

Case C1: The ingress densities for N2 and CH4 measured by theINMS are used. The minor neutral densities are all derived fromthe Vuitton et al. (2006, 2007) mixing ratios for the T5 flyby foran altitude of 1100 km. Although these mixing ratios were fornightside conditions (T5), they should still provide someguidance for neutral composition for a dayside flyby. Thedensities for altitudes other than 1100 km are found by scalingthe relative Toublanc et al. (1995a,b) curve to fit the Vuittonet al. density at 1100 km. If a species was not given byVuitton et al. we used values from INMS if available (thederivation of these minor densities is described in case B1).Finally, Toublanc et al. (1995a,b) density profiles were directlyused for the five species still left after this procedure (NH, CO,C3H2, C4N2, and CH3).

2.

Case B1: The same N2 and CH4 densities were used as in the C1case (i.e., INMS). However, densities of most minor specieswere derived from the measured INMS neutral mass spectrafor the relevant flybys using a spectrum deconvolutiontechnique to unravel the ‘‘cracking pattern’’ as described byMagee et al. (2009), Waite et al. (2005) and Cui et al. (2009b).In many cases this procedure provided mixing ratios foraltitudes near 1100 km, and at other altitudes we again usedthe relative altitude variations from Toublanc et al. For a fewspecies only the mixing ratios obtained by Vuitton et al. for T5were available. Since the minor neutrals can only be derivedaccurately during ingress at present, we only use the ingressdata for the major species in all our runs to maintainconsistency.

Note that we had separate B1 and C1 cases for each flyby wemodeled. Density profiles for a few neutral species for the B1 caseand T18 are shown in Fig. 2. Also see Table 1.

The determination of T17 minor INMS neutral densities is carriedout over a limited low altitude region during the inbound leg of theflyby. It was only possible to obtain one data point in some cases,and the remaining data along our altitude grid points weredetermined by fitting the Toublanc et al. (1995a, b) curve for thatminor neutral to this one data point, introducing a large uncertainty.Our approach to the densities of minor neutrals used in ourionospheric model has its limitations. A more careful considerationof the minor neutral composition is ultimately needed for all flybysincluding those on the dayside. For example, a Vuitton et al. typeanalysis should be part of such an effort. Note that Vuitton et al.(2009) reported on such an analysis for T40. This is a significant taskwith many uncertainties and for the current paper we will employthe atmosphere models just described.

2.2. The solar extreme ultraviolet and soft X-ray flux and

photoionization

Ions and photoelectrons are generated when atmosphericneutrals absorb solar photons with energies exceeding the

relevant ionization potentials (cf., Schunk and Nagy, 2000). Weused two solar flux models, although we adjusted the wavelengthbins in these models (our model uses 320 bins). For one case weused the SOLAR2000 (version 2.34) irradiance model (Tobiskaet al., 2000) for wavelengths exceeding 4.2 nm. For lowerwavelengths, we used a solar corona collisional model spectrumtuned to soft X-ray observations made by the Yohkoh observatoryas described by Acton et al. (1999). Solar conditions for the day ofthe T17, T18, and Ta flybys were adopted, although we did notcorrect for the Sun-to-Titan time lag. For the Ta flyby we adoptedF10.7 ¼ 136.7, for T17 we used F10.7 ¼ 86.7, and for T18 we usedF10.7 ¼ 70.4.

In order to test the sensitivity of our ionospheric model to ourchoice of solar flux model, we also used the EUVAC solar fluxmodel (Richards et al., 1994a, b). The EUVAC model generatesfluxes at larger wavelength bin sizes than our model, so weinterpolated these fluxes to our bin-structure. And the EUVACspectra did not extend to wavelengths as low or as high as weused (EUVAC ranged from 50–1050 A; the Tobiska model allowedus to use a range from 2 to 1306 A), so we adopted either Acton orSOLAR2000 fluxes to supplement these spectra. The solar fluxspectra used in our model for the T18 flyby are shown in Fig. 3.The SOLAR2000 and EUVAC fluxes agree reasonably well except inthe 100–300 A range where the EUVAC fluxes are a bit higher.

Given neutral densities in the atmosphere and an incidentsolar flux, altitude profiles for ion production rates and photo-electron production rates (which depend on electron energy)can be found using wavelength-dependent photoabsorption andphotoionization cross sections (see Keller et al., 1992; Cravenset al., 2004). The solar zenith angle (SZA) is another importantparameter because the opacity of the atmosphere to incomingsolar radiation at a given wavelength and altitude depends on thisangle.

Some of the photoelectrons generated by photoionization aresufficiently energetic to ionize neutrals. Photoelectrons areproduced over a large range of energies (less than 1 eV up tohundreds of eV) but typical energies are in the range 10–30 eV.The energy deposition and transport of these electrons were takeninto account with a two-stream electron transport code includ-ing both elastic and inelastic electron impact cross sections.The transport is constrained by the magnetic field lines drapedaround Titan and we adopted a draped field-line geometry withparabolic-shaped field lines (see discussion in Cravens et al.,2009) for this purpose. However, most photoelectrons are

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Fig. 3. The solar flux for solar activity conditions appropriate for the T18 flyby. For

wavelengths greater than 42 A (i.e., 4.2 nm) two spectra are shown—one from the

SOLAR2000 model and the other from the EUVAC model. The soft X-ray flux was

deduced from Yohkoh observations (see text) but was not tuned to this specific

date.

Fig. 4. Calculated electron flux versus energy for an altitude of 1210 km and for

T18 flyby conditions. Both upward and downward fluxes are shown. Suprathermal

electron fluxes measured by the Cassini CAPS ELS for T18 are also shown as the

points with error bars. The CAPS ELS 1-count background level is also shown.

Fig. 5. Ion production rate plot versus altitude for some major ion species. Total

ion production rates (i.e., direct production from both photoionization and from

photoelectrons) are shown for a few ion species. T18 outbound conditions are

adopted and the C1 neutral model is used. The solar zenith angle used was 80.51.

I.P. Robertson et al. / Planetary and Space Science 57 (2009) 1834–1846 1837

produced below an altitude of E1400 km where they deposit theirenergy locally (cf., Gan et al., 1992). The total ion production ratesdetermined with and without photoelectron transport are almostthe same. Coulomb collisions with the thermal electron popula-tion are taken into account, but mainly affect the photoelectronenergy distribution for energies less than 10 eV or so.

In the 2-stream code, initially developed by Nagy and Banks(1970) for the terrestrial ionosphere, electrons are allowed tomove in both directions along the magnetic field (defined asupward flux, F+(E,s), and downward flux, F�(E,s), where s is thedistance traveled along the field line). This code was previouslyused to model the dayside ionosphere for the Ta flyby (Cravenset al., 2004, 2005) and the nightside ionosphere for the T5 flyby, inwhich precipitating magnetospheric electrons were imposed as anupper boundary condition (Cravens et al., 2009). For the currentpaper we do not include magnetospheric electron inputs. We onlyconsidered photoionization and the 2-stream code was only usedto determine the effects of photoelectrons (i.e., secondaryionization). We found (as will be shown later) that there issufficient solar illumination even out to a solar zenith angle ofabout 1051 to produce the observed ionosphere at the altitudes ofthose observations.

Fig. 4 shows calculated and measured photoelectron fluxes forthe T18 flyby for an outbound (i.e., dayside) altitude of 1210 km.For this case, the C1 neutral densities and the Tobiska SOLAR2000solar fluxes are used. The electron spectrum measured by theCAPS ELS instrument is also shown. The spacecraft potential forT18 at this time was estimated to be about �1.5 V. The CAPSspectrum shown in Fig. 4 was shifted by 1.5 eV in order tocompensate for this potential. Overall, our modeled photoelectronfluxes agree quite well with the fluxes measured by CAPS. Thedisagreement at very low energies is partly due to the large(0.5 eV) bin size in our model. The structure in the spectra near25–30 eV is well-known (e.g., Nagy and Banks, 1970; Gan et al.,1992) and is associated with the photoabsorption of photonsat the strong solar HeII resonance line at 30.4 nm (see Fig. 3).

Fig. 5 shows ion production rates (the sum of both primaryand secondary production rates) versus altitude for several ionspecies. Opacity effects were accurately determined only for ionspecies resulting from ionization of the major species, N2 and CH4

(including production of N2+, N+, CH4

+, CH3+, CH2

+, CH+, and H+). Theoptical depths for the other species were derived only at about30 nm and were not determined as a function of wavelength;

however, the ionization associated with the major neutral speciesstrongly dominates all the ion chemistry.

The production rates in Fig. 5 are for the C1 neutral densitiesand the Tobiska SOLAR2000 solar flux. A solar zenith angle (SZA)of 80.51 was used corresponding to an altitude of 1210 km on theoutbound leg of the T18 flyby. The ion production rates havemaxima near an altitude of 1050 km for this SZA. The highestproduction rate, not surprisingly, is for N2

+, followed by theproduction rates for N+, CH4

+ and CH3+.

Fig. 6 shows primary and secondary production rate profilesfor N2

+ at a 1208 km outbound altitude (SZA ¼ 80.51) for T18.The secondary production is due to electron impact ionization bythe photoelectrons produced by the photoionization of neutrals.Note that no magnetospheric electrons were included for thisstudy—all the ionization is associated in one way or the otherwith the absorption of solar radiation.

The primary ionization clearly dominates near the peak and athigher altitudes, whereas secondary ionization (impact ioniza-tion by photoelectrons only, not by magnetospheric electrons)dominates at lower altitudes (i.e., on the ‘‘bottom side’’). The‘‘ledge’’ in the ion production rates evident in Figs. 5 and 6 at analtitude of 800 km is associated with absorption of the soft X-ray

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Fig. 6. Primary, secondary (i.e., by photoelectrons) and total production rates of N2+

at the solar zenith angle (80.51) corresponding to Cassini located at an altitude of

1208 km on the T18 egress.

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photons and ‘‘harder’’ EUV photons in the solar spectrum. Thephotoabsorption cross section for higher energy photons is lowerthan it is for lower energy EUV photons, and consequently theradiation penetrates deeper into the atmosphere. This part of thesolar spectrum is known to be especially variable and sensitive tosolar activity (cf., Schunk and Nagy, 2000), and these productionrates (and the ionosphere they produce) should be consideredto be more uncertain. These high energy photons create higherenergy photoelectrons, thus explaining the dominance of second-ary ionization in this production rate ledge.

Fig. 7. Inbound and outbound electron temperatures measured during T17 and

T18 by the RPWS/LP. Data are available only down to about 1000 km and the

temperatures are extrapolated to lower altitudes and put onto the grid used in our

model.

2.3. Photochemical model of Titan’s ionosphere

Many ionospheric models have been constructed for Titan (seethe introduction) and each of these models has a differentemphasis—some on the chemistry and others on the transporteffects. Our photochemical model is very similar to what we usedfor our T5 nightside ionosphere study (Cravens et al., 2009),although obviously the primary ion sources and the neutralatmospheres used are different. Our photochemical model onlysolves for the ion composition and currently has 73 ion species.Approximately 600 ion–neutral chemical reactions are includedas well as electron-ion dissociative recombination reactions. Wethink that we have included the key chemistry for at least themajor ions with mass numbers less than 100, but much workremains to be done on the ion–neutral chemistry, particularly atthe higher mass numbers. The number of neutral species adoptedwas 38 but our model does not solve for the densities of theseneutral species. The Keller et al. (1992, 1998) studies were thestarting point for our ion chemistry, but we now include most ofthe new ion chemistry described by Vuitton et al. (2006, 2007) asneeded to explain the ion composition measured by the INMSduring the T5 flyby (Cravens et al., 2006, 2009). In particular,Vuitton et al. described how nitrogen-bearing species such asCH2NH2

+ are created. The chemistry is complex and the reader isreferred to the Vuitton et al. papers or to the book Titan after

Cassini-Huygens (to be published in 2009 or 2010).The altitude range of the model is 725–2725 km, although

above about 1500 km transport processes begin to be important(e.g., Ma et al., 2006) and a photochemical model becomesincreasingly inaccurate; and below about 960 km in-situ Cassinidata are not available.

Dissociative recombination coefficients depend on the electrontemperature, Te (cf., McClain et al., 2004, 2009) with therecombination coefficient decreasing as some power of Te

(roughly Te�1/2 typically). We use temperatures measured by the

RPWS/LP experiment (cf., Wahlund et al., 2005) for the threeCassini flybys that we considered. Measured electron tempera-tures are about TeE600–1200 K below an altitude of 1500 km orso, but Te increases rapidly at higher altitudes (Wahlund et al.,2005). Fig. 7 shows the measured inbound and outbound T17 andT18 electron temperatures used in our models.

Ion production rates strongly depend on the solar zenith angleand for each flyby we ran our model for a large number of solarzenith angles corresponding to the location of the spacecraft.Cravens et al. (2005) published the results of similar calculationsin order to determine electron densities for the Ta flyby. The Taionosphere was also strongly dominated by photoionization. Themodeled results were compared with electron densities measuredby RPWS and very reasonable agreement was obtained. Since thattime the chemical model has been improved, mainly due to thechemistry introduced by Vuitton et al. (2006, 2007) but the detailsof the chemistry do not appear to have a strong effect on the totalion density (i.e., on the electron density).

At lower altitudes in most planetary ionospheres ion–neutralcollisions and high neutral densities limit the flow of plasma, so thatchemistry dominates the density structure. However, at sufficientlyhigh altitudes such that chemical lifetimes exceed transport times,chemical processes alone are not sufficient to predict ionosphericdensity profiles (cf., Schunk and Nagy, 2000). This is also true forTitan (e.g., Cravens et al., 1998; Ma et al., 2006; Cui et al., 2009a). Aproper consideration of transport requires plasma codes (e.g., MHD)that include the magnetic field as well as plasma sources and sinks,and such codes have not been run with the hundreds of chemicalreactions that are needed to fully understand the complex ioncomposition. The Ma et al. MHD code (and model of the magneto-spheric interaction with Titan) only includes 7 ion species as well assome rudimentary chemistry. For Ta conditions Ma et al. (2006)demonstrated that transport effects are more important thanchemistry above an altitude of roughly 1400–1500 km. Cravenset al. (2009) made some empirical estimates of transport times (i.e.,distance scale divided by flow speed) for the T5 ionosphere bydetermining the thermal and magnetic pressures from Cassini dataand then adopting either a rough vertical length scale (100 km) orhorizontal length scale (500 km). The Cravens et al. transport timeestimates were lower limits in that it was assumed that themagnetic and thermal pressure gradient forces were in the samedirection, which is not always the case. Cravens et al. also estimatedchemical lifetimes for several ion species for T5. For example, CH5

+

mainly reacts with C2H2, C2H4, and H2 and the lifetime is the inverseof the sum of the rate coefficients times the densities of theseneutral species. Cravens et al. concluded, in agreement with the Ma

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Fig. 8. Chemical lifetimes versus altitude for several ion species for T18 (out-

bound—dayside). Also shown are estimated vertical and horizontal transport times

versus altitude. Neutral winds are not included in the transport time constants.

I.P. Robertson et al. / Planetary and Space Science 57 (2009) 1834–1846 1839

et al. (2006) MHD model results, that transport effects start tobecome important for altitudes above about 1400–1500 km. Usingthese same methods we have estimated chemical and transporttimes versus altitude for the T18 outbound (dayside) ionosphere(shown in Fig. 8). Transport effects should at least be consideredabove about 1300 km for long-lived species and above about1500 km for short-lived species.

3. Cassini measurements: INMS and RPWS/LP

The introduction provided a brief review of Cassini ionosphericmeasurements but we provide a few more details here. The INMSmeasures both neutral and ion species with charge to mass ratios(m/q) ranging from 0.5 to 8.5 daltons and from 11.5 to 99.5daltons, using a radio-frequency quadrupole mass analyzer. In itsopen source ion (osi) mode, ions enter a narrow aperture beforebeing guided to the mass analyzer (see the instrument descrip-tions by Kasprzak et al., 1996, and Waite et al., 2004). The osiinstrument mode was used to measure the composition ofSaturn’s ring ionosphere (Waite et al., 2005) as well as Titan’sT5 ionosphere (Cravens et al., 2006).

In this paper, we emphasize measurements of the compositionof the dayside ionosphere as measured by the INMS in its opensource mode, but we will also compare our model results to theionospheric electron densities measured by the Langmuir probe(LP) part of the Cassini radio and plasma wave spectrometer(RPWS) experiment (see Wahlund et al., 2005; Agren et al., 2009;Gurnett et al., 2004). The sum of all ion densities should equal theelectron density due to quasi-neutrality unless significant den-sities of negative ions are present (this is a possibility below about1000 km—Coates et al., 2007). The INMS only measures specieswith mass to charge ratios up to 99.5 daltons and does notmeasure negative ion species, both of which can affect the validityof this statement in practice.

Fig. 9. The black line shows the electron density measured by the RPWS/LP and

the green line shows the total ion density measured by the INMS for T18. The

abscissa shows time, altitude, and solar zenith angle. The model densities (total

ion density) for the C1 atmosphere are shown by the dark blue line (Tobiska solar

flux) and the red line (EUVAC solar flux). The model densities for the B1

atmosphere and the Tobiska solar flux are shown with the yellow line. The light

blue line shows densities for a model with the Tobiska solar flux and a neutral

model with all densities increased by a factor of 2.5. None of the models include

magnetospheric ionization sources (i.e., they are ‘‘solar only’’). (For interpretation

of the references to color in this figure legend, the reader is referred to the web

version of this article.)

4. Results of the ionospheric model and comparisons withcassini data—the Ta flyby

Cravens et al. (2005) generated electron densities from anearlier version of the ionospheric model and compared thesedensities with the density time history measured by the RPWS/LPexperiment (Wahlund et al., 2005). Since INMS only measuredneutral densities and ion composition was not measured duringthis flyby, the modeled results could only be compared with RPWS

measurements. Cravens et al. considered several cases including:(a) ionization just from solar radiation and (b) ionization bothfrom solar radiation and from precipitation of 100 eV magneto-spheric electrons. Cravens et al. previously concluded that on theinbound leg of the trajectory (on the dayside) the solar source wasdominant, but on the outbound leg (on the ‘‘near’’ nightside) solarradiation was still important, but the magnetospheric sourcemade some contribution. Our current model has many more ionspecies and additional chemistry in comparison with this earlierstudy, as well as slightly different solar flux inputs. To reconsiderTa, we adopted the same neutral densities as used by Cravenset al. (2005). The electron densities produced by our updatedmodel are almost the same as those presented by Cravens et al.(2005). The new model’s densities are slightly higher (about 10%)near the time of �80 s (before closest approach). And again itseems that some small extra source of ionization (i.e., magneto-spheric) might be needed on the nightside beyond SZA ¼ 1001.But overall the Cravens et al. (2005) conclusions for the Taionosphere stand.

5. Model and data—the T18 flyby

The T18 flyby took place on September 23, 2006, during which theCassini spacecraft traveled over the north polar region of Titan,entering the atmosphere on the nightside at a latitude of �631 northand 2551 west longitude. Closest approach (CA) was at an altitude of960 km and was at high latitudes near the terminator.

5.1. T18—total ion density and electron density

Fig. 9 shows both measured and modeled total ion (or electron)densities versus time for the T18 flyby. The electron density measuredby the RPWS/LP and the total ion density (i.e., sum of densities of allion species up to 99.5 Da) measured by the INMS are shown. Thespacecraft not only changes its altitude during the flyby but the solarzenith angle changes too. We display the densities versus time inorder to emphasize that altitude is not the only relevant variable; infact near closest approach (CA) solar zenith angle is a more important

ARTICLE IN PRESS

Fig. 10. Calculated total ion density (i.e., electron density) altitude profiles for

many solar zenith angles and for the T18 model using the C1 atmosphere and the

Tobiska solar flux. The curves are labeled with the solar zenith angle and also with

the corresponding spacecraft time relative to the closest approach time. The

dashed curve shows the spacecraft ‘‘trajectory’’ in altitude and solar zenith angle.

The time interval is 100 s.

I.P. Robertson et al. / Planetary and Space Science 57 (2009) 1834–18461840

variable. For T18, Cassini moved from the nightside, through theterminator, and onto the dayside.

Both the RPWS/LP and INMS density time histories exhibit twomaxima—the first on the nightside at zE1150 km and SZAE991and the second on the dayside at zE1140 km and SZAE821. Thedayside peak has a density (neE2100 cm�3), about twice thenightside peak density (neE1000 cm�3). The INMS and RPWS/LPpeak densities and also the topside (i.e., altitudes above the peak)densities all agree within about 50%, except on inbound for timesbetween �275 and �375 s (with respect to closest approach).In this time interval (altitudes of about 1350–1700 km) on thenightside, it is possible that the plasma could have been flowingrapidly enough (i.e., in excess of roughly 300 m/s) such that iontransmission into the INMS aperture was affected. Spacecraftpotential can also affect the INMS densities. As mentioned earlier,the Cassini spacecraft can acquire an electrical charge (producingan electrical potential) which can affect ion transmission into theINMS mass analyzer for osi measurements. The data shown in thispaper, including the T18 data, were corrected for this spacecraftpotential, but the corrections might not be perfect, particularly athigher values, and this introduces uncertainties into the densitydeterminations. But, overall, the INMS and RPWS densities nearand above the peaks agree reasonably well.

The INMS and RPWS density time histories each show aminimum near closest approach, but the INMS total densities areabout a factor of 2.5 less than the RPWS densities near CA. Theeffects on INMS measurements of spacecraft potential or ion flowsare not likely to be important at lower altitudes; however, it hasbeen suggested that the densities of ion species with massnumbers exceeding 100 amu (and not observable by INMS) couldbe quite high below about 1100 km (J.-E. Wahlund, privatecommunication; Crary et al., 2009). In particular, the IBS sensoron the CAPS instrument obtained mass spectra showing familiesof ion species similar to those observed by INMS but at lowermass resolution and extending to mass numbers of at least200 amu. Chemical models have predicted a few ion species withsuch high-mass numbers (e.g., Keller et al., 1998; Vuitton et al.,2008), although a detailed and complete knowledge of suchchemistry is still lacking.

Another complication in the lower altitude ionosphere is thepresence of negative ions. The CAPS ELS (electron spectrometer)measured densities as high as 100 cm�3, with mass numbersextending up to hundreds of amu (Coates et al., 2007). However,due to charge neutrality, the electron density has to be less thanthe total positive ion density if negative ions are present, which isthe opposite of the INMS–RPWS comparison that was observed.

The results of our model are now used to interpret INMS andRPWS T18 measurements (Figs. 9 and 10). Model profiles for anumber of solar zenith angles are shown in Fig. 10, as well as thespacecraft location for this flyby.

In order to produce Fig. 9, the model was run for 33 solarzenith angles (except for the factor of 2.5 case for which fewerruns were used). The model profiles in Fig. 10 show that thedensity at the peak increases with decreasing solar zenith angleand that the height of the peak increases with increasing solarzenith angle—ionospheric behavior that is also found at otherplanets including the Earth (Schunk and Nagy, 2000). Also evidentin Fig. 10 is a lower altitude density ledge located below the mainpeak. This ledge is produced by harder solar EUV photons and softX-ray photons (also see Figs. 5 and 6). Cassini entered theionosphere at high altitudes on the nightside and simultaneouslymoved to both lower altitudes and lower solar zenith angles,resulting in an increasing density until a time before CA of 200 s.From this time until CA, the density decreases as the spacecraftmoves to lower altitudes below the peak. From about CA untilabout 200 s after CA, the density is mainly determined by the

decreasing solar zenith angle as the spacecraft moves from nightto day and the density increases.

The time history of the modeled electron density correspond-ing to that in Fig. 10 is also shown in Fig. 9 (C1-TOB blue curve).The model results and the measured densities from RPWS/LP arein reasonable agreement, although it appears that the altitude ofthe peak is somewhat higher for the data than for the model. Thetwo model cases with different solar flux inputs but the sameneutral atmosphere are also shown in Fig. 9 (C1-TOB for theSOLAR2000 and C1-EUV for the EUVAC). At the lowest altitudes,near CA, the EUVAC solar flux produces a modestly higher density,and in the topside ionosphere the opposite is true. The explana-tion is that the EUVAC solar flux is somewhat greater than theSOLAR2000 flux in the 10–20 nm part of the solar spectrum,which is absorbed in the atmosphere near 1000 km for solarzenith angles near 901.

The curves in Fig. 9 labeled C1-TOB and B1-TOB are for modelsthat have somewhat different densities for minor neutral species(see Table 1), but are otherwise the same. The differences arenegligible near the peak and at higher altitudes, and thedifferences are modest even at lower altitudes. This comparisonshows that the minor neutral composition does not affect theelectron density very much, although, as will be shown later,the ion composition is highly dependent on the neutral composi-tion. Different ion species have somewhat different dissociativerecombination rate coefficients, and this results in somewhatdifferent electron densities for the two cases. Because of the largenumber of runs that are made to generate a single fly–through wecreated one additional fly–through where only solar zenith angleswere run that relate to the peak electron densities and where allneutral densities (major and minor) were increased by a factor of2.5 (B1 TOB�2.5). The peak densities are not too different butthe peaks shift to higher altitudes (by an interval about equalto the neutral scale height of about 60 km) on both inbound andoutbound, as would be expected due to opacity effects. This peakshift seems to help the model–data agreement for the peakheights on both inbound and outbound (i.e., night and day,respectively).

The electron densities calculated with the photochemicalmodel for the topside ionosphere on the outbound leg (Fig. 9)agree rather well with both the RPWS and the INMS densities upto 1800 km in spite of the expectation that that transport shouldstart to be important at altitudes of 1400–1500 km and higher.

ARTICLE IN PRESS

Fig. 11. Modeled and observed CH5+ ion densities for the T18 flyby. The dashed lines

at �350 s and +350 s indicate an altitude of 1500 km above which ion transport

effects should start to become important.

Fig. 12. Modeled C2H5+ and observed mass 29 ion densities for the T18 flyby.

Fig. 13. Model HCNH+ and measured mass 28 ion densities for the T18 flyby.

I.P. Robertson et al. / Planetary and Space Science 57 (2009) 1834–1846 1841

The transport-controlled electron density must evidently mimicthe photochemical density profile (proportional to the square rootof the total neutral density for the optically thin topsideionosphere).

5.2. T18-ion composition—CH5+

Dozens of ion species were measured by the INMS, and thephotochemical model has 73 species. Before showing massspectra, we present in the next few subsections altitude/timeprofiles for a few individual interesting species. CH5

+ is mainlyproduced by the reaction of CH4

+ with CH4 and is mainly destroyedby reactions with H2, C2H2, and C2H4 (cf., Keller et al., 1998). Thepeak densities in the models agree with the measured peakdensities outbound (E200 cm�3), but all model densities (includ-ing the factor of 2.5 enhanced densities) have peaks at altitudesmuch lower than the measured peak altitude (Fig. 11). The densityfor the C1 case agrees with the data near CA (i.e., 960–1000 km).The B1 and C1 model densities differ and the differences canlargely be explained by the higher acetylene and ethyleneabundances in the C1 model than in the B1 model. The overallpoor model–data agreement for this species suggests that perhapsa key loss process is missing or that acetylene and ethyleneabundances significantly exceed the B1 or C1 neutral modelvalues. Since ion transport effects become important above analtitude of about 1500 km, the dashed vertical lines indicate thesealtitudes; nevertheless, because there is still agreement betweenRPWS and our model results, we are also showing measured andmodeled results for higher altitudes or out to a time of 450 s.

5.3. T18-ion composition—C2H5+

C2H5+ is mainly produced by the reaction of CH3

+ ions withmethane with a production rate almost equal to the N2

+ productionrate. This species has several loss processes, the most importantof which is reaction with hydrogen cyanide (HCN) producinganother important species (HCNH+). Fig. 12 shows the modeledC2H5

+ density and the mass 29 densities for the T18 flyby. Themodel–data agreement is better for this species than for CH5

+, butis still not great for the ‘‘standard’’ B1 or C1 models for either solarflux. The best agreement between model and observations is forthe B1-TOBX2.5 model case (B1 neutral atmosphere, SOLAR2000solar flux, and overall enhanced densities). Enhancing all theneutral densities moves the peaks to higher altitudes, which helpswith the agreement.

5.4. T18-ion composition—HCNH+

Protonated hydrogen cyanide (HCNH+) is mainly formed bychemical reaction between C2H5

+ and HCN. The C1 model adopts aT5 mixing ratio for the HCN neutral density, while the B1 minordensities are derived from fitting techniques applied to theobserved INMS major densities (Waite et al., 2005). As shown inFig. 13, for both solar flux cases, below about 1100 km the C1model case gives a better fit to the observed mass 28 ion densitiesthan does the B1 case. However, the B1 case does better than C1near the peaks and at higher altitudes. Again, the model peaksappear to be shifted to lower altitudes in comparison with theINMS data, and enhancing all neutral densities appears to improvethis situation.

5.5. T18-ion composition—CH2NH2+ and a complete mass spectrum

Pre-Cassini models (e.g., Keller et al., 1998) did not predictsignificant abundances of ion species at a number of massnumbers that were later observed by INMS to be present in theionosphere (Cravens et al., 2006). For example, species at massnumbers 18 and 30 were observed to be quite abundant. Mass 18was suggested to be NH4

+, produced by reaction of major ionspecies (e.g., HCNH+) with ammonia (Cravens et al., 2006; Vuittonet al., 2006) and mass 30 was suggested to be CH2NH2

+, mainlyproduced by reaction of major species with the minor neutral

ARTICLE IN PRESS

Fig. 15. A mass spectrum measured by the INMS in its open source mode for

altitudes on the outbound T18 flyby from 960 to 984 km. Also shown is the model

spectrum for an altitude of 972 km and a SZA of 891. The C1 neutral atmosphere

case was used as well as the SOLAR2000 solar flux.

Fig. 16. Mass spectrum calculated for the C1 case and the SOLAR2000 solar flux.

The altitude is 1208 km outbound with a SZA of 80.51 (see Fig. 15 caption).

I.P. Robertson et al. / Planetary and Space Science 57 (2009) 1834–18461842

species, CH2NH (Vuitton et al., 2006, 2007). Vuitton et al. deducedthe abundances of this and many other nitrogen-bearing neutralspecies using the T5 INMS data and a photochemical model andVuitton et al. (2009) have done this for the T40 flyby. The neutraldensities of ammonia and CH2NH were low enough that the INMSin its closed source mode could not measure them. Model resultsfor the T5 ionosphere agree with INMS data for mass 30, but forour dayside models we simply adopted the mixing ratios fromVuitton et al. (2006, 2007), although this neglects any day–nightand latitudinal variations in the mixing ratios of these minornitrogen-bearing species. The Vuitton et al. approach really shouldbe applied to the T18 (or any other) data to obtain new minorneutral densities, but this is a nontrivial task and for the currentpaper we use the T5 Vuitton minor neutral densities. Vuitton et al.(2009) has done this for T40.

Dissociative recombination is the main loss process forCH2NH2

+ (and for similar ‘‘terminal’’ ion species including NH4+)

and the chemical lifetimes are rather long (Cravens et al., 2009,and also see Fig. 8), although the recombination time for the T5ionosphere (nightside) is about twice what it is for the outboundportion of the T18 ionosphere.

As can be seen in Fig. 14, all model cases overestimate thedensity of CH2NH2

+ in comparison with the INMS values, exceptright near closest approach where the C1 case is not too bad. Theagreement for the outbound ionosphere and for altitudes greaterthan 1150 km is also not too bad for all model cases. Given thatthe production rate of this ion derives directly from CH2NHabundances calculated to give data–model agreement for T5 at1100 km (Vuitton et al., 2006, 2007) or other altitudes (Cravenset al., 2009, 2008), the simplest solution is to just adjustdownward the density profile of CH2NH. Just focusing on theoutbound (dayside) ionosphere, it appears that adopting anabundance of CH2NH of half what it was during T5 would beappropriate.

We have discussed only a few of the many ion speciesobserved. Figs. 15 and 16 show complete mass spectra for theT18 flyby. All mass numbers from 1 through 99 daltons aremeasured by INMS, with the exception of mass numbers 9–11. Ifno measured or modeled density is shown in these two figures, itmeans that the density was lower than the minimum scale value.Fig. 15 presents a spectrum found by averaging densities for a timeperiod on the outbound corresponding to a 24 km altitude rangeextending above CA (at 960 km). The model results werecalculated for an altitude of 972 km on the outbound trajectory.

Fig. 16 presents a spectrum for higher altitudes (1183–1233 km)on the dayside (solar zenith angle of 80.51). These spectra allow usto compare INMS observations to model results for all ion species

Fig. 14. Model CH2NH2+ and observed mass 30 ion densities for the T18 flyby.

with mass numbers less than 100 amu. The agreement betweenmodel densities and the INMS-measured densities is better overallat higher altitudes (Fig. 16) than it is at lower altitudes (Fig. 15),although even at higher altitudes the model does better for somespecies than for others. Considering species with densities greaterthan 10 cm�3 at an altitude of 1200 km (Fig. 16), and not discussedpreviously, the model densities are a factor of 4–10 too low formass numbers m ¼ 39, 40, and 53. In our model, mass 39 ismainly c-C3H3

+, m ¼ 40 is C3H4+, and m ¼ 53 consists of C4H5

+ orC2H3CN+. For m ¼ 40, Vuitton et al. also have CHCNH+ in theirmodel, and our model does not, which can probably explain theproblem. But the mass 53 species in our model and the Vuitton etal. T5 model are the same. C2H3CN+ is created by reaction of majorion species with C2H3CN; for our dayside model we just adoptedthe T5 mixing ratio from Vuitton et al. for this neutral species. Butthe abundance of this species (too low to be measured by theINMS in its csn mode) could justifiably be adjusted upward forthe dayside atmosphere. For higher mass number species near1200 km, the biggest model–data discrepancy is for m ¼ 91, whichis C7H7

+ in our model. The source of this species in our modelsinvolves C6H5

+ and C5H5+ (see Keller et al., 1998), but we did not

include the reaction of C3H5+ with C7H8 as suggested by Vuitton et

al. in their T5 model. This would help boost the density at thismass number. Again, a Vuitton-type analysis is called for, but formany altitudes and not just at 1100 km as carried out by Vuittonet al. (2007).

ARTICLE IN PRESS

Table 1Cases B1 and C1 mixing ratios for the T18 flyby for an altitude of 1100 km.

T18 mixing ratios

C1 B1

N2 0.950 0.948

CH4 0.040 0.040

H 979a 978a

H2 3923a 6167a

N 65.6a 65.5a

NH 592a 591a

CH3 2609a 2604a

NH3 6.84a 6.82a

H2O 0.29a 0.29a

C2H2 275a 460a

C2H4 979a 240a

HCN 196a 604a

CO 7.68a 7.67a

CH2NH 10.2a 10.2a

C2H6 118a 29a

C3H2 43.2a 43.1a

C3H4 3.9a 5.1a

C3H6 1.26a 1.25a

C3H8 2.25a 1.09a

C4H2 9.79a 1.81a

HC3N 39a 0.95a

C2N2 1.77a 1.77a

C4N2 22.5a 22.4a

The major species for the two cases have very similar mixing ratios, and even the

minor species abundances agree within factors of 2 or 3 for the most part.

a parts per million.

I.P. Robertson et al. / Planetary and Space Science 57 (2009) 1834–1846 1843

Overall, the model–data agreement near closest approach(Fig. 15) is worse than at higher altitudes and this is especially thecase for species with mass numbers greater than 50. The modeldensity for the m ¼ 52 species (mainly HC3NH+) is about 20 timestoo high at 972 km, although the agreement is not bad at higheraltitudes. The main source of this species is the reaction of themajor ion species with HC3N. The abundance of this speciesadopted near 1200 km must be reasonable. The problem below1000 km could be that the adopted mixing ratio for this species atthis altitude is too high or that there is a new chemical lossprocess not included at all in our model. In fact, a general problemwith the current photochemical model for lower altitudes is that itincludes very few ion species with mass numbers exceeding100 amu (m ¼ 106, 141 are included). The CAPS IBS experimenthas demonstrated that the ‘‘family’’ structure evident in thespectra shown in Figs. 15 and 16 extends up to at least 200 daltons(Crary et al., 2009–paper submitted to PSS). Clearly, reactionsmust be taking place between lower mass number species andneutrals that produce species with mass number exceeding100 amu. Such chemistry would not only populate the highermass spectrum, but would also reduce the densities of lower massspecies (which would improve our model–data comparisons).Note though that the model electron density below 1000 kmagrees rather well with the density measured by the RPWS/LP,which implies that there is a trade-off between lower and highermass ion species and that the dissociative recombination ratecoefficients for the lower and higher mass ion species cannot betoo different.

6. Model and data—the T17 flyby

The September 7, 2006, T17 flyby covered a region just north ofthe equator and the part of the flyby with altitudes of 2000 km orless was all on the dayside. The solar zenith angle ranged fromabout 301 to 701, making T17 a purely dayside flyby.

6.1. T17—total ion density/electron density

Fig. 17 shows modeled and observed electron densities forthe T17 flyby. Again, the INMS ‘‘electron density’’ at a particularaltitude is just the total density for all ion species observed in thataltitude bin. The INMS observations were again corrected for theeffects of spacecraft potential. The RPWS and INMS densities agreequite well except for altitudes below E1050 km, where the RPWS

Fig. 17. Electron density versus time, altitude and solar zenith angle for the T17

flyby, which was entirely on the dayside. The electron density measured by the

RPWS/LP and the total ion density measured by the INMS are shown, as are model

densities for several model cases as described in the text.

densities again exceed the INMS densities. The most probableexplanation is again that ion species with mass numbersexceeding 99.5 daltons are present at lower altitudes and werenot measured by the INMS.

The neutral densities for the major species in the model werefrom the INMS closed source neutral experiment. These T17densities were very close to the densities measured by the INMSclosed source for T18 and we will not display them. For the minorspecies we used the T18 mixing ratios for the B1 case, as given inTable 1. We used the electron temperatures measured by theRPWS/LP experiment for T17. The RPWS T17 inbound electrontemperatures differed by a factor of 3 from the RPWS T18 daysideelectron temperatures at higher altitudes; the outbound electrontemperatures were about the same (see Fig. 7). For instance, at1700 km inbound the T17 temperature was about 2000 K, whichwas similar to the T18 outbound temperature at that altitude;however at 2500 km the T17 inbound temperature was �104 Kwhile the T18 dayside temperature was still around 2000 K. Weused both the SOLAR2000 solar flux (Tobiska et al., 2000;designated in our figures as TOB) and the EUVAC solar fluxadjusted for the solar conditions on the date of the T17 flyby. Weonly carried out high time resolution model runs for the B1-TOBcase; for the other cases we sampled just a select few solar zenithangles.

Now consider the model–data comparisons. The model andobserved densities are very similar above 1100 km, but at loweraltitudes the model densities are significantly higher than thedensities measured either by the INMS or RPWS/LP. Just as for T18,the modeled electron density profile in the topside ionospherefor T17 agrees rather well with the measured density profiles,even above 1500 km where transport effects are expected to beimportant.

Although the C1 neutral densities were not used in thesimulations, we did do a number of model runs for B1 densitiesthat were increased by a factor of 2.5. This change made thedata–model agreement worse, even at higher altitudes where

ARTICLE IN PRESS

Fig. 18. Ion mass spectrum measured by the INMS during T17 and a mass

spectrum from the model for the SOLAR2000 solar flux. The model altitude is

1181 km on the outbound with a SZA of 33.891.

Fig. 19. Comparison of ion mass spectrum measured by INMS for T5 (nightside,

polar region) (Cravens et al., 2006), for T17 (dayside, equatorial region) and for T18

outbound (dayside, polar region). The spectra were found for data collected

between altitudes of 1075 and 1125 km and all spectra are normalized to the mass

29 density.

I.P. Robertson et al. / Planetary and Space Science 57 (2009) 1834–18461844

there was previous agreement. To generate lower densities atlower altitudes we could use a lower ion production rate or ahigher loss rate (i.e., more dissociative recombination). Thedata–model comparisons above 1100 km are good, so that it isunlikely that the solar flux or the ion production rates we usedcould be too very wrong, except perhaps at the lowest altitudeswhere the relevant solar flux at lower wavelengths is known to behighly variable (Tobiska et al., 2000). Perhaps the hard EUV or softX-ray fluxes in our solar spectra for both the SOLAR2000 and theEUVAC cases are too large. However, essentially the same solarfluxes did seem to work well for T18, even at lower altitudes.

To check the second idea (increasing the loss rate), we changedthe electron temperature at all altitudes from RPWS/LP values(TeE600–200 K) to 200 K. Dissociative recombination rate coeffi-cients decrease with temperature approximately as Te

�1/2 and theelectron density varies as neE[P/a]1/2 (where P is the ionproduction and a the dissociative recombination rate coefficient)and hence ne varies roughly as Te

1/4. The lower electrontemperature should result in an electron density of about 40%lower, which is in line with the numerical results shown in Fig. 17.However, this 40% reduction is not sufficient to bring the modeldensities down to the measured densities. Perhaps the problemlies in the details of the ion composition. Another possibility isthat the very high-mass species (m4100 amu) that are thought tobe important at lower altitudes (and not included in our model)have much higher dissociative recombination rate coefficientsthan lower mass ion species.

Fig. 18 shows the mass spectrum for T17 outbound at analtitude of 1181 km. The spectra measured by INMS on the daysidenear 1200 km for T17 and T18 are very similar overall even thoughthe latitudes for the two observations were different. Overall, atthis altitude the model agrees with the T17 INMS data, althoughthe model underestimates the densities of some species (forinstance, masses 31 and 91). The agreement is better for low-massion species than for high-mass ion species. However, the higherdensity ions still have reasonably good agreement with themeasured densities. Most of the discussion given earlier for theT18 dayside spectra also applies to the T17 (dayside) spectrum.

7. Comparison of INMS data for T5, T17, and T18 at 1100 km

The mass spectra measured by the INMS during T5 (highlatitude nightside), T17 (low latitude dayside) and T18 (highlatitude dayside) are compared in Fig. 19. Note that the measuredT17 low latitude dayside spectrum near this altitude is similar to

the T18 spectrum. Qualitatively, the T5, T17 and T18 spectra aresimilar in that the same major species (mass numbers 28, 29, 30,for example) are present, as is the ‘‘family’’ structure (spacing ofabout 12 daltons) extending up to 100 daltons. On the other hand,quantitative differences exist. The most obvious difference is thatthe relative abundances of all high-mass species (m4 50 daltons)are about a factor of 10 less for T17 and T18 than for T5.

The primary in-situ source of the ionosphere is very differentfor T5 (magnetospheric electron precipitation—Agren et al., 2007;Cravens et al., 2008) than for T17 and T18 (as discussed in thispaper—solar radiation) with lower ionization rates (and lowerelectron densities) being found on the nightside. One possibleexplanation for the relatively lower densities of high-mass ionspecies (relative to low-mass species) on the dayside comparedto the nightside evident in Fig. 19 could be linked to the lowernightside electron densities. Dissociative recombination is themain loss process for higher mass species than lower mass specieswith high-mass species more likely to be ‘‘terminal’’ ion specieswhich react much less with neutrals than do lower massspecies. The chemical lifetimes of lower mass species are mainlycontrolled by ion–neutral reactions (see time constants in Fig. 8)and should be about the same during the day as during the night ifthe neutral densities are the same. In addition, higher massspecies are more likely to have larger dissociative recombinationrate coefficients than low-mass species. The dayside electrondensity at 1100 km is about a factor of 2 greater than the T5electron density, so that dissociative recombination removeshigher mass species twice as fast during the day (T17 and T18)as during the night but the densities of low-mass species remainunaffected. However, this effect alone seems unlikely to be able toaccount for all of the T5–T17 relative density differences evidentin Fig. 19.

Differences in the minor neutral composition between thepolar night (T5) and the dayside (T17 and T18) could also affectthe relative densities of high-mass and low-mass ion species.Transport of ionospheric plasma might also help to explain theday–night ion composition differences. Cui et al. (2009a) usedINMS data to determine average ion density profiles for the dayand nightside ionospheres, and showed that the higher massspecies, with longer chemical lifetimes, had smaller diurnalvariations than did lower mass ion species, and this is consistentwith the results shown in Fig. 19. Cui et al. (2009a) also used anionospheric model that incorporated day-to-night transporteffects, and concluded that day-to-night transport was a majorsource of high-mass species in the nightside ionosphere.

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8. Discussion and conclusions

This paper has presented ion densities versus time and ionmass spectra as measured by the Cassini ion and neutral massspectrometer for Titan’s dayside ionosphere for two flybys. Thedayside data have also been compared with the nightside T5ionospheric data and with Cassini RPWS Langmuir probe data. Inorder to help interpret the dayside ion density data we used aphotochemical model that included both primary and secondaryionization and a large number of ion–neutral reactions. The modeldid not fully address the issue of the ionospheric effects of minorneutral composition or dynamical effects, leaving plenty of roomfor other types of models to address these effects and to revisit thedata that was presented here.

In this paper we have come to the following conclusions:

�

Solar radiation is the main source of the dayside ionosphereout to solar zenith angles of about 1001, at least for the Ta, T17,and T18 passes. Ionization by magnetospheric electrons wasnot needed in the model in order to explain the T17 and T18data. � Near the ionospheric peak and at higher altitudes, the total ion

densities measured by the INMS agree reasonably well withthe electron densities measured by the RPWS Langmuir probefor the dayside passes examined, which is consistent with theagreement found earlier on the nightside (T5).

� The INMS total ion densities are less than the RPWS electron

densities near an altitude of about 1000 km and lower, and thishas been attributed to the existence of chemically complex ionspecies with mass numbers exceeding 100 amu (not measuredby the INMS).

� The overall structure of the ionosphere (location and densities

of the peaks, and the shape of the profile) can be explainedwith a photochemical model that includes Cassini-baseddensities of major neutral species, electron temperatures, andboth primary and secondary ionization associated with theabsorption of solar radiation. This is true even at altitudesabove about 1500 km, where transport effects are expected tobe important.

� The ion composition measured by the INMS in the dayside

ionosphere is qualitatively very similar to that measured onthe nightside and at high latitudes (T5). The major ion species(e.g., CH5

+, C2H5+, HCNH+, CH2NH2

+, C3H5+y) had similar relative

abundances on both the dayside and nightside, but highermass species (m450 amu) had lower relative abundances onthe dayside passes than during T5, which can perhaps beexplained in several ways including differences in neutralcomposition, less dissociative recombination on the nightsidethan on the dayside (affecting heavier ion species more thanlighter ones), and transport of longer-lived high-mass speciesfrom day-to-night.

� The photochemical model did a reasonable job semi-quantita-

tively for most species and for altitudes near 1100 km andhigher, but did poorly at lower altitudes. We attribute this tomissing chemistry that leads to high-mass ions (m4100 amu)and also to our using some minor neutral abundances that areperhaps more suitable for the T5 nightside atmosphere (i.e.,Vuitton et al., 2006, 2007) than for the dayside atmosphere.

Titan’s ionosphere is both chemically and dynamically complexand much future work will be needed to fully understand thisionosphere. The current paper emphasized the ion compositionmeasured by the Cassini INMS during two flybys, T17 and T18. Inthe future, a more statistical approach should be undertaken andshould incorporate more flybys. The data in such a study should

be sorted by variables such as local time (or solar zenith angle)and latitude. Cui et al. (2009a) made some progress in this respectby considering global averages of dayside and nightside iondensity profiles and Agren et al. (2009) examined many electrondensity profiles measured by the RPWS Langmuir probe.

Models are needed to interpret the ionospheric data, anddifferent models make different simplifying assumptions andemphasize different aspects of the ionosphere. For example, theDe La Haye et al. (2008) model combined neutral and ionosphericchemistry. Ion–neutral chemistry was emphasized by Vuittonet al. (2007), Carrasco et al. (2007, 2008), Krasnopolsky (2009),and by the current paper. Vuitton et al. (2009) explored negativeion chemistry in order to explain the Cassini CAPS measurementsof negative ion species in the lower ionosphere (Coates et al.,2007). Primary ionization sources were considered by Agren et al.(2007), Cravens et al. (2009), and by the current paper. Simpledynamical and/or time-dependent effects (i.e., day-to-night flowor time-dependent magnetospheric ionization) were taken intoaccount by Cravens et al. (2009), De La Haye et al. (2008), and byCui et al. (2009a). The global ionospheric dynamics, including howthe ionosphere couples to the external magnetospheric flow, wasstudied by Ma et al. (2006, 2008) using a magnetohydrodynamicmodel with simple ion chemistry and 7 ion species. These modelshave all provided different and complementary insights intoTitan’s ionosphere, but major gaps remain in our understanding,particularly for the high-mass ion chemistry and concerning therelative role of dynamics and chemistry.

Acknowledgements

Support from the NASA Cassini project (Grant NFP45280via subcontract from Southwest Research Institute) is acknowl-edged. Model development at the University of Kansas was alsosupported by the NASA Planetary Atmospheres Grant NNX07AF47G.Solar Irradiance Platform historical irradiances are providedcourtesy of W. Kent Tobiska and Space Environment Technologies.These historical irradiances have been developed with partialfunding from the NASA UARS, TIMED, and SOHO missions.

References

Acton, L.W., Weston, D.C., Bruner, M.E., 1999. Deriving solar X-ray irradiance fromYohkoh observations. J. Geophys. Res. 104, 14,827–14,832.

Agren, K., et al., 2007. On magnetospheric electron impact ionisation anddynamics in Titan’s ram-side and polar ionosphere—a Cassini case study. Ann.Geophys. 25, 2359–2369.

Agren, K., et al., 2009. The ionospheric structure of Titan. Planet. Space Sci., inpress, doi:10.1016/j.pss.2009.04.012.

Atreya, S.K., 1986. Atmospheres and Ionospheres of the Outer Planets and theirSatellites. Springer, New York.

Banaskiewicz, M., Lara, L.M., Rodrigo, R., Lopez-Moreno, J.J., Molina-Cuberos, G.J.,2000. A coupled model of Titan’s atmosphere and ionosphere. Icarus 147,386–404.

Bird, M.K., Dutta-Roy, R., Asmar, S.W., Rebold, T.A., 1997. Detection of Titan’sionosphere from Voyager 1 radio occultation observations. Icarus 130,426–436.

Carrasco, N., Dutuit, O., Thissen, R., Banaszkiewicz, M., Pernot, P., 2007. Uncertaintyanalysis of bimolecular reactions in Titan ionosphere chemistry model. Planet.Space Sci. 55, 141–157.

Carrasco, N., Plessis, S., Dobrijevic, M., Pernot, P., 2008. Toward a reduction of thebimolecular reaction model for Titan’s ionosphere. Int. J. Chem. Kinet.,699–709, doi:10.1002/kin20374.

Coates, A.J., Crary, F.J., Lewis, G.R., Young, D.T., Waite Jr., J.H., 2007. Discovery ofheavy negative ions in Titan’s ionosphere. Geophys. Res. Lett. 34,L22103–L22107, doi:10.1029/2007GL030978.

Crary, F.J., Magee, B.A., Mandt, K., Waite Jr., J.H., Westlake, J., Young, D.T., 2009.Heavy Ions, temperatures and winds in Titan’s ionosphere: Combined CassiniCAPS and INMS observations, submitted to Planet. Space Sci. (this issue).

Cravens, T.E., Lindgren, C.J., Ledvina, S.A., 1998. A two-dimensional multifluid MHDmodel of Titan’s plasma environment. Planet. Space Sci. 46, 1193.

ARTICLE IN PRESS

I.P. Robertson et al. / Planetary and Space Science 57 (2009) 1834–18461846

Cravens, T.C., Vann, J., Clark, J., Yu, J., Keller, C.N., Brull, C., 2004. The ionosphere ofTitan: an updated theoretical model. Adv. Space Res. 33, 212, doi:10.1016/J.asr.2003.02.012.

Cravens, T.E., et al., 2005. Titan’s ionosphere: model comparisons with Cassini Tadata. Geophys. Res. Lett. 32, L12108, doi:10.1029/2005GL023249.

Cravens, T.E., et al., 2006. The composition of Titan’s ionosphere. Geophys. Res.Lett. 33, L07105, doi:10.1029/2005GL025575.

Cravens, T.E., et al., 2008. Energetic ion precipitation at Titan. Geophys. Res. Lett.35, doi:10.1029/2007GL032451.

Cravens, T.E., et al., 2009. Model–data comparisons for Titan’s nightsideionosphere. Icarus 199, 174.

Cui, J., 2008. Analysis of Titan’s neutral upper atmosphere from Cassini Ion neutralmass spectrometer measurements in the Closed Source Neutral mode, Ph.D.dissertation, University of Arizona, Tucson.

Cui, J., et al., 2009a. Diurnal variations of Titan’s ionosphere. J. Geophys. Res. 114,A06310, doi:10.1029/2009JA014228.

Cui, J., Yelle, R.V., Vuitton, V., Waite Jr., J.H., Kasprzak, W.T., Gell, D.A., Niemann,H.B., Muller-Wodarg, I.C.F., Borggren, N., Fletcher, G.G., Patrick, E.L., Raaen, E.,Magee, B., 2009b. Analysis of Titan’s neutral upper atmosphere from Cassiniion neutral mass spectrometer measurements. Icarus 200, 581.

De La Haye, V., Waite Jr., J.H., Cravens, T.E., Robertson, I.P., Lebonnois, S., 2008.Coupled ion and neutral rotating model of Titan’s upper atmosphere. Icarus197, 1.

Fox, J.L., Yelle, R.V., 1997. Hydrocarbon ions in the ionosphere of Titan. Geophys.Res. Lett. 24, 2179.

Galand, M., Lilensten, J., Toublanc, D., Maurice, S., 1999. The ionosphere of Titan:ideal diurnal and nocturnal cases. Icarus 104, 92.

Galand, M., et al., 2006. Electron temperature of Titan’s sunlit ionosphere.Geophys. Res. Lett. 33, L21101, doi:10.1029/2006GL027488.

Gan, L., Keller, C.N., Cravens, T.E., 1992. Electrons in the ionosphere of Titan.J. Geophys. Res. 97, 12136.

Gurnett, D.A., et al., 2004. The Cassini radio and plasma wave investigation. SpaceSci. Rev. 114, 395.

Horst, S.M., Vuitton, V., Yelle, R.V., 2008. Origin of oxygen species in Titan’satmosphere. J. Geophys. Res. 113, E10006, doi:10.1029/2008JE003135.

Kasprzak, W.K., et al., 1996. Cassini orbiter ion and neutral mass spectrometer.Proc. SPIE 2803, 129.

Keller, C.N., Cravens, T.E., Gan, L., 1992. A model of the ionosphere of Titan.J. Geophys. Res. 97, 12117–12135.

Keller, C.N., Cravens, T.E., Gan, L., 1994. One dimensional multispecies magnetohy-drodynamic models of the ramside ionosphere of Titan. J. Geophys. Res. 99, 6511.

Keller, C.N., Anicich, V.G., Cravens, T.E., 1998. Model of Titan’s ionosphere withdetailed hydrocarbon chemistry. Planet. Space Sci. 46, 1157–1174.

Kliore, A.J., et al., 2008. First results from the Cassini radio occultations of the Titanionosphere. J. Geophys. Res. 113, A09317, doi:10.1029/2007JA012965.

Krasnopolsky, V.A., 2009. A photochemical model of Titan’s atmosphere andionosphere. Icarus 201, 226–256.

Lavvas, P.P., Coustenis, A., Vardavas, I.M., 2008. Coupling photochemistry with hazeformation in Titan’s atmosphere: part I: model description. Planet. Space Sci.56, 27–66.

Ma, Y.-J., Nagy, A.F., Cravens, T.E., Sokolov, I.U., Hansen, K.C., Wahlund, J.-E., Crary,F.J., Coates, A.J., Dougherty, M.K., 2006. Comparisons between MHD model

calculations and observations of Cassini flybys of Titan. J. Geophys. Res. 111,A05207, doi:10.1029/2005JA011481.

Ma, Y.-J., et al., 2007. 3D global multi-species Hall-MHD simulation of the CassiniT9 flyby. Geophys. Res. Lett. 34, 24, doi:10.1029/2007GL031627.

Ma, Y.-J., et al., 2009. Time-dependent global MHD simulations of Cassini T32flyby: from magnetosphere to magnetosheath. J. Geophys. Res. 114, A03204,doi:10.1029/2008JA013676.

Magee, B., Bell, J., Waite Jr., J.H., Mandt, K., Westlake, J., Gell, D., 2009. INMS derivedcomposition of Titan’s upper atmosphere: analysis methods and modelcomparison, Planet. Space Sci., in press, doi:1016/j.pss.2009.06.016.

McClain, J.L., Poterya, V., Molek, C.D., Babcock, L.M., Adams, N.G., 2004. Flowingafterglow studies of the temperature dependencies for dissociative recombi-nation of O2+, CH5+, C2H5+, and C6H7+ with electrons. J. Phys. Chem. A 108,6706–6708.

Nagy, A.F., Banks, P.M., 1970. Photoelectron fluxes in the ionosphere. J. Geophys.Res. 75, 6260–6270.

Richards, P.G., Fenelly, J.A., Torr, D.G., 1994a. EUVAC: a solar EUV flux model foraeronomic calculations. J. Geophys. Res. 99 (A5), 8981.

Richards, P.G., Fenelly, J.A., Torr, D.G., 1994b. Correction to EUVAC: A solar EUV fluxmodel for aeronomic calculations. J. Geophys. Res. 99 (A7), 13,283.

Schunk, R.W., Nagy, A.F., 2000. Ionospheres. Cambridge Univ. Press, Cambridge.Tobiska, W.K., Woods, T., Eparvier, F.G., Viereck, R., Floyd, L., Bouwer, D., Rottman,

G., White, O.R., 2000. The SOLAR2000 empirical solar irradiance model andforecast tool. J. Atmos. Sol. Terr. Phys. 62 (14), 1233.

Toublanc, D., Parisot, J.P., Brillet, J., Gautier, D., Raulin, F., McKay, C.P., 1995a.Photochemical modeling of Titan’s atmosphere. Icarus 113, 2.

Toublanc, D., Parisot, J.P., Brillet, J., Gautier, D., Raulin, F., McKay, C.P., 1995b.Erratum: Photochemical modeling of Titan’s atmosphere. Icarus 117, 218.

Vervack, J.R., Sandel, B.R., Strobel, D.F., 2004. New perspectives on Titan’s upperatmosphere from a reanalysis of the Voyager 1 UVS solar occultations. Icarus170, 91.

Vuitton, V., Yelle, R.V., Anicich, V.G., 2006. The nitrogen chemistry of Titan’s upperatmosphere revealed. Astrophys. J. 647, L175–L178.

Vuitton, V., Yelle, R.V., McEwan, M., 2007. Ion chemistry and N-containingmolecules in Titan’s upper atmosphere. Icarus 191, 722–742.

Vuitton, V., Yelle, R.V., Cui, J., 2008. Formation and distribution of benzene on Titan.J. Geophys. Res. 113, E05007 10.1029/2007JE002997.

Vuitton, V., Lavvas, P., Yelle, R.V., Galand, M., Wellbrock, A., Lewis, G.R., Coates, A.J.,Wahlund, J.-E., 2009. Negative ion chemistry in Titan’s upper atmosphere.Planet. Space Sci. in press, doi:10.1016/j.pss.2009.04.004.

Wahlund, J.-E., et al., 2005. Cassini measurements of cold plasma in theionosphere of Titan. Science 308, 986–989.

Waite, J.H., et al., 2004. The Cassini ion and neutral mass spectrometer (INMS)investigation. Space Sci. Rev. 114 (1), 113.

Waite Jr., J.H., et al., 2005. Ion neutral mass spectrometer (INMS) results from thefirst flyby of Titan. Science 308, 982–986.

Wilson, E.H., Atreya, S.K., 2004. Current state of modeling the photochemistry ofTitan’s mutually dependent atmosphere and ionosphere. J. Geophys. Res. 109,E06002, doi:10.1029/2003JE002181.

Yelle, R.V., 1991. Non-LTE models of Titan’s upper atmosphere. Astrophys. J. 383,380.

Yung, Y.L., 1987. An update of nitrile photochemistry on Titan. Icarus 72, 468.