1 Icarus 201 (2009) Contents lists available at ScienceDirect Icarus A photochemical model of Titan s atmosphere and ionosphere Vladimir A. Krasnopolsky Department of Physics, Catholic University of America, Washington, DC 20064, United States article info abstract Article history: Received 27 January 2008 Revised 16 December 2008 Accepted 18 December 2008 Availableonline10January2009 Keywords: Titan Photochemistry Atmospheres, composition Atmospheres, chemistry Ionospheres A global-mean model of coupled neutral and ion chemistry on Titan has been developed. Unlike the previous coupled models, the model involves ambipolar diffusion and escape of ions, hydrodynamic escape of light species, and calculates the H 2 and CO densities near the surface that were assigned in some previous models. We tried to reduce the numbers of species and reactions in the model and remove all species and reactions that weakly affect the observed species. Hydrocarbon chemistry is extended to C 12 H 10 for neutrals and C 10 H + 11 for ions but does not include PAHs. The model involves 415 reactions of 83 neutrals and 33 ions, effects of magnetospheric electrons, protons, and cosmic rays. UV absorption by Titan s haze was calculated using the Huygens observations and a code for the aggregate particles. Hydrocarbon, nitrile, and ion chemistries are strongly coupled on Titan, and attempt to calculate them separately (e.g., in models of ionospheric composition) may result in significant error. The model densities of various species are typically in good agreement with the observations except vertical profiles in the stratosphere that are steeper than the CIRS limb data. (A model with eddy diffusion that facilitates fitting to the CIRS limb data is considered as well.) The CO densities are supported by the O + flux from Saturn s magnetosphere. The ionosphere includes a peak at 80 km formed by the cosmic rays, steplike layers at and kmandapeakat1060km(SZA= 60 ). Nighttime densities of major ions agree with the INMS data. Ion chemistry dominates in the production of bicyclic aromatic hydrocarbons above 600 km. The model estimates of heavy positive and negative ions are in reasonable agreement with the Cassini results. The major haze production is in the reactions C 6 H + C 4 H 2,C 3 N + C 4 H 2,and condensation of hydrocarbons below 100 km. Overall, precipitation rate of the photochemical products is equal to 4 7 kg cm 2 Byr 1 (50 90 m Byr 1 while the global-mean depth of the organic sediments is 3 m). Escape rates of methane and hydrogen are 2.9 and 1.4 kg cm 2 Byr 1, respectively. The model does not support the low C/N ratio observed by the Huygens ACP in Titan s haze Elsevier Inc. All rights reserved. 1. Introduction The atmosphere and ionosphere of Titan present a challenging task for photochemical modeling. Titan is one of three bodies in the Solar System that have nitrogen-methane atmospheres. However, while the atmospheres of Pluto and Triton are at a level of 10 μbar, Titan has a dense atmosphere with a surface pressure of 1.5 bar. The column mass of Titan s atmosphere is 10.8 kg cm 2 and exceeds that on the Earth by a factor of 10; its ratio to mass of the solid body is higher than on the Earth by a factor of 78. A breakthrough in our knowledge of Titan s atmosphere was related to the Voyager 1 flyby in Chemical composition of the atmosphere was studied using IRIS, an infrared spectrometer for a range of 180 to 2500 cm 1 (4 55 μm) with an apodized spectral resolution of 4.3 cm 1, and UVS, an ultraviolet spectrometer at 52 to 170 nm with resolution of 1.5 nm. Later observations * Address for correspondence: 6100 Westchester Park Dr. #911, College Park, MD 20740, United States. address: from ground-based and Earth-orbiting observatories in the millimeter and infrared ranges significantly contributed to the study of the chemical composition, especially for species like HCN, HC 3 N, CH 3 CN, CO, and H 2 O. The most detailed observations of Titan s chemical composition started in October 2004 and are being done currently by the Cassini mission. The Huygens probe landed on Titan and studied the atmosphere below 150 km. Visits of the Cassini orbiter into the upper atmosphere of Titan down to 950 km gave numerous data from the remote and in situ observations. To some extent, the basic tools to study the chemical composition are similar to those on Voyager 1, i.e., infrared and ultraviolet spectroscopy. The Composite Infrared Spectrometer (CIRS) covers a range from 10 to 1400 cm 1 ( μm) at an apodized spectral resolution between 0.5 and 15.5 cm 1. The Ultraviolet Imaging Spectrometer (UVIS) operates at 110 to 190 nm, and the most interesting data are obtained in the stellar occultation mode. Except the improved capabilities of CIRS and UVIS and the long-term orbiter observations, the major progress relative to the Voyager 1 is in using the Ion Neutral Mass Spectrometer (INMS) for in situ studies of the neu /$ see front matter 2009 Elsevier Inc. All rights reserved. doi: /j.icarus
2 Titan s photochemistry 227 tral and ion composition. The Cassini Plasma Spectrometer (CAPS) measures densities of heavy positive and negative ions, the Langmuir probe (LP) and Magnetospheric Imaging Instrument (MIMI) study the plasma environment around Titan. The collected data required theoretical interpretation, and photochemical modeling is the basic tool for this purpose. Using densities of N 2 and CH 4 near the surface, this modeling makes it possible to calculate vertical profiles of many dozen of photochemical products, neutral and ionized, throughout the atmosphere. Photochemical models may be full and partial. The full models involve as initial data only densities of parent species at a lower boundary which is a surface for bodies that have it. The partial models involve the background atmosphere to calculate densities of some species and neglect possible effects of the calculated species on the background atmosphere. Most of the ionospheric models for Titan are of this type. Here we will briefly discuss the published full models for Titan. The first post-voyager model was made by Yung et al. (1984) with further improvements by Yung (1987). The model was extended up to 1160 km and did not include a significant part of the upper atmosphere and the ionosphere. That was an accurate work with detailed descriptions of the basic processes. The model by Toublanc et al. (1995) covered the atmosphere up to 1250 km, also leaving out some of the upper atmosphere and ionosphere. They erroneously assumed that absorption of the EUV photons by N 2 results in dissociation into 2 N( 2 D) while the true basic processes are N 2 + hν(λ < 80 nm) N e, N CH 4 CH H + N 2, with no production of atomic nitrogen. That assumption meant a significant overestimate of the production of atomic nitrogen and nitriles. Toublanc et al. (1995) neglected all reactions that involve CH, and this also adversely affects the conclusions of their model. The model by Lara et al. (1996) extends up to 1432 km and does not include the ionospheric processes. They repeated the error from Toublanc et al. (1995) concerning the production of atomic nitrogen. The first model which coupled the neutral atmosphere and the ionosphere of Titan was made by Banaszkiewicz et al. (2000). In their model the vertical distributions of the hydrocarbon and oxygen compounds were calculated without including the nitrogen species (except for N 2 of course). These distributions were then used to calculate the concentrations of the nitrogen compounds. This neglect of nitriles onto the balance of hydrocarbons may be a significant source of error. The ion densities were calculated in Banaszkiewicz et al. (2000) neglecting the ambipolar diffusion and escape of ions. This neglect is a reasonable approximation below 1350 km (see below). The latest pre-cassini photochemical model for Titan was made by Wilson and Atreya (2004, hereafter WA04). Thisisaratherde- tailed and accurate model. Some comments may be done to the lower boundary conditions, where H 2 wasassignedinsteadofbeing calculated by the model. Densities of all other photochemical products are adopted to be zero at the surface. This actually means a very effective removal of all products at the surface. 32 photochemical cycles are described in WA04. However, no quantitative assessment of these cycles is given, and it is not clear which cycles are important and which are negligible. The haze absorption is overestimated in the model (see Section 2) making the photolytic processes ineffective below 250 km. The model did not involve ambipolar diffusion and ion escape. Recently Sekine et al. (2008a) studied heterogeneous reactions of H with an analog of Titan s haze (tholin) and then effects of these reactions on Titan s photochemistry (Sekine et al., 2008b). We will also include these reactions into our model. A detailed model of Titan s neutral chemistry was recently made by Lavvas et al. (2008a, 2008b, hereafter LCV). It is probably the first model that couples photochemistry with thermal balance, radiative transfer calculations, and the haze microphysics. It involves 520 reactions of 68 species. Along with WA04, LCV will be a standard reference throughout the paper, and we will consider in more detail its advantages and some weak points related to boundary conditions, properties of the haze, and neglect of ion chemistry. Here we discuss some weaknesses of the published models because avoiding of these weaknesses is among the major motivations of the current work. Actually the published models resulted in a gradual and significant progress in the field and present many interesting findings. A two-dimensional model that coupled photochemistry and transport processes was made by Lebonnois et al. (2001). General circulation models with photochemistry may be developed as well. However, two- and three-dimensional models are mostly aimed at geographic variations of photochemical products, while the basic chemistry of the atmosphere may be properly analyzed using onedimensional models. We intend in this work to respond to the recent Cassini observations, to remove some weaknesses of the previous models, and to update some kinetic data in the model. 2. Ultraviolet absorption by haze Photochemically active photons extend up to λ 300 nm on Titan, and the haze extinction significantly affects the photolysis rates, especially in the spectral range of 200 to 300 nm. Pre-Cassini models for Titan s haze (see Rannou et al., 2003, and references therein) involve the cloud microphysics, transport, and radiative transfer calculations. Those models were based on the Voyager, ground-based, IUE and HST observations. Two spectrometers at the Huygens probe (Tomasko et al., 2005, 2008) that covered the range of 480 to 1700 nm in the upward- and downward-looking modes provided the most detailed data on the haze structure from the surface to 150 km. The haze particles on Titan are aggregates of small spherical monomers with radius of 0.02 to 0.1 μm. The Huygens spectral and polarization data (Tomasko et al., 2008) arefittedusingthe monomer radius of 0.05 μm, the number of monomers in the particle N = 3000, and the particle number density is equal to n = 5 cm 3 below 80 km and decreases with a scale height of 65 km from 80 to 250 km. Real and imaginary refractive indices of the particle material were adopted from Khare et al. (1984). We will extend this scale height up to 500 km. The haze is optically thin above 250 km, and this extension weakly affects our calculations. We calculate optical properties of the aggregate particles (extinction cross section, single scattering albedo a, and asymmetry factor g) using a code from Rannou et al. (1999). Following Rannou et al. (2003), we double the imaginary index from Khare et al. (1984) below 300 nm and use the index of i. Forachosen wavelength of 250 nm we calculate optical depths, a, and g for ten layers from 0 to 200 km and two layers at km and above 250 km. Optical depth of the total atmosphere for Rayleigh scattering is equal to 34 at 250 nm using the cross section for N 2 from Allen (1976). The Rayleigh extinction exceeds that of the haze below 50 km. Therefore we add the Rayleigh scattering and correct the values of a and g for each layer. The next step is to calculate reflection and transmission of each layer. Our calculations show that the atmosphere is opaque in the
3 228 V.A. Krasnopolsky / Icarus 201 (2009) Fig. 1. Calculated reflections r and transmissions t of the layers and the ratio of total to incident illuminations. This ratio involves the reflected sunlight and exceeds one above 350 km. These data refer to λ = 250 nm. The illumination ratios for λ = 200 and 300 nm are very similar. UV below 80 km. The asymmetry factor is within above 80 km. Van de Hulst (1980) published detailed tables of reflection and transmission for multiple scattering in layers with various optical depth and single scattering albedo a for g = 0.5 and Actually a and g may be combined in ( ) 1/2 1 a s =. 1 ag For example, spherical albedo of a semi-infinite atmosphere is equal to A = (1 s) s s with uncertainty of (van de Hulst, 1980). Values of the layer reflections r i and transmissions t i (Fig. 1) were obtained by interpolation from the tables in van de Hulst (1980). Then the layer adding technique was applied to calculate downward and upward fluxes at the layer boundaries using the two-stream approximation (van de Hulst, 1980). If a layer r i, t i is between levels i and i + 1, then reflection and transmission of this system for diffuse light are R i+1 = r i + R it 2 i 1 r i R i and T i = T i+1t i 1 r i R i. The results (Fig. 1) are insensitive to the surface albedo that was adopted at The profile at 250 to 500 km was calculated using a single scattering approximation and the haze scale height of 65 km. The similar calculations were made for 200 and 300 nm, and the results are very close to those for 250 nm. Rayleigh scattering, which is strongly wavelength dependent, becomes weak above 100 km, while the haze extinction is almost constant for the constant refractive index. Transmission of the haze calculated for our model at 300 nm is compared to similar data from the previous models in Fig. 2. A simple exponential approximation adopted by Yung et al. (1984) and then applied by Lara et al. (1996) and Banaszkiewicz et al. (2000) agrees better with our calculations than the radiative transfer modeling in Toublanc et al. (1995) and WA04. Our calculations presume more intense photochemistry at 100 to 300 km than that for the haze absorption in WA04. LCV involve detailed modeling of the haze production, microphysics, sedimentation, absorption, and its effect on Titan s thermal balance. However, we have not found in LCV an altitude profile of the UV absorption by the haze which is the most important for Fig. 2. Haze transmission plus reflection at 250 nm calculated in our model, approximated by Yung et al. (1984), andcalculatedbytoublanc et al. (1995) and WA04. Lara et al. (1996) and Banaszkiewicz et al. (2000) adopted the approximation by Yung et al. (1984). photochemistry. Therefore, we cannot compare the results of LCV with our calculations. LCV assumed eddy diffusion for the haze that exceeds that for the gas species by two orders of magnitude and adopted spherical particles for their calculations of the haze. It is not clear how eddy diffusion may be different for gas and dust, and the spherical haze particles disagree with the observed polarization. However, these assumptions did not affect the photochemical calculations in LCV. 3. Initial data Observations from the Huygens probe confirmed Titan s radius of 2575 km and gravity acceleration of cm s 2 near the surface. Titan is at a mean distance of 9.54 AU from the sun. These values are among the basic input data for photochemical modeling Background atmosphere Temperature and density profiles of the atmosphere below 1400 km were measured during the descent of the Huygens probe (Fulchignoni et al., 2005) using the Huygens Atmospheric Structure Instrument (HASI). The most detailed measurements were made at the parachute phase of the descent below 150 km. The INMS measurements of N 2 and CH 4 cover a range of 950 to 1500 km (Yelle et al., 2008). Temperature profiles at 200 to 500 km are retrieved from the CIRS spectra (Vinatier et al., 2007). Strobel (2008) calculated a temperature profile for slow hydrodynamic escape from Titan. The temperature profile adopted in our model is shown in Fig. 3. It is mostly based on the HASI measurements below 150 km and the INMS measurements above 1000 km and reflects the main features from the other observations. To reduce the model uncertainties, it is better to calculate densities of both the parent species and the photochemical products with the same technique. Then the finite difference method results in a minor error of absolute densities in each step which is accumulated to the upper boundary. To compensate for this error, our mean temperatures in the upper atmosphere are slightly higher than those measured. It follows from the atmospheric density measured by HASI near the surface and the Gas Chromatograph Mass Spectrometer (GCMS) observations of 4.9% of CH 4 (Niemann et al., 2005) that [N 2 ] 0 = cm 3 and [CH 4 ] 0 = cm 3.TheN 2 profile in Fig. 3 has been corrected for a true mean molecular mass
4 Titan s photochemistry 229 We will consider two models with different eddy diffusion. Eddy diffusion coefficient varies from cm 2 s 1 below 100kmto10 8 cm 2 s 1 above 700 km in our basic model, and the relationship for K from Hörstetal.(2008)is adopted for our second model. This model is discussed in Appendix A. Eddy diffusion in the basic model is rather close to those in the previous models. Both profiles of K are shown in Fig. 3. Coefficients of molecular diffusion are taken from Marrero and Mason (1972) and WA04. They are given in the form of D = AT s /n. A = and s = 0.75 for the diffusion of CH 4 in N 2 (Fig. 3). Homopause D = K is at 900 km for this diffusion. Diffusion coefficients for many species are unknown, and we use for them the coefficient for CH 4 corrected for molecular mass μ (Krasnopolsky and Cruikshank, 1999): 0.75 ( ) 16 T /μ 1/2 D = cm 2 s 1. n /16 Fig. 3. Vertical profiles of temperature, N 2 density, eddy diffusion, and molecular diffusion coefficient of CH 4 in N 2. The calculated profile of N 2 is based on the Huygens HASI observations near the surface and perfectly fits the INMS data at km (Yelle et al., 2008). Profiles of eddy diffusion for two models are shown by the solid and dashed lines, respectively, and the latter is taken from Hörst et al. (2008). These profiles of eddy diffusion result in positions of homopause K = D at 900 and 820 km, respectively. calculated by the model. It is in excellent agreement with the INMS measurements at km. A typical collision cross section for species at the planetary exobases is cm 2 (Chamberlain and Hunten, 1987), then the exobase is at a level where [N 2 ] ( H) cm 3, that is, at 1360 km. (H is the scale height.) Thermal escape is rather insensitive to a chosen upper boundary, and we choose it at 1600 km, similar to WA04. Rate coefficients of ion recombination are proportional to T α e in our model. Here T e is the electron temperature which is equal to 720 K above 1200 km and gradually decreases to the neutral temperature at and below 950 km (Galand et al., 2006) Eddy, molecular, and ambipolar diffusion Eddy diffusion coefficient K typically varied from 10 4 cm 2 s 1 below 150 km to 10 8 cm 2 s 1 above 700 km in the previous models. Measurements of the Ar mixing ratio in the lower atmosphere by the Huygens Gas Chromatograph Mass Spectrometer (Niemann et al., 2005) are coupled with those in the upper atmosphere by INMS (Yelle et al., 2008) to get values of eddy diffusion in the upper atmosphere. The derived eddy diffusion coefficients are (2 5) 10 7 cm 2 s 1 (Yelle et al., 2008) and 10 8 cm 2 s 1 (LCV). Hörst et al. (2008) argue that large and slow varying with height abundances of major photochemical products retrieved from the CIRS limb observations at km require very low eddy diffusion below 100 km with a steep increase above this altitude. They recommend K (z) = ( /400)(p/p 0 ) 2 cm2 s 1, where p 0 = 177 μbar and K = 400 cm 2 s 1 for p > p 0. This means that eddy diffusion increases from 75 km to 300 km by almost five orders of magnitude and K n 2 (n is the number density) in this altitude range. For example, eddy diffusion generated by tides and gravity waves is proportional to n 1/2 if the waves propagate in the atmosphere without absorption. It is not clear which mechanism may be responsible for such a steep increase in eddy diffusion. Our model is low sensitive to the chosen diffusion coefficients. Coefficient of ambipolar diffusion of ion i is equal to D i = k(t e + T i ) m i ν i, where T e and T i are the electron and ion temperatures, m i = μ i /A is the ion mass, A is the Avogadro number, k is the Boltzmann constant, and v i is the collision frequency between ion i and the neutral gas molecules. According to the Langevin formula, this frequency is [ ( 1 ν i = 2πen Aα + 1 )] 1/2. μi μ 0 Here e is the electron charge, μ 0 = 28 is the molecular mass of N 2, and α is the polarizability which is equal to cm 3 for N 2 (Atreya, 1986). Substituting the numerical values with T e = 720 K and T i = 200 K (Galand et al., 2006; Vuitton et al., 2007), D i = nμ i (1/μ 0 + 1/μ i ) 1/2 cm2 s Polymerization and condensation Recent measurements of reactions between H atoms and a Titan s haze analog (Sekine et al., 2008a) are the only reliable data on heterogeneous chemistry on Titan. We suspect that many radicals might react with the haze on Titan but do not include these reactions because of the lack of the data. Formation of polymers in reactions between some radicals and molecules was considered in WA04 and LCV. We apply their reactions and rate coefficients in our model. Many species in our model are subject to condensation near the tropopause at 44 km with T = 70.4 K (Fulchignoni et al., 2005). These species are CH 4,C 2 H 2,C 2 H 4,C 2 H 6,C 3 H 4,C 3 H 8,C 4 H 2, C 4 H 4,C 4 H 6,C 4 H 8,C 4 H 10,C 6 H 6,C 7 H 8,HCN,CH 3 CN, HC 3 N, C 2 H 3 CN, C 2 H 5 CN, C 4 H 5 N, C 5 H 5 N, C 2 N 2,H 2 O, and CO 2. Saturated vapor densities for these species (Table 1) are either taken from Atreya (1986), Moses et al. (1992), and Lara et al. (1996) or calculated using data from Lide (2006) and Khanna (2005a, 2005b). We apply a three-term approximation for saturated vapor pressures which corresponds to p s = at α e A/T where a, α, and A are the fitting parameters and subscript s means saturation. A condensation term is added to the loss term if n i /n is = S i > 1. Yung et al. (1984) and WA04 used L si = A(1 1/S i ). LCV recommend L si = A(S i 1) exp( 0.5/ ln2 (S i + 1)) ln 2 (S i + 1)
5 230 V.A. Krasnopolsky / Icarus 201 (2009) Table 1 Saturated vapor pressures used in the model. a CH b 4 (67 91 K) /T /T /T /T 4 C 2 H b 2 ( K) /T ln T C 2 H b 4 ( K) /(T 2.555) C 2 H b 6 (30 90 K) /(T 0.561) C 3 H c 4 ( K) /T ln T C 3 H 8 ( K) /T ln T C 4 H b 2 ( K) /T ln T C 4 H 4 ( K) /T ln T C 4 H 6 ( K) /T ln T C 4 H 8 ( K) /T ln T C 4 H 10 ( K) /T ln T C 6 H 6 ( K) /T C 7 H 8 ( K) /T ln T HCN c ( K) /T HC 3 N c /T CH 3 CN c ( K) /T C 2 N 2 ( K) /T ln T C 4 N c 2 ( K) /T C 2 H 3 CN ( K) /T ln T C 2 H 5 CN d ( K) /T ln T C 4 H 5 N ( K) /T ln T C 5 H 5 N ( K) /T CO 2 ( K) /T ln T H 2 O e ( K) /T ln T T a These relationships give log 10 P S (Pa). Expressions without references were calculated using vapor pressures from Lide (2006). b Moses et al. (1992). c Lara et al. (1996). d Khanna (2005a, 2005b). e Murphy and Koop (2005). to avoid oscillations of the solution in the region of condensation. We have found that their term is rather similar to L si = A ln S i which we used in our model. A = 10 7 s 1 in our basic model and 10 8 s 1 in the model with eddy diffusion from Hörst et al. (2008). A is reduced for some species with very low loss terms by a factor of 10 or 100 to deplete oscillations Sources of energy and chemistry Solar UV and EUV photons are the main source that drives photochemistry on Titan. Data on the solar UV radiation are taken from Woods et al. (1996); those for the EUV radiation from Richards et al. (1994). The photon fluxes are chosen from these papers for the medium solar activity. The model is calculated for the global mean conditions, that is, the solar flux is half of that at 9.54 AU to account for the night side, and cosine of the solar zenith angle (SZA) is 0.5. Sources of ionization in Titan s atmosphere are shown in Fig. 4. Photons with λ<25 nm produce energetic photoelectrons that increase ionization and dissociation of the atmospheric species. Electron energy deposition in N 2 was considered in detail by Fox and Victor (1988). To account for the photoelectron impact ionization, we assume that effective photoionization yield exceeds one for photons with λ<25 nm (Krasnopolsky, 1986). According to Nicolas et al. (2003), dissociative photoionization of N 2 yields mostly N + + N( 2 D) (hereafter N ). We assume that this also refers to dissociative ionization of N 2 by electrons, protons, and cosmic rays. Precipitation of magnetospheric electrons is another source of ionization and chemistry that maintains the nighttime ionosphere on Titan. The nighttime ionosphere was observed at the Cassini encounters T5 and T21. According to Cravens et al. (2006, 2008, 2009), the T5 observations were conducted during a strong precipitation event when the ionization rate exceeded that for T21 by a factor of 10. The conditions for T21 may be considered as mean. The ion densities were comparatively high during T5 and were dis- Fig. 4. Sources of ionization in Titan s atmosphere. Ionization by magnetospheric electrons is shown for mean conditions and a strong precipitation event T5. cussed in detail and subject to modeling by Vuitton et al. (2007) and Cravens et al. (2008, 2009). We adopt a profile of ionization by magnetospheric electrons given as ( ( ) P N + ) h h0 2 = a/ exp + e (h h 0)/H, H with a = 0.8 and 0.08 cm 3 s 1 for T5 and the mean conditions, respectively, h 0 = 1050 km, and H = 150 km (Fig. 4). Electron impact dissociation of N 2 for E > 50 ev forms nitrogen atoms in the ground and metastable states in almost equal quantities. Using the data of Fox and Victor (1988), productions of ions and atomic nitrogen by magnetospheric electrons are in the proportion N + 2 :N+ :N:N = 1:0.17:0.84:1.0 fore = 100 ev. The adopted profile for the mean conditions results in the column ionization rate of cm 2 s 1. (Here and below all column rates are reduced to the surface, that is, scaled by (1 + h(km)/2575) 2.) Ionization of N 2 by magnetospheric protons was calculated by Cravens et al. (2008) for a strong precipitation event and for typical conditions. The latter is used in our model, and the column rate is cm 2 s 1. We assume that production ratios of N + 2 :N+ :N:N for protons are similar to those for magnetospheric electrons. Processes of ionization and dissociation by the cosmic rays are essential in the lower atmosphere. We take a profile of the cosmic ray ionization from Molina-Cuberos et al. (1999). ThesolarUV photons do not reach the atmosphere below 100 km because of the haze absorption (Fig. 1), and the cosmic rays are the important source of chemistry in the lower atmosphere. Production of N + 2 :N+ :N:N = 1:0.22:0.73:0.95 assumed for the cosmic rays corresponds to that for electrons at 200 ev from Fox and Victor (1988). This energy is much higher than the ionization potential of N 2 (15.6 ev), and the data may be considered as a high-energy limit (Born approximation). This profile peaks at 70 km with the column ionization rate of cm 2 s 1. We neglect direct ionization of CH 4 and other species by the magnetospheric electrons, protons and cosmic rays because this ionization is much weaker than the effects of ion-neutral reactions. Column ionization rates by the magnetospheric electrons, protons and the cosmic rays are significantly smaller than that for the solar EUV photons ( cm 2 s 1 ). Delivery of water into Titan s atmosphere occurs by precipitation of the interplanetary dust and meteorites. Vertical profile of the H 2 O deposition rate for this source was calculated by English et al. (1996) and used in our model. This profile peaks near 700 km with the column production of cm 2 s 1.
6 Titan s photochemistry 231 The Cassini Plasma Spectrometer (CAPS) and Voyager detected O + ions precipitating from Saturn s magnetosphere into Titan s atmosphere (Hartle et al., 2006) withafluxof 10 6 cm 2 s 1.These ions are the major source of CO on Titan (Hörst et al., 2008). These ions are stopped and form a flux of neutral oxygen atoms near 1100 km (Hörst et al., 2008). The low energy O + ions mostly charge transfer with neutrals forming fast O atoms. They may be effective in sputtering. Ionization by neutral neutral collisions is comparatively weak and neglected in our model Species and their reactions The number of the C H N O species is very large, and the number of their reactions is huge. Our goal is to restrict both numbers and include in the model the observed species, species that may be essential in balances of the observed species, and some species that have not been observed but may be rather abundant and photochemically important. Fitting to the observed ion mass spectrum at 1100 km by means of the ion chemistry (Vuitton et al., 2007) involves many species that have not been observed previously. Table 2 is a list of species involved in the model. Comparing hydrocarbons in this list and in WA04, we remove C 6 H 7 and add C 7 H 4,C 7 H 8, and four bicyclic hydrocarbons. We do not distinguish isomersinourmodel.thenumberofnitrilesisincreasedrelative WA04 by inclusion of five species from Vuitton et al. (2007). Three species are metastable: N( 2 D), 1 CH 2, and excited C 4 H 2. They are marked by asterisks. LCV involves C 5 H 3, C 6 H 3, C 6 H 7, C 2 H 5 N, and H 2 C 3 N that are lacking in our model but does not include seven hydrocarbons (C 7 H 4, C 7 H 8, C 8 H 2, C 9 H 8, C 10 H 8, C 10 H 10, C 12 H 10 ) and seven nitrogen compounds (N 2 H, CH 2 CN, C 2 H 5 CN, C 4 H 3 N, C 4 H 5 N, HC 5 N, C 5 H 5 N) from our list. LCV does not cover oxygen and ion species and their chemistries. Photodissociation and basic photoionization processes are given in Table 3, reactions of neutral species in Table 4, and ion reactions in Table 5. Rate coefficients of three-body reactions are usually given by their low- and high-pressure limits k 0 and k.we use a limiting density n = k /k 0. Then the effective two-body rate coefficient for three-body reaction is equal to k 0 n/(1 + n/n ). Though this expression is an approximation, it is rather accurate and fits both limits. Values of n are given in Table 4 for T = 170 K that is close to the mean temperature at 200 to 1100 km. Column production of hydrocarbons is cm 2 s 1, production of nitriles is cm 2 s 1, production of oxygen species is cm 2 s 1, and production of ions is cm 2 s 1. Evidently oxygen species weakly affect chemistry of hydrocarbons and nitriles. The previous models predict very low abundances of all oxygen species except CO, CO 2, and H 2 O. Therefore we include only those oxygen species and reactions that look Table 2 Species calculated in the model. a H H 2 CCHCH 2 CH 2 CH 3 CH 4 C 2 C 2 H C 2 H 2 C 2 H 3 C 2 H 4 C 2 H 5 C 2 H 6 C 3 H 2 C 3 H 3 C 3 H 4 C 3 H 5 C 3 H 6 C 3 H 7 C 3 H 8 C 4 H C 4 H 2 C 4 H 2 C 4H 3 C 4 H 4 C 4 H 5 C 4 H 6 C 4 H 8 C 4 H 10 C 6 H C 6 H 2 C 6 H 4 C 6 H 5 C 6 H 6 C 7 H 4 C 7 H 8 C 8 H 2 C 9 H 8 C 10 H 8 C 10 H 10 C 12 H 10 polymers NN N 2 NH NH 2 NH 3 N 2 HN 2 H 2 N 2 H 3 N 2 H 4 CN C 2 NC 3 N HCN H 2 CN CH 2 NH CH 3 NH 2 CHCN CH 2 CN CH 3 CN HC 3 N C 2 H 3 CN C 2 H 5 CN C 4 H 3 N C 4 H 5 N HC 5 NC 5 H 5 N C 2 N 2 C 4 N 2 OOHCO H 2 O HCO CO 2 H 2 CO CH 3 OH CH 2 CO CH 3 CO CH 3 COH H 5 C 2 O CH + 2 CH+ 3 CH+ 4 CH+ 5 C 2H + 2 C 2H + 3 C 2H + 4 C 2H + 5 C 2H + 7 C 3H + 3 C 3H + 4 C 3H + 5 C 4H + 3 C 4 H + 5 C 5H + 5 C 5H + 7 C 6H + 5 C 6H + 7 C 7H + 7 C 9H + 11 C 10H + 9 C 10H + 11 N + N + 2 NH+ 4 N 2H + HCN + HCNH + CH 3 CNH + HC 3 NH + C 2 H 3 CNH + CH 2 NH + 2 C 4 H 3 NH + e a Observed species are bold, species derived by the ion chemistry fitting to the ion mass spectrum (Vuitton et al., 2007) are italic. Asterisks mark metastable species. The model involves 85 neutral species and 33 ions. All ions in the list except C x > 8 H y have been detected using the INMS (Vuitton et al., 2007). essential for the balances of CO, CO 2, and H 2 O. Our model involves 31 reactions of 12 oxygen species (135 reactions of 15 species in WA04 and 44 reactions of 14 species in Hörstetal.,2008). 21 reactions of 10 oxygen species in our model are similar to those in Hörstetal.(2008). The numbers of hydrocarbon and nitrile reactions in our model may be compared with those in WA04 and LCV (172, 224, and 370 reactions of hydrocarbons and 101, 91, and 154 reactions of nitriles, respectively). Overall, our model involves 304 reactions of 83 neutral species. While the reaction rate coefficients are given, there are very few data on the reaction rates in the previous models, and it is difficult to estimate relative importance of different reactions. Therefore, though the number of reactions in our model is significantly smaller than those in WA04 and LCV (415 versus 855 and 524 reactions), a further reduction of this number is possible. All products in our reaction lists are identified as chemical species or polymers. There are 73 reactions in WA04 and 75 reactions in LCV with unidentified products. This actually means irreversible loss of the reactants in the models, similar to that by polymerization. The so lost abundances are not discussed in those papers. We intend to model the most abundant ions in Titan s ionosphere. Therefore our model will not cover some minor mass peaks in the INMS ion spectrum (Vuitton et al., 2007). Mixing ratios of the photochemical products at the ionospheric altitudes are smaller than 10 3 while their photoionization cross sections are comparable to those of N 2 and CH 4. Therefore the primary ionization processes are essential only for N 2 and CH 4, and all numerous ions on Titan are mostly formed by ion-neutral reactions. Again, if a reaction of an ion proceeds with CH 4 with a rate coefficient comparable to those with other species, then the reactions of this ion with the other species may be neglected. Many of the ion-neutral reactions on Titan are actually the proton exchange reactions. These reactions proceed from ions with lower proton affinities to those with the higher affinities (Hunter and Lias, 1998; Vuitton et al., 2007). Our approach results in a reduction of the number of ion reactions from 404 in WA04 to 111 in our model. The numbers of ions are 33 in both models, though some ions are different. This may be also compared to 150ionsand 1250 ions reactions in Vuitton et al. (2007). The mass spectrum of ions observed by INMS contains 57 mass peaks. We believe that our choice is adequate to our goal to model densities of the most abundant ions on Titan. After the model was calculated, we figured out that 40 reactions might be further removed from the model without any loss of its accuracy. However, we did not do this because their insignificance was not obvious a priori. Each reaction typically involves four species as reactants and products. Therefore the mean number of reactions per species is fourteen in our model, and this number is adequate for careful presentation of the formation and loss processes Boundary conditions Only two parent species, N 2 and CH 4, are given by their densities at the lower boundary near the surface in our model. Densities of H 2 and CO, which were also fixed at the lower boundary in WA04 and some other models, are calculated by our model, and we do not use these degrees of freedom for better fitting to the observational data. Some species are restricted by their saturated densities near the lower boundary. A chemically passive surface is adopted for all other species, that is, their fluxes are zero at the surface. (All densities except those of N 2,CH 4,H 2, and CO were adopted equal to zero at the surface in WA04, and this means an extremely effective loss on the surface.) According to Strobel (2008), H 2 and CH 4 are subject to slow hydrodynamic escape on Titan. The escape velocities are different for
7 232 V.A. Krasnopolsky / Icarus 201 (2009) H 2 and CH 4 while the escape of N 2 and heavier species does not occur. However, there are some objections to the hydrodynamic escape on Titan, and we will discuss this problem later (see below). Therefore we use the escape velocities of H 2 and CH 4 as free parameters. Then we adopt V H = V H cm s 1 and escape velocities of all CH x, NH x, and OH x equal to V CH4. Here 5000 and cm s 1 are the thermal escape velocities of H 2 and H, respectively. Fluxes of all other neutral species are zero at the upper boundary. Ions above the upper boundary may be swept by the rotating magnetosphere of Saturn. We assume their velocities at the upper boundary equal to half the diffusion velocity: V i = D i /2H i.substitution of the numerical values gives ( 1 V i = ) 1/2 cm s 1. μ0 μ i Our model results in a set of 116 second-order ordinary nonlinear differential equations that are the continuity equations for production, loss, and transport of each species in a one-dimensional spherical atmosphere. Finite difference analogs to these equations are solved using a method described in Krasnopolsky and Cruikshank (1999). The atmosphere is divided in 289 layers with a step gradually increasing from 2.5 km near the surface to 10 km at the upper boundary, that is, each step is 1/8 of the scale height. LCV adopt first derivatives of mixing ratios equal to zero at the lower boundary (surface) and second derivatives of mixing ratios equal to zero at the upper boundary (1350 km). The upper boundary conditions look unusual and generally do not rule out unspecified escape fluxes and/or influxes. The problem of boundary conditions for 1D steady-state models was considered by Krasnopolsky (1995). 4. Results of the model The calculated column reaction rates, CR, and mean altitudes h m for all reactions are given in Tables 3 5. Allthesevaluesare corrected for sphericity, that is, CR = r(h)(1 + h/r 0 ) 2 dh, 0 h m = r(h)(1 + h/r 0 ) 2 hdh 0 / 0 r(h)(1 + h/r 0 ) 2 dh. Table 3 Photodissociation and basic photoionization processes. # Reaction Yield a, reference Column rate b c h m 1 N 2 + hν N + N nm, Krasnopolsky and Cruikshank (1999) N e λ<80 nm d N + + N + e λ<51 nm d, Nicolas et al. (2003) N 2 + e N + N + e Krasnopolsky and Cruikshank (1999) CO+ hν C + O nm, Fox and Black (1989) CH 3 + hν CH 2 + H s 1, Gladstone et al. (1996) CH 4 + hν CH e λ<93 nm, γ 0.5 d CH H + e λ<86 nm, γ 0.5 d CH 3 + H 0.29, Mountetal.(1977), Mount and Moos (1978) CH 2 + H , Au et al. (1993), Wang et al. (2000) CH 2 + 2H CH + H + H C 2 H 2 + hν C 2 H + H 0.3, Wu et al. (2001) C 2 + H 2 0.1, Seki and Okabe (1993) C 2 H 3 + hν C 2 H 2 + H s 1, Fahr et al. (1998) C 2 H 4 + hν C 2 H 2 + H ( nm) C 2 H 2 + 2H , Holland et al. (1997) C 2 H 5 + hν CH 3 + CH s 1, Adachi et al. (1979) C 2 H 6 + hν C 2 H 4 + H Ly α 0.56, Au et al. (1993) C 2 H 4 + 2H , Mount and Moos (1978) C 2 H 2 + 2H CH 4 + CH CH C 3 H 2 + hν C 2 H + CH J C 3 H 3 + hν C 3 H 2 + H s 1, Fahr et al. (1997) C 3 H 4 + hν C 3 H 3 + H 0.56, Ho et al. (1998), Chen et al. (2000) C 3 H 2 + H C 3 H 5 + hν C 3 H 4 + H s 1, Jenkin et al. (1993) C 3 H 6 + hν C 3 H 5 + H ( nm) C 3 H 4 + H , Fahr and Nayak (1996) C 2 H 4 + CH , Samson et al. (1962) C 2 H 3 + CH C 2 H 2 + CH C 3 H 8 + hν C 3 H 6 + H ( nm) C 2 H 6 + CH , Au et al. (1993) C 2 H 5 + CH C 2 H 4 + CH C 4 H 2 + hν C 4 H + H ( nm) C 2 H , Okabe (1981), Smith et al. (1998) C 2 H 2 + C C 4 H C 4 H 4 + hν C 4 H 2 + H s 1, Fahr and Nayak (1996) C 2 H C 4 H 6 + hν C 4 H 4 + H , Hamai and Hirayama (1979) C 2 H C 2 H 4 + C 2 H C 3 H 3 + CH C 4 H 5 + H
8 Titan s photochemistry 233 Table 3 (continued) # Reaction Yield a, reference Column rate b h m c 49 C 4 H 8 + hν C 4 H 6 + 2H ( nm) C 3 H 5 + CH , Samson et al. (1962) C 3 H 4 + CH C 2 H 5 + C 2 H C 4 H 10 + hν C 4 H 8 + H , Au et al. (1993) C 2 H 6 + C 2 H C 3 H 6 + CH 3 + H C 2 H 5 + C 2 H 4 + H C 6 H 2 + hν C 6 H + H s 1, Benilan et al. (1995) C 6 H 6 + hν C 6 H 5 + H 0.8, Rennie et al. (1998) C 6 H 4 + H C 6 H 4 + hν C 6 H 2 + H s 1 (assumed) C 7 H 4 + hν C 4 H 2 + C 3 H s 1 (assumed) C 7 H 8 + hv C 3 H 4 + C 4 H s 1 (assumed) C 8 H 2 + hν 2C 4 H J HCN + hν CN + H Nuth and Glicker (1982), West and Berry (1974) HC 3 N + hν C 3 N + H 0.09, Connors et al. (1974), Benilan et al. (1994) CN + C 2 H 0.05, Seki et al. (1996a), Bruston et al. (1989) CH 3 CN + hν CH 3 + CN Suto and Lee (1985) CH 2 CN + hν CH 2 + CN J C 2 H 3 CN + hν C 2 H 3 + CN Eden et al. (2003) C 2 H 5 CN + hν C 2 H 5 + CN J C 4 H 3 N + hν C 3 H 3 + CN s 1, Bruston et al. (1989) C 4 H 5 N + hν C 3 H 5 + CN J HC 5 N + hν C 4 H + CN J C 5 H 5 N + hν C 4 H 5 + CN s 1, Walker et al. (1989), Wan et al. (2001) 75 C 2 N 2 + hν 2CN Connors et al. (1974) C 4 N 2 + hν C 3 N + CN Connors et al. (1974), Benilan et al. (1996) NH 2 + hν NH + H d NH 3 + hν NH 2 + H s 1d N 2 H 4 + hν N 2 H 3 + H 10 6 s 1, Syage et al. (1992) CH 3 NH 2 + hν CH 3 + NH s 1, Hubin-Franskin et al. (2002) H 2 O + hν OH + H ( nm) O + H , Cheng et al. (1999) CO 2 + hν CO + O Lewis and Carver (1983) H 2 CO + hν CO + H s 1d CO + 2H s 1d HCO + H s 1d CH 2 CO + hν CO + CH s 1, Okabe (1978) CH 3 OH + hν H 2 CO + H 2 d a Yields are taken from WA04. References refer to cross sections. Photolysis frequencies (in s 1 ) are given for species with the main photolytic effect at λ>200 nm and equal to halves of the daytime values above the haze. b Column rates are scaled to the surface and given in cm 2 s 1 : = c Weighted-mean altitude in km. d Cross sections from Table 4 Chemical reactions of neutrals. a # Reaction Rate coefficient Reference Column rate h m 89 H + haze haze 2.1e 1000/T Sekineetal.(2008a) H 2 + haze e 300/T Sekineetal.(2008a) N N + hν s 1 Krasnopolsky and Cruikshank (1995) N + N 2 N + N Lin and Kaufman (1971) N + CH 4 NH + CH e 750/T Takayanagi et al. (1999) CH 2 NH + H e 750/T Takayanagi et al. (1999) N + H 2 NH + H e 880/T Suzuki et al. (1993) N + C 2 H 2 CHCN + H e 270/T Takayanagi et al. (1998) N + C 2 H 4 CH 3 CN + H e 500/T Sato et al. (1999) NH + C 2 H N + C 2 H 6 NH + C 2 H Herron (1999) N + C 3 H 4 C 2 H 3 CN + H k N + C 3 H 6 C 2 H 5 CN + H Herron (1999) N + C 4 H 4 C 4 H 3 N + H k N + C 4 H 6 C 4 H 5 N + H k CH 2 + N 2 CH 2 + N T Ashfold et al. (1981) CH 2 + CH 4 2CH Bohland et al. (1985) CH 2 + H 2 CH 3 + H Langford et al. (1983) C 4 H 2 C 4H 2 + hν 10 s 1 Vuitton et al. (2003) C 4 H 2 + N 2 C 4 H 2 + N Zwier and Allen (1996) C 4 H 2 + C 4H 2 C 6 H 2 + C 2 H Zwier and Allen (1996) (continued on next page)
9 234 V.A. Krasnopolsky / Icarus 201 (2009) Table 4 (continued) # Reaction Rate coefficient Reference Column rate h m 110 C 4 H 2 + C 2H 2 C 6 H 2 + 2H Zwier and Allen (1996) C 4 H 2 + C 2H 4 C 6 H 4 + 2H Zwier and Allen (1996) C 4 H 2 + C 3H 4 C 7 H 4 + H Zwier and Allen (1996) C 4 H 2 + C 3H 6 C 6 H 4 + CH 3 + H Zwier and Allen (1996) C + H 2 + M CH 2 + M /T Husain and Young (1975) CH + H C + H Becker et al. (1989) CH 2 + H CH + H T 0.32 Fulle and Hippler (1997) CH 2 + H + M CH 3 + M , Gladstoneetal.(1996) CH 2 C 2 H 2 + 2H e 400/T Tsang and Hampson (1986) CH 3 + H + M CH 4 + M e 260/T, Pratt and Wood (1984) CH 3 + CH 2 C 2 H 4 + H Tsang and Hampson (1986) CH 3 + M C 2 H 6 + M T 8.75 e 980/T,10 14 LCV CH 4 + CH C 2 H 4 + H /Te 36/T Canosa et al. (1997) C 2 + CH 4 C 3 H 3 + H T 0.42 e 13/T Canosa et al. (2007) C 2 + C 2 H 2 C 4 H + H T 1.14 e 77/T Canosa et al. (2007) C 2 + C 2 H 4 C 4 H 3 + H T 0.93 e 58/T Canosa et al. (2007) C 2 + C 2 H 6 C 3 H 3 + CH T 0.94 e 44/T Canosa et al. (2007) C 2 + C 3 H 8 C 3 H 2 + C 2 H T 1.31 e 94/T Canosa et al. (2007) C 2 H + CH 4 C 2 H 2 + CH e 490/T Opansky and Leone (1996a) C 4 H + CH 4 C 4 H 2 + CH 3 k C 2 H 2 + H + M C 2 H 3 + M e 740/T, Baulch et al. (1992) C 2 H 2 + CH C 3 H 2 + H T 0.23 e 16/T Canosa et al. (1997) C 2 H 2 + C 2 H C 4 H 2 + H T 1.8 e 474/T Opansky and Leone (1996b) C 2 H 2 + C 4 H C 6 H 2 + H T 1.06 e 66/T Berteloite et al. (2008) C 2 H 3 + H C 2 H 2 + H Monks et al. (1995) C 2 H 3 + H + M C 2 H 4 + M T 2.2, Monks et al. (1995) C 2 H 3 + CH 2 C 2 H 2 + CH Tsang and Hampson (1986) C 2 H 3 + CH 3 C 2 H 2 + CH Fahr et al. (1999) C 2 H 3 + CH 3 + M C 3 H 6 + M ,10 15 Assumed C 2 H 3 + C 2 H 2 C 4 H 4 + H e 2520/T Fahr and Stein (1989) C 2 H 3 + C 2 H 2 + M C 4 H 5 + M T 5.8 e 2364/T,10 18 Wang and Frenklach (1994) C 2 H 3 C 2 H 4 + C 2 H Laufer and Fahr (2004) C 2 H 3 + M C 4 H 6 + M k C 2 H 4 + H + M C 2 H 5 + M e 380/T, Baulch et al. (1994) C 2 H 4 + C C 2 H + CH T 0.07 Chastaing et al. (1999) C 2 H 4 + CH C 3 H 4 + H See text C 2 H 4 + C 2 H C 4 H 4 + H Vakhtin et al. (2001) C 2 H 4 + C 4 H C 2 H 3 + C 4 H T 0.4 e 10/T Berteloite et al. (2008) C 6 H 4 + H T 0.4 e 10/T Berteloite et al. (2008) C 2 H 5 + H CH 3 + CH Pimentel et al. (2004) C 2 H 4 + H Tsang and Hampson (1986) C 2 H 5 + CH 2 C 2 H 4 + CH Tsang and Hampson (1986) C 2 H 5 + CH 3 C 2 H 4 + CH Baulch et al. (1994) C 2 H 5 + CH 3 + M C 3 H 8 + M T 16.1 e 1900/T, Laufer et al. (1983) C 2 H 5 + C 2 H 3 + M C 4 H 8 + M k C 2 H 5 + C 2 H 3 C 2 H 4 + C 2 H Laufer and Fahr (2004) C 2 H 5 C 2 H 4 + C 2 H Baulch et al. (1992) C 2 H 5 + M C 4 H 10 + M T 6.4 e 300/T, Laufer et al. (1983) C 2 H 6 + CH C 2 H 4 + CH T 0.86 e 53/T Canosa et al. (1997) C 3 H 6 + H T 0.86 e 53/T Galland et al. (2003) C 2 H 6 + C 2 H C 2 H 5 + C 2 H Opansky and Leone (1996b) C 2 H 6 + C 4 H C 2 H 5 + C 4 H T 1.24 e 26/T Berteloite et al. (2008) C 3 H 2 + H + M C 3 H 3 + M , Laufer et al. (1983) C 3 H 3 + H + M C 3 H 4 + M , Laufer et al. (1983) C 3 H 3 + CH 3 + M C 4 H 6 + M k C 3 H 3 + M C 6 H 6 + M 10 26, Assumed C 3 H 4 + H C 2 H 2 + CH e 960/T Fernandes et al. (2005) C 3 H 4 + H + M C 3 H 5 + M T 2 e 1220/T, Whytock et al. (1976) C 3 H 4 + C 2 H C 3 H 3 + C 2 H T 0.3 Carty et al. (2001) C 3 H 5 + H C 3 H 4 + H Tsang (1991) C 3 H 5 + H + M C 3 H 6 + M 10 24, Hanning-Lee and Pilling (1992) C 3 H 5 + CH 2 C 4 H 6 + H Tsang (1991) C 3 H 5 + CH 3 C 3 H 4 + CH T 0.32 e 66/T Tsang (1991) C 3 H 5 + CH 3 + M C 4 H 8 + M k C 3 H 5 + C 2 H 3 C 3 H 6 + C 2 H Tsang (1991) C 3 H 4 + C 2 H Tsang (1991) C 3 H 5 + C 2 H 5 C 3 H 6 + C 2 H e 66/T Tsang (1991) C 3 H 4 + C 2 H e 66/T Tsang (1991) C 3 H 6 + H + M C 3 H 7 + M , Laufer et al. (1983) C 3 H 6 + CH 2 C 3 H 5 + CH e 2660/T Tsang (1991) C 3 H 7 + H C 3 H 6 + H Tsang (1988) C 3 H 7 + H + M C 3 H 8 + M k
10 Titan s photochemistry 235 Table 4 (continued) # Reaction Rate coefficient Reference Column rate h m 182 C 3 H 7 + CH 2 C 3 H 6 + CH Tsang (1988) C 2 H 5 + C 2 H Tsang (1988) C 3 H 7 + CH 3 C 3 H 6 + CH T 0.32 Tsang (1988) C 3 H 7 + C 2 H 3 C 3 H 8 + C 2 H Tsang (1988) C 3 H 6 + C 2 H Tsang (1988) C 3 H 7 + C 2 H 5 C 3 H 8 + C 2 H Tsang (1988) C 3 H 6 + C 2 H Tsang (1988) C 3 H 8 + C 2 H C 3 H 7 + C 2 H Laufer and Fahr (2004) C 4 H + C 3 H 4 C 7 H 4 + H T 0.82 e 47/T Berteloite et al. (2008) C 4 H + C 3 H 8 C 7 H 8 + H T 1.35 e 56/T Berteloite et al. (2008) C 4 H 2 + H + M C 4 H 3 + M 10 28, Gladstone et al. (1996) C 4 H 2 + C 2 H C 6 H 2 + H Frank and Just (1980) C 4 H 2 + C 4 H C 8 H 2 + H k C 4 H 3 + H C 4 H 2 + H Wang and Frenklach (1997) C 4 H 3 + CH 3 C 4 H 2 + CH 4 k C 4 H 4 + H + M C 4 H 5 + M 10 7 T 7 e 1390/T, Schwanebeck and Warnatz (1975) C 4 H 5 + H C 4 H 4 + H 2 k C 4 H 5 + H + M C 4 H 6 + M k C 4 H 5 + C 2 H 2 C 6 H 6 + H T 1.18 e 1880/T Westmoreland et al. (1989) 201 C 6 H 2 + C 2 H C 8 H 2 + H e 31/T WA C 6 H 5 + H+M C 6 H 6 + M k C 6 H 6 + C 4 H 5 C 10 H 10 + H e 3220/T Fascella et al. (2004) C 6 H 6 + C 6 H 5 C 12 H 10 + H e 2010/T Fahr and Stein (1989) C 6 H 2 + C 4 H polymer k C 6 H 4 + C 4 H polymer k C 4 H 2 + C 6 H polymer k C 6 H 2 + C 6 H polymer k C 8 H 2 + C 6 H polymer k C 6 H 5 + C 2 H 2 polymer T 0.2 e 2520/T Wang and Frenklach (1994) C 6 H 5 + C 2 H 2 + M polymer T 4 e 400/T, Wang and Frenklach (1994) C 6 H 6 + C 2 H polymer Wang and Frenklach (1994) C 6 H 6 + C 6 H 5 polymer e 2170/T Park et al. (1999) C 8 H 2 + C 2 H polymer k C 8 H 2 + C 4 H polymer k C 7 H 4 + C 2 H polymer k C 7 H 8 + C 2 H polymer k N + CH CN + H T 0.1 Brownsword et al. (1996) N + CH 2 HCN + H e 250/T Yung et al. (1984) N + CH 3 H 2 CN + H e 420/T Marston et al. (1989) N + H 2 CN HCN + NH e 200/T Nesbittetal.(1990) H + H 2 CN HCN + H 2 2k 221 Nesbittetal.(1990) H 2 CN + HCN polymer e 900/T LCV H 2 CN + HC 3 N polymer e 900/T LCV N + C 2 N 2CN Whyte and Phillips (1983a) CH 2 CN + H + M CH 3 CN + M Assumed N + C 2 H 3 CH 2 CN + H Payneetal.(1996) C 2 H 2 + NH Payneetal.(1996) N + C 2 H 3 + M CH 3 CN + M Payneetal.(1996) N + C 2 H 5 C 2 H 4 + NH Stief et al. (1995) H 2 CN + CH Stief et al. (1995) CH 3 CN + H Assumed NH 3 + C 2 H See text N + CHCN C 2 N 2 + H Yung (1987) N + C 3 H 6 C 2 H 5 CN + H Herron (1999) NH 3 + C 3 H 3 k 235 See text N + C 4 H 4 C 4 H 3 N + H Assumed NH + H N + H T 1.5 e 100/T Adam et al. (2005) NH + N N 2 + H Hack et al. (1994) NH + NH NH 2 + N T 2.9 e 1000/T Zu et al. (1997) NH + CH 3 CH 4 + N Lellouch et al. (1994) NH + CH 2 CH 3 + N k NH + C 2 H 3 C 2 H 4 + N k NH + C 2 H 5 C 2 H 6 + N k NH + C 2 H 2 CHCN + H T 1.07 Mullen and Smith (2005) NH 2 + H + M NH 3 + M , Schofield (1973) NH 2 + N N 2 + 2H Whyte and Phillips (1983b) NH 2 + NH 2 + M N 2 H 4 + M T 3.9, Fagerstrom et al. (1995) NH 2 + NH 2 N 2 H 2 + H Stothard et al. (1995) NH 2 + CH 3 + M CH 3 NH 2 + M T 3.85, Jodkowski et al. (1995) N 2 H 2 + H N 2 H + H T 2.6 e 115/T Linder et al. (1996) N 2 H 2 + NH 2 NH 3 + N 2 H T 4.05 e 810/T Linder et al. (1996) N 2 H + H N 2 + H Bozzelli and Dean (1995) (continued on next page)
11 236 V.A. Krasnopolsky / Icarus 201 (2009) Table 4 (continued) # Reaction Rate coefficient Reference Column rate h m 254 N 2 H 4 + H N 2 H 3 + H e 1200/T Stief and Payne (1976) N 2 H 3 + H 2NH von Gehring et al. (1971) CH 2 NH + H CH 3 + NH Assumed CN + CH 4 HCN + CH e 675/T Sims et al. (1993) CN + C 2 H 2 HC 3 N + H T 0.52 e 19/T Sims et al. (1993) CN + C 2 H 4 HCN + C 2 H T 0.7 e 31/T Sims et al. (1993) C 2 H 3 CN + H T 0.7 e 31/T Monks et al. (1993) CN + C 2 H 6 HCN + C 2 H T 0.22 e 58/T Sims et al. (1993) CN + C 3 H 4 HCN + C 3 H Carty et al. (2001) CN + C 4 H 2 C 5 HN + H Seki et al. (1996b) CN + HCN C 2 N 2 + H T 1.7 e 770/T Yang et al. (1992) CN + HC 3 N C 4 N 2 + H Halpern et al. (1989) CN + CH 3 CN C 2 N 2 + CH e 1190/T Zabarnick and Lin (1989) CN + CH 2 CN C 2 N 2 + CH Assumed CN + C 2 H 3 CN polymer e 103/T Butterfield et al. (1993) CN + C 2 N 2 polymer T 2.7 e 325/T Yang et al. (1992) CN + C 4 N 2 polymer Seki et al. (1996b) HC 3 N + C 6 H 5 + M polymer + M T 4 e 400/T, LCV C 4 H + HC 3 N polymer T 1.8 e 474/T LCV C 3 N + C 4 H 2 polymer T 1.8 e 474/T LCV C 4 H 3 + HC 3 N + M polymer + M e 740/T LCV CH + HCN CHCN + H e 500/T Zabarnick et al. (1991) CHCN + H C 2 N + H Osamura and Petrie (2004) CHCN + CH 3 C 2 H 3 CN + H Osamura and Petrie (2004) C 2 N + CH 4 CHCN + CH Zhu et al. (2003) C 2 N + C 2 H 6 C 4 H 3 N + H 2 + H Zhu et al. (2003) C 3 N + CH 4 HC 3 N + CH Yung (1987) C 3 N + C 2 H 6 C 5 H 5 N + H See text C 3 N + HCN C 4 N 2 + H Assumed CH 2 CN + CN C 2 N 2 + CH See text HC 3 N + C 2 H HC 5 N + H Assumed HCN + C 2 H 3 C 2 H 3 CN + H e 900/T Monks et al. (1993) CH 3 CN + H CN + CH e 1500/T Jamieson et al. (1970) O + CH 3 H 2 CO + H Hack et al. (2005) CO + H 2 + H Hack et al. (2005) O + C 2 H 4 CH 3 CO + H T 1.9 e 92/T Baulch et al. (1994) CH 3 + HCO k /0.3 Baulch et al. (1994) H 2 CO + CH 2 k /0.3 Endo et al. (1986) OH + CH 3 CH 2 + H 2O T 0.91 e 275/T Pereira et al. (1997) H 2 CO + H T 0.12 e 209/T Pereira et al. (1997) OH + CH 3 + M CH 3 OH + M T 6.21 e 671/T, Pereira et al. (1997) OH + CH 4 H 2 O + CH e 1775/T Sander et al. (2006) OH + CO CO 2 + H e 176/T Frost et al. (1993) OH + H 2 H 2 O + H e 1530/T Orkin et al. (2006) OH + C 2 H 2 + M CH 3 CO + M , Sander et al. (2006) OH + C 2 H 4 + M H 5 C 2 O + M T 4.5,10 16 Sander et al. (2006) HCO + H CO + H Friedrichs et al. (2002) HCO + CH 3 CH 4 + CO Krasnoperov et al. (2005) CH 3 CO + H HCO + CH Bartels et al. (1991) CH 2 CO + H Ohmori et al. (1990) CH 3 CO + CH 3 CH 2 CO + CH Hassinen et al. (1990) CO + C 2 H Adachi et al. (1981) CH 2 CO + H CO + CH T 1.45 e 1400/T Senosiain et al. (2006) CH 3 COH + H CH 3 CO + H T 1.16 e 1210/T Baulch et al. (1992) H 5 C 2 O + H CH 3 COH + H Edelbuttel et al. (1992) a Reaction probabilities are given for r89 and r90. Rate coefficients of two- and three-body reactions are in cm 3 s 1 and cm 6 s 1, respectively. Densities in cm 3 for high-pressure limits are given for three-body reactions at 170 K. Column rates are scaled to the surface and given in cm 2 s 1, mean altitudes h m are in km. Here r(h) is the reaction rate, and R 0 = 2575 km is the surface radius. Column reaction rates and their mean altitudes are helpful in analyses of photochemical balances in the model. Designations for comparison of the model results with the observations are shown in Fig. 5. Similar to Table 2, the model data refer to species that have been observed, to those derived by the ion chemistry fitting to the INMS ion spectrum (Vuitton et al., 2007), and to species with no observational data. Of all pre- Cassini observations, we use only the Voyager UVS occultation data (Vervacketal.,2004), millimeter observations of HCN, HC 3 N, and CH 3 CN (Marten et al., 2002; Gurwell, 2004; Tanguy et al., 1990), and the ISO observation of H 2 O(Coustenis et al., 1998). We consider the Cassini CIRS observations as superior to the Voyager IRIS because of the better properties of CIRS. Results of the Cassini UVIS occultations (Shemansky et al., 2005) are given as number densities n i of a species i at some altitudes. Densities of N 2 are not available in Shemansky et al. (2005), and the species mixing ratios are obtained from the retrieved densities using a relationship ( ) ( ni [CH 4 ] f i = [CH 4 ] UVIS [N 2 ]+[CH 4 ] ). model
12 Titan s photochemistry 237 Table 5 Ion reactions. a # Reaction Rate coefficient Column rate h m 309 N 2 + cosmic rays N e N + + N + e N + N N 2 + e N e + e N + + N + e + e N + N + e N 2 + p N e + p N + + N + e + p N + N + p N CH 4 CH H + N CH H 2 + N N H 2 N 2 H + + H N C 2H 4 C 2 H H + N C 2 H H 2 + N HCN + + HCN + H HCNH + + HCN + H N 2 H + + C 2 H N + + CH 4 CH NH CH N HCNH + + H HCN + + H 2 + H CH CH 4 CH CH CH CH 4 C 2 H H CH CH 4 C 2 H H C 2 H H N 2 H + + CH 4 CH N HCN + + CH 4 C 2 H NH HCNH + + CH C 2 H CH 4 C 3 H H C 2 H CH 4 C 3 H H C 3 H H CH C 2H 2 C 2 H CH CH C 2H 4 C 2 H CH CH C 2H 6 C 2 H CH 4 + H C 2 H CH CH HCN HCNH+ + CH C 2 H C 2H 2 C 3 H CH C 4 H H C 2 H C 2H 2 C 3 H CH C 4 H H C 2 H C 2H 4 C 3 H CH C 2 H HCN HCNH+ + C 2 H C 2 H HC 3N HC 3 NH + + C 2 H C 2 H CH 3CN CH 3 CNH + + C 2 H C 2 H C 2H 3 CN C 2 H 3 CNH + + C 2 H C 2 H CH 2NH CH 2 NH C 2H C 2 H C 4H 3 N C 4 H 3 NH + + C 2 H C 2 H C 4H 2 C 4 H C 2H C 2 H NH 3 NH C 2H C 3 H HCN C 4H 3 NH + + hν C 3 H C 2H 4 C 5 H H C 5 H hν C 3 H C 2H 2 C 5 H H C 3 H C 2H 2 C 5 H H C 4 H C 2H 2 C 6 H hν C 4 H C 2H 4 C 6 H H C 4 H C 2H 2 C 6 H H C 5 H C 2H 2 C 7 H hν C 6 H CH 4 C 7 H H C 6 H H 2 C 6 H hν C 7 H C 2H 4 C 9 H 11 + hν C 7 H C 3H 4 C 10 H hν HCNH + + C 4 H 2 C 4 H HCN HCNH + + HC 3 N HC 3 NH + + HCN HCNH + + CH 3 CN CH 3 CNH + + HCN HCNH + + C 2 H 3 CN C 2 H 3 CNH + + HCN HCNH + + CH 2 NH CH 2 NH HCN HCNH + + C 4 H 3 N C 4 H 3 NH + + HCN HCNH + + NH 3 NH HCN HCNH + + e HCN + H , HC 3 NH + + e C 3 N + H , C 2 H 3 CNH + + e HCN + C 2 H , (continued on next page)
13 238 V.A. Krasnopolsky / Icarus 201 (2009) Table 5 (continued) # Reaction Rate coefficient Column rate h m 381 CH 3 CNH + + e CH 2 CN + H , C 4 H 3 NH + + e HCN + C 3 H , C 2 H e C 2H 6 + H , C 3 H e C 3H 3 + H , CH 2 NH e CH 2 + NH , C 4 H e C 3H 2 + CH , C 5 H e C 4H 4 + CH , NH e NH 2 + 2H , C 6 H e C 6H 4 + H , C 6 H e C 6H 6 + H 10 6, C 7 H e C 6H 6 + CH 10 6, C 5 H e C 2H 4 + C 3 H , C 3 H e C 3H 2 + H , N e N + N , N 2 H + + e N + NH 10 6, HCN + + e CN + H , CH e C + 2H , CH e CH + 2H , CH e CH 3 + H , CH e CH 3 + 2H , C 2 H e C 2H + H , C 2 H e C 2H + 2H , C 2 H e C 2H 2 + 2H , C 2 H e C 2H 4 + H , C 3 H e C 3H 3 + H , C 4 H e C 3H 4 + CH , C 4 H 4 + H b C 4 H C 6H 6 C 10 H hν b C 4 H C 6H 4 C 10 H hν b C 6 H C 4H 4 C 10 H hν b C 6 H C 4H 2 C 10 H hν b C 7 H C 3H 2 C 10 H hν b C 10 H e C 10H 8 + H 10 6, b C 9 H e C 9H 8 + H 2 + H 10 6, b C 10 H e C 10H 8 + H 2 + H 10 6, a Rate coefficients are from McEwan and Anicich (2007); recombination rate coefficients are from Woodall et al. (2007). TheyareequaltoA(300/T e ) α ; A and α are given. All rate coefficients are in cm 3 s 1. Column rates are scaled to the surface and given in cm 2 s 1, mean altitudes h m are in km. b Assumed, see text. Fig. 5. Designations for comparison of the model results with the observations. For the Voyager UVS occultations (Vervack et al., 2004), we compare the observed densities of a species with the N 2 and CH 4 densities from the model. Absorption of the solar UV and EUV radiation in the atmosphere of Titan is calculated interactively in the model. The absorption curves for some wavelengths are shown in Fig. 6. The absorption is very strong at λ<80 nm due to N 2.AbsorptionbyCH 4 is mostly Fig. 6. Absorption of the solar UV and EUV radiation in the atmosphere of Titan. Numbers near the curves give wavelengths in nm. at nm, and photochemical products absorb at nm. Absorption by gases becomes weak at λ>200 nm, and the haze extinction dominates in this spectral range. Mean altitudes of the absorption by the parent species are 1010 km for N 2 and810kmforch 4.C 2 H 2,C 2 H 4,C 4 H 2, and HC 3 N are the most effective absorbers among the photochemical prod-
14 Titan s photochemistry 239 ucts (Table 3). For example, the absorption by acetylene peaks at 380 km Hydrodynamic escape of CH 4 and H 2? Yelle et al. (2006) concluded from the INMS data in the first Cassini encounter that either very high eddy diffusion K cm 2 s 1 or high escape rate of methane cm 2 s 1 were required to explain the observed profiles of N 2 and CH 4. Strobel (2008) applied the equations for hydrodynamic escape to H 2 and CH 4 on Titan and concluded that a slow hydrodynamic escape of these species occurs. Yelle et al. (2008) analyzed the mean profiles of N 2, CH 4, and Ar from all encounters and derived K = (2 5) 10 7 cm 2 s 1 and hydrodynamic escape of CH 4 at the diffusion limited rate of cm 2 s 1. Cui et al. (2008) studied the INMS data for H 2 and obtained the escape rate of cm 2 s 1 that exceeds thermal escape by a factor of 3 and requires hydrodynamic escape. However, direct Monte Carlo simulations by Tucker and Johnson (2008) do not support the slow hydrodynamic escape. Furthermore, Sittler et al. (2008) reported no evidence for the methane group ions in Saturn s outer magnetosphere. The observed profiles of N 2 and CH 4 in Yelle et al. (2008) are rather close to those for diffusive equilibrium. For example, the N 2 densities at 1000, 1250 and 1500 km are , and cm 3, respectively, and those of CH 4 are , and cm 3, respectively. Then ln(n 1000km 1500 km 2 /N2 ) ln(ch 1000km 1500 km 4 /CH4 ) = 1.56, which is close to 28/16 = 1.75 required for the diffusive equilibrium, and corrections by 3% and of the proper signs to the densities could fit exactly to 28/16. Actually the conclusions in Yelle et al. (2008) are based on comparison of the INMS observations in the upper atmosphere with the Huygens GCMS measurements in the lower atmosphere (Niemann et al., 2005). Comparing the data for Ar, they found a position of homopause and the proper eddy diffusion. Applying this eddy diffusion to the GCMS data on CH 4 and neglecting the effects of chemistry, they found f CH4 = at 700 km and the flow of CH 4 requiredtofittheinmsobservations. The chemical lifetime of methane is longer than its mixing time above the troposphere, and the neglect of chemistry looks reasonable. However, the methane density varies from the troposphere to the upper atmosphere by so many orders of magnitude, that accumulation of the weak chemical effect to a significant value of something like 2 is generally not ruled out. Mean temperatures in the intervals of and km may be calculated from the N 2 densities using the standard relationship n = a exp(γ Mm/rkT). Hereγ and k are the gravitational and Boltzmann constants, M and m are the masses of Titan and the N 2 molecule. The calculated temperatures are 148 and 152 K, respectively, and do not show a decrease with altitude typical of hydrodynamic escape. Therefore the existence of hydrodynamic escape on Titan looks rather uncertain. The different escape velocities for H 2 and CH 4 and no escape of N 2 (Strobel, 2008) are also disputable. We may check the problem with our model. The best fit to theinmsobservationsofch 4 and H 2 is for their escape velocities of 2000 and cm s 1 that result in the escape rates of and cm 2 s 1, respectively (Fig. 7, upper left panel). These values are close to those in Strobel (2008), Yelle et al. (2008), and Cui et al. (2008). Assuming thermal escape of H 2 with V H2 = 5000 cm s 1 and no escape of CH 4,wemeetsignificant disagreements between the model and the INMS observations (Fig. 7). Therefore, our model supports the slow hydrodynamic escape of the light species from Titan Hydrocarbons and H 2 The hydrocarbon chemistry is initiated by photolysis of methane which is one of two parent gases on Titan. The CH 3 H bond energy corresponds to the dissociation limit at 277 nm. However, the effective dissociation starts at 140 nm because of the high symmetry of the CH 4 molecule. Photochemical products of CH 4 are less symmetric, and their photolyses begin at longer wavelengths. For example, the molecule of acetylene C 2 H 2 is linear, that is, highly symmetric but not so high as CH 4 which is almost spherical, and the effective dissociation of C 2 H 2 starts at 200 nm, though the C 2 H H bond energy exceeds that of CH 3 H. More and more solar photons are involved into the atmospheric chemistry by the photochemical products of CH 4, and this makes the hydrocarbon chemistry complicated and interesting. To avoid possible confusion, we define total production or loss as sum of all production or loss processes for a species, and column production or loss in a specific process as integrated production or loss in the whole atmosphere or above a specified level Methane CH 4 The calculated profile of CH 4 (Fig. 7) perfectly matches the Huygens GCMS observations (Niemann et al., 2005) below150km. The profile corresponds to the saturation conditions at 8 to 30 km and is constant below and above this interval. Small deviations from the constant mixing ratio start at 400 km. The calculated densities agree with the INMS observations (Yelle et al., 2008, 2006; Waite et al., 2005, 2007), mostly due to the hydrodynamic escape adopted in the model. The upward flux of CH 4 at the surface is equal to cm 2 s 1 (Table 6). Photochemical production of CH 4 is smaller by an order of magnitude and mostly due to CH 3 + C 4 H 3 C 4 H 2 + CH 4. Losses of methane are primarily caused by photolysis which results in CH 4 + hν CH 2 + H 2, 2H (16%), CH 3 + H (7%), CH + H + H 2 (2%), CH 3 + hν CH 2 + H, CH 2 + N 2 CH 2 + N 2, CH 2 + CH 4 2CH 3 (4%), CH 2 + H CH + H 2, CH 4 + CH C 2 H 4 + H (19%). The numbers in parentheses indicate shares in the total photochemical (without the escape) loss of CH 4. Other losses of CH 4 are related to photolyses of acetylene and diacetylene: C 2 H 2 + hν C 2 H + H, C 2 + H 2, C 4 H 2 + hν C 2 H 2 + C 2, C 2 H + CH 4 C 2 H 2 + CH 3 (19%), C 2 + CH 4 C 3 H 3 + H (11%).
15 240 V.A. Krasnopolsky / Icarus 201 (2009) Fig. 7. Calculated vertical profiles of CH 4,H 2,andthemajorC 2 H x hydrocarbons. Box in the upper left panel shows H 2 observed by the Voyager IRIS (Samuelson et al., 1981; Courtin et al., 1995). Solid line in the upper left panel is the CH 4 profile measured by GCMS (Niemann et al., 2005), and thin lines show the H 2 and CH 4 profiles without hydrodynamic escape. The other designations are from Fig. 5. Table 6 Column production and loss rates, escape and precipitation rates, and lifetimes for some species. a Species Production (cm 2 s 1 ) Loss (cm 2 s 1 ) Flux (cm 2 s 1 ) Flux (g cm 2 Byr 1 ) Lifetime (yr) H H b CH C 2 H C 2 H C 2 H C 3 H C 3 H c C 4 H C 4 H C 4 H C 4 H C 4 H C 6 H HCN HC 3 N CH 3 CN C 2 H 3 CN C 2 H 5 CN C 2 N H 2 O d CO CO C x H haze y C x H y N haze Ion haze a All values are scaled to the surface. Lifetime is a ratio of column abundance to column production or loss = Flux is positive for escape and negative for precipitation. Superscript haze refers to the polymerization processes. Some of them result in C x H y N z > 1 but their rates are low. Ion haze means formation of the haze by recombination of heavy ions. b Flux at the surface is cm 2 s 1 and gcm 2 Byr 1. c Excitation and quenching of C 4 H 2 without chemical reaction do not included. d Production by meteorites is cm 2 s 1. The direct photolysis of CH 4 is effective up to 140 nm while the photolyses of C 2 H 2 and C 4 H 2 extend beyond 200 nm and result in indirect photolysis of CH 4 in the above processes (Yung et al., 1984). Reactions with N,C 4 H, CN, C 2 N, and C 3 N add 16% and the ion reactions 7% to the loss of methane. Hydrodynamic escape is another significant loss process Molecular hydrogen H 2 Unlike in some previous models (Toublanc et al., 1995; Lara et al., 1996; Banaszkiewicz et al., 2000; WA04), H 2 is not fixed at the lower boundary, and its densities are calculated by the model throughout the atmosphere. The calculated vertical profile of H 2 (Fig. 7) agrees with the INMS data (Yelle et al., 2006; Cui et al., 2008) in the upper atmosphere and with the Voyager IRIS observations (Samuelson et al., 1981; Courtin et al., 1995) in the stratosphere. Hydrodynamic escape of H 2 (Strobel, 2008) lowers the H 2 densities in the upper atmosphere and facilitates the agreement with the INMS measurements. A step-like increase in the H 2 mixing ratio from below 300 km to above 500 km reflects a peak in the H 2 production at 550 km and very low production below 300 km. Loss of H 2 in reactions with the cosmic ray induced ions below 100 km adds to this feature. Loss of H 2 is much smaller than its production (Table 6) which is spent to supply the escape flux of cm 2 s 1.Photolyses of hydrocarbons provide 34% of the column production, with the major contribution from CH 4,C 2 H 4,C 2 H 2, and C 4 H 4.Reactions with radicals CH 2 + H CH + H 2 (15%), C 4 H 3 + H C 4 H 2 + H 2 (26%), are the important sources of H 2. The remaining 25% of H 2 are formed by reactions of other radicals and ions. Formation of H 2 in the heterogeneous reaction of H with the haze (Sekine et al., 2008a) is insignificant.
16 Titan s photochemistry 241 Fig. 8. Calculated vertical profiles of C x > 2 H y hydrocarbons. See designations in Fig. 5. Measurements refer to the closest curve. CIRS observations in the upper left panel refer to C 3 H 4 and C 3 H Acetylene C 2 H 2 The calculated profile of C 2 H 2 (Fig. 7) is near the INMS values and the UVIS occultation measurements at km. It agrees with the CIRS nadir observations but does not fit the CIRS limb observations that revealed the constant C 2 H 2 mixing ratio at 100 to 400 km. The CIRS limb data and the Voyager UVS occultations form a continuous profile with a break at km. The steep slope of the Voyager UVS profile means a strong downward flux of C 2 H 2 that disappears below 450 km where the profile becomes constant. Therefore a strong sink should exist at the break point while actually both the production and loss of C 2 H 2 peak near this point (at 450 km) in our model. 4% of the total production of cm 2 s 1 (Table 6) is lost by condensation below 100 km. Photolyses of C 2 H 4 and C 4 H 2x (x = 1, 2, 3) give 38% of the total production, and the reactions of C 2 HwithCH 4 and C 2 H 6 add 32%. Another important source of acetylene is the reactions C 2 H 3 + H C 2 H 2 + H 2, C 3 H 4 + H C 2 H 2 + CH 3, which gives 26% of the column production in our model. The loss of C 2 H 2 by photolysis is equal to 46% of its production. Other losses are from reactions with radicals H, CH, C 2,C 2 H, C 4 H, CN, N, and NH (48%) and ions (2%) Ethylene C 2 H 4 The calculated profile of C 2 H 4 perfectly fits the INMS data near 1000 km (Waite et al., 2007), the value from Vuitton et al. (2007) and the UVIS data below 1000 km, agrees with the Voyager UVS observations at km and is close to the CIRS nadir value. However, the profile does not support the decrease in the C 2 H 4 mixing ratio at km observed with CIRS at the limb. Formation of ethylene on Titan is mostly due to CH 4 + CH C 2 H 4 + H. This reaction gives 75% of the total production of C 2 H 4 which is equal to cm 2 s 1 (Table 6). The rate of this reaction is near the rate of the direct photolysis of methane. Therefore, the quantum yield of the direct photolysis of CH 4 is 2, and the yield of C 2 H 4 in this photolysis is 1 (Krasnopolsky and Cruikshank, 1995, 1999). Significant productions of C 2 H 4 are in photolysis of C 4 H 6 (6%) and in CH 3 + CH 2 C 2 H 4 + H (4%). The ion reactions add 9% to the production of ethylene. Photolysis dominates in the loss of C 2 H 4 (73%). Other significant loss processes are the reactions with C 2 H (6%), CN (9%), and with H, CH, C 2, and C 4 H (6%). Ethylene does not condense in our basic model (Table 6). The total production and loss peak at 780 and 580 km, respectively, and are perfectly balanced below 400 km Ethane C 2 H 6 The calculated profile of C 2 H 6 (Fig. 7) agrees with the INMS measurements and the CIRS nadir value. Comparison with the UVIS occultations and the CIRS limb data is not so good. Almost all ethane is formed in the reaction CH 3 + CH 3 + M C 2 H 6 + M. Ethane is lost (Table 6) in condensation and precipitation (48%). Other loss processes are the reactions with C 2 H(18%),C 2, CH, and C 4 H (8%), and C 3 N(10%),C 2 N (5%), and CN (7% ). Loss of C 2 H 6 by photolysis is low (2%). Vinatier et al. (2007) calculated a vertical profile of C 2 H 6 at 50 to 500 km assuming no production and loss in this altitude range. However, both the production and loss of C 2 H 6 peak near 400 km in our model, and this assumption is not valid Other hydrocarbons Vertical profiles of some other hydrocarbons are shown in Fig. 8. ReactionC 2 H 4 + CH C 3 H 4 + H is the main high-altitude source of methylacetylene. Its rate coefficient relative to a similar reaction with CH 4 is very different in measurements by Canosa et al. (1997) and McKee et al. (2003). We choose a value of
17 242 V.A. Krasnopolsky / Icarus 201 (2009) cm 3 s 1 which fits the INMS observations of C 3 H 4.The C 3 H 4 profile below 600 km reflects a balance between the production in C 3 H 3 + H + M and loss in the reaction with H. The INMS data for propylene C 3 H 6, and propane C 3 H 8 are in excellent agreement with the model abundances. The model results for diacetylene C 4 H 2 near 1000 km exceed the INMS value but are close to the value from the ion chemistry fitting to the observed ion mass spectrum (Vuitton et al., 2007) and the UVIS occultation data. Benzene C 6 H 6 near 1000 km agrees with the INMS measurements, and both benzene and triacetylene C 6 H 2 are close to the estimates by Vuitton et al. (2007). Our model agrees with the CIRS nadir measurements of C 3 H 4,C 3 H 8,C 4 H 2, and C 6 H 6, while the vertical profiles from the CIRS limb observations are not as steep as the model profiles. The estimates of Vuitton et al. (2007) are near the model values for tetraacetylene C 8 H 2 but differ for C 7 H 4 and C 7 H 8. Contributions of various reactions to production and loss of each species may be analyzed using the data of Tables 3 6. Indirect recombination of H 2 and CH 4 is the most effective via C 4 H 2 : C 4 H 2 + H + M C 4 H 3 + M, C 4 H 3 + H C 4 H 2 + H 2, C 4 H 3 + CH 3 C 4 H 2 + CH 4, Net H + H H 2, H + CH 3 CH 4. Diacetylene recycles in these processes with no net loss. Photoexcitation of C 4 H 2 is almost completely quenched to C 4H 2, and the column production/loss of cm 2 s 1 corrected for this indirect recombination and photoexcitation/quenching of C 4 H 2 is equal to cm 2 s 1.ProductionofC 4 H 2 is mostly in the reactions C 2 H 2 + C 2 H C 4 H 2 + H (76%), C 4 H 4 + hν C 4 H 2 + H 2 (17%). Photolysis (54%) and polymerization (40%) are the basic loss processes. Benzene mostly recycles in the main photolysis branch in our model: C 6 H 6 + hν C 6 H 5 + H, C 6 H 5 + H + M C 6 H 6 + M. The second photolysis branch that yields C 6 H 4 + H 2 results in removal of benzene. C 6 H 6 is produced by the reactions: C 6 H e C 6H 6 + H (20%), C 7 H e C 6H 6 + CH (31%), 2C 3 H 3 + M C 6 H 6 + M (49%). The first two reactions proceed in the upper atmosphere and make a peak of the benzene mixing ratio of at 960 km. Figs reflect the calculated vertical profiles of 55 neutral species and the basic observational data for these species. It is not easy to show all these data in four figures. However, the figures are readable and clear even in the black-and-white option using the designations from Fig Nitriles The N N triplebondinthen 2 molecule is very strong with dissociation energy of 9.76 ev. Ionization of N 2 by the solar photons λ<80 nm, photoelectrons, magnetospheric electrons, protons, and cosmic rays forms mostly N + 2 ions that return N 2 after charge exchange with hydrocarbons. Predissociation of N 2 at λ = nm, dissociative ionization at λ <51 nm, and electron, proton and cosmic ray impact dissociations originate N and N atoms and N + ions with a total column rate of cm 2 s 1. Most of them (67%) form nitriles C x H y CN. The basic reaction of the nitrile formation is N + CH 3 H 2 CN + H (72% of the column production of nitriles). The C N bonds are strong (7.85 ev in CN and 9.58 ev in HCN); therefore nitriles do not recombine to N 2 in Titan s atmosphere, and condensation and polymerization with precipitation to the surface are the ultimate fate of nitriles on Titan Hydrogen cyanide HCN Almost all H 2 CN molecules quickly convert into HCN in the reactions with H and N. They give 25% of the HCN production, reactions of CN with CH 4,C 2 H 4,C 2 H 6, and ion reactions give 50 and 24%, respectively. Losses of HCN by photolysis (7%), reactions with CH (11%), C 2 H 3 (11%), C 3 N (40%), ion reactions (17%), and condensation near the tropopause (13%) balance this production. The production and loss of HCN have three maxima at 100, 400, and 950 km. The peaks at 100 and 950 km are initiated by the reaction of N + CH 3 and the ion reactions in the lower and upper ionospheres, respectively. Reactions of CN with hydrocarbons are almost completely responsible for the peak at 400 km. The calculated vertical profile of HCN (Fig. 9) is near the UVIS occultation data and the CIRS nadir observations. Two profiles derived from the CIRS limb observations and three profiles from the measurements in the millimeter range are also shown. The agreement with the profiles by Tanguy et al. (1990) and Gurwell (2004) is very good. The calculated abundance at 1100 km is still within the uncertainty of the value from Vuitton et al. (2007) Cyanoacetylene HC 3 N The model profile (Fig. 9) agrees with the HC 3 N abundance at 1100 km from Vuitton et al. (2007) but exceeds the CIRS limb and IRAM millimeter measurements. Maybe some polymerization and heterogeneous loss processes are missing in our scheme that is based on WA04 and LCV. Production of HC 3 N in the reactions of CN + C 2 H 2 (62%) and C 3 N + CH 4 (38%) is balanced by the loss in photolysis (94%), polymerization with C 4 H (0.7%), condensation (0.6%), and ion reactions (4%). The abundances of HC 3 N and HCN are sensitive to partition of CN which is produced by photolysis of HCN and either forms HC 3 N in the reaction with C 2 H 2 or returns HCN in the reactions with CH 4,C 2 H 4, and C 2 H Acetonitrile CH 3 CN This is another species that has been detected in the millimeter range, and its profile was retrieved from the observed line shape (Marten et al., 2002). As for the most of the observed stratospheric profiles, the calculated profile of CH 3 CN does not fit the observation while the mean values at km are in reasonable agreement. The species abundance at 1100 km from Vuitton et al. (2007) also agrees with the model. Acetonitrile is formed by the reactions of N + C 2 H 4 (19%), N + C 2 H 3 + M (11%), and CH 2 CN + H + M (69%) and lost in the ion reactions (45%), reaction with H (10%), photolysis (5%), and condensation near the tropopause (41%).
18 Titan s photochemistry 243 Fig. 9. Nitriles in Titan s atmosphere. See designations in Fig. 5. Dotted line is the microwave observations of HCN by Gurwell (2004) Cyanogen C 2 N 2 We have not found a rate coefficient for the reaction of CH 2 CN + CN (r271). Furthermore, enthalpy of CH 2 CN is lacking in the NIST database, and HCN + CHCN may be products of this reaction as well. This is a reaction between two radicals, and its rate coefficient may be even higher than those of the reactions of C 3 H 4 and C 4 H 2 with CN (r255 and r256). Then the reaction of CH 2 CN + CN is the major source of cyanogen C 2 N 2 and makes it possible to reach a good agreement with the C 2 N 2 abundance of 4 ppm observed by INMS (Waite et al., 2007) at 1000 km. Production of C 2 N 2 in the reactions CN + CH 2 CN (53%) and N + CHCN (47%) is balanced by photolysis (45%) and condensation in the lower atmosphere (55%) Other nitriles Some other nitriles from our model are shown in Fig. 9. Abundances of acrylonitrile C 2 H 3 CN, cyanoethane C 2 H 5 CN, cyanomethylacetylene C 4 H 3 N, cyanodiacetylene HC 5 N, and pyridine C 5 H 5 N were derived by Vuitton et al. (2007) using the ion chemistry fitting to the observed mass spectrum of ions at 1100 km. Our model is in excellent agreement with these values. Balances for each of these species may be studied using the data from Tables 3 6. C 2 H 3 CN is formed by the reactions of HCN + C 2 H 3 (49%), CH 3 + CHCN (37%), and CN + C 2 H 4 (14%) and lost by photolysis (82%), the ion reactions (12%), and condensation (6%). C 2 H 5 CN is produced in the reaction of N + C 3 H 6 and removed by photolysis. Rate coefficients of the reactions of C 2 N with hydrocarbons were measured by Zhu et al. (2003). They did not observe the reaction products but argued in favor of the hydrogen abstraction. However, formation of C 4 H 3 N in C 2 N + C 2 H 6 is not ruled out, and this is the only way in our model to fit the abundance of this species from Vuitton et al. (2007). C 4 H 3 Nislostbyphotolysis (84%) and in the ion reactions (16%). HC 5 N is balanced between production in CN + C 4 H 2 (81%) and C 2 H + HC 3 N (19%) and loss by photolysis. Reaction C 3 N + C 2 H 6 is the only source of pyridine in our model. The adopted rate coefficient is smaller by a factor of 3.5 than that adopted by Yung (1987) for the same reaction with products HC 3 N + C 2 H 5. Maybe, an excited adduct forms in this reaction and then is stabilized by release of H. Pyridine is lost by photolysis N x H y,ch 2 NH, and CH 3 NH 2 Our model involves nine species with N H bonds, and the calculated vertical profiles for five of them are shown in Fig. 10. The N H bonds are 3.5 ev and comparable to the C H bonds. Therefore N x H y may return N and recombine to N 2,forexample, NH + CH 3 CH 4 + N, NH 2 + N N 2 + H + H. That is why N x H y are much less abundant than nitriles on Titan. Vuitton et al. (2007) give densities of CH 2 NH and NH 3 at 1100 km. Our model agrees with the former and much smaller than the latter, although we adopted very high yields of NH 3 in the reactions of N with C 2 H 5 and C 3 H 6 that may be significantly overestimated Oxygen species We have discussed in Section 3.5 that the total production rate of oxygen species is smaller than those of hydrocarbons and nitriles by a few orders of magnitude. Therefore the numbers of oxygen species and their reactions are much smaller than those in WA04. Unlike some previous models, our model does not fix the CO mixing ratio at the surface but calculates it interactively. Precipitation of the interplanetary dust and meteorites results in a production of water in Titan s atmosphere with a peak near 700 km and a column rate of cm 2 s 1 (English et al., 1996). Chemistry of oxygen species on Titan begins with dissociation of H 2 O to OH and H. Wong et al. (2002) argue using the results of Pereira et al. (1997) that H 2 O is returned in the reaction OH + CH 3, and only minor branches of this reaction may further form CO and CO 2. They calculated a CO mixing ratio of
19 244 V.A. Krasnopolsky / Icarus 201 (2009) Fig. 10. Some NH, oxygen species, and radicals in the model. See designations in Fig ppm for a source of oxygen from the meteoritic water, which is much smaller than the observed value of 50 ppm. Wong et al. (2002) explained the observations by primordial CO or a surface outgassing of CO. Detection of a flux of O + ions from Saturn s magnetosphere by the Cassini Plasma Spectrometer (Hartle et al., 2006) resolvedthe problem of a source for CO (Hörstetal.,2008). The O + ions are stopped near 1100 km as the O atoms and quickly react with CH 3 to form CO and H 2 CO that dissociates to CO. The model for oxygen species by Hörst et al. (2008) involves 44 reactions of 15 species. Of 31 reactions of 12 species in our model, 21 reactions of 10 species are similar to those in Hörst et al. (2008). Weremovedall reactions with metastable O( 1 D) and O( 1 S). Our model estimates the column production of O( 1 D) in photolyses of H 2 O and CO 2 at 10 5 cm 2 s 1. 95% of this production is quenched by N 2 to the ground state O, and only 5000 cm 2 s 1 form OH and add 0.1% to its production by photolysis of H 2 O. We also remove reactions of HOCO and CH 3 COCH 3 and include CH 3 COH. The required flux of O + in our model is near the observed value and equal to cm 2 s 1 at 1100 km and cm 2 s 1 scaled to the surface. Photolysis of H 2 O forms the OH radicals that either return H 2 O in the reactions with CH 3,CH 4, and H 2 (50%), or convert CO into CO 2 (46%), or release products that form CO (4%). The CO mixing ratio would be 4 ppm without the flux of O + in our model. Almost all oxygen atoms react with CH 3 and are readily converted into CO and H 2 CO with subsequent photolysis to CO. The final product of the oxygen chemistry is CO 2 which almost completely condenses and precipitates to the surface. The CO molecule has the highest dissociation energy (11.16 ev) of all known molecules, and the CO lifetime is very long on Titan (170 million years, see below). Therefore, in spite of the weak source, CO is very abundant on Titan with mixing ratio of 47 ppm in our model. This mixing ratio is constant throughout the atmosphere. It is in excellent agreement with the CIRS observations (Flasar et al., 2005) and ground-based observations in the millimeter range that were reviewed by de Kok et al. (2007). The model data on H 2 O and CO 2 also agree with the ISO and CIRS observations, respectively. Our model favors the photochemical formation of CO on Titan without sources and sinks at the surface. This situation is very different from those on Triton and Pluto where CO originates in the atmosphere from the surface ice (Cruikshank et al., 1993; Owen et al., 1993). Vertical profiles of the most important and/or abundant radicals are shown in Fig. 10. These data may be helpful for technical calculations and evaluations Ionosphere It was discussed in Section 3.5 that our model is aimed at simulation of major ions in the ionosphere of Titan. Therefore, these ions and their reactions should be carefully chosen, while their numbers are significantly smaller than those in WA04 and Vuitton et al. (2007). Finally, the model involves 111 reactions of 33 ions. Rate coefficients of ion-neutral reactions and dissociative recombination are taken from McEwan and Anicich (2007) and Woodall et al. (2007), respectively. Most of the recombination processes have many channels. For example, recombination of C 2 H + 3 has 12 channels. We adopt only the main channel for each reaction to avoid overloading of the model by details that may be not very important. We assume C 3 N + H 2 as the products of HC 3 NH + + e. We have not found rate coefficients and products for recombination of C 5 H + 7 and C 7H + 7 and assume the values given in Table 5. Radiative association reactions between hydrocarbon ions and hydrocarbons may form large hydrocarbon ions, and six reactions of this type are taken from McEwan and Anicich (2007). Wealso include five radiative association reactions ( ) that form C 10 H + 9 and are not given in McEwan and Anicich (2007). Wehave checked enthalpies and proton affinities (Hunter and Lias, 1998) of the species in these reactions and found that these reactions are exothermic. (C 3 H 2 wasnotfoundinthenistdatabase.)their rate coefficients are adopted at cm 3 s 1 which is smaller than the values for the similar reactions in McEwan and Anicich (2007).
20 Titan s photochemistry 245 Fig. 11. Electron density profiles in the daytime (solid line) and nighttime (dashed lines) ionospheres of Titan. The nighttime profiles refer to mean and strong magnetospheric electron precipitations. The observed mean dusk and dawn profiles for solar zenith angles of 87 and 95 (Kliore et al., 2008) andthetotalinmsion density profile at 130 during the T5 encounter (Cravens et al., 2009) areshown (thin lines) Electron density profiles The calculated electron density profile of the daytime ionosphere is shown in the linear scale in Fig. 11. The daytime ionosphere consists of a main peak at 1060 km, steplike features at and km, and a lower peak at 70 km which is due to the cosmic ray ionization. The steplike feature at km is mostly caused by the solar photons with λ<20 nm. Their flux (Fig. 6) is comparable to that at nm. This feature appears because the absorption cross sections of N 2 at λ<20 nm are significantly smaller than those at nm. The feature at km is mostly formed by magnetospheric protons. The calculated profiles of the nighttime ionosphere in Fig. 11 refer to the mean magnetospheric precipitation and to the T5 conditions of strong precipitation. They reflect the adopted ionization by magnetospheric electrons and protons and by the cosmic rays. Four electron density profiles have been observed by the Cassini radio occultations (Kliore et al., 2008) at the low latitudes. Mean dusk and dawn profiles for the solar zenith angle of 87 and 95, respectively, are copied in Fig. 11. Our model is calculated for SZA = 60, and the peak altitude at 1060 km is similar to that calculated by Cravens et al. (2005) for SZA = 59. According to Cravens et al. (2005), the peak altitude increases from 59 to 90 by 140 km while the peak electron density becomes smaller by a factor of 2. With these values taken into account, we may state that the main ionospheric peak in our model is in reasonable agreement with the observed radio occultation profiles. These profiles also have the features formed by the solar photons with λ<20 nm and some structures that may be attributed to magnetospheric protons. Another published profile (Cravens et al., 2006) refers to the T5 encounter that was on April 16, 2005, with the closest approach at 1027 km at SZA = 127 and a local time of 23:15, i.e., at the nighttime conditions. The profile in Fig. 11 is an approximate sum of all ions observed by INMS. The electron density profile observed by RPWS is very similar to this sum below 1400 km. The T5 encounter was during a strong precipitation event. Our model for this event fits rather well the observed electron density profile. The fitting could be further improved by adjustment of the adopted ionization profile (Fig. 4) Effective recombination rate coefficient This coefficient is defined as a ratio of ionization rate to n 2 e. The calculated profile of this coefficient is shown in Fig. 12 Fig. 12. Effective recombination rate coefficient (solid line) calculated as a ratio of the total ionization rate to n 2 e.itiscomparedwith10 6 (300/T e ) 1/2 cm 3 s 1.The electron temperature profile based on Galand et al. (2006) is also shown. along with the electron temperature and the value of 10 6 (300/T e ) 1/2 cm 3 s 1 that is typical of heavy ions (Table 5). This value closely matches the model coefficient at 200 to 900 km. The model coefficient at 1100 to 1300 km is smaller than this value by afactorof 2.7 and corresponds to (300/T e ) 1/2 cm 3 s 1 which is typical of light ions with mass less than 30 (Table 5). The strong increase in the effective recombination coefficient above 1300 km reflects an increasing role of the ambipolar diffusion in the loss of ions. Actually the effective recombination coefficient does not make sense above 1350 km. The stronger is ionization rate, the higher is the role of primary ions that are typically light. This is the explanation of the difference between the model recombination coefficient and the value for heavy ions below 200 km, where the cosmic ray ionization becomes strong. Using the recombination coefficient of cm 3 s 1 at km from Fig. 12 and the electron density at these altitudes of 1000 cm 3 measured during the T5 encounter, the required ionization rate may be estimated as 0.25 cm 3 s 1 which is close to that adopted in our model (Fig. 4) Ion composition The calculated profiles of fourteen most abundant ions are shown in Figs. 13a, 13b, and 13c. HCNH +,C 2 H + 5,C 3H + 5, and CH+ 5 dominate above 1000 km. HC 3 NH +,C 9 H + 11,C 2H + 7, and CH 3CNH + are the most abundant ions below 800 km. The currently published data on the ion composition refer to the T5 encounter, and HCNH +,C 2 H + 5,C 3H + 5, and CH+ 5 are the most abundant in the observations (Cravens et al., 2006, 2009) and in our model (Fig. 13d). Ions in our model are subject to ambipolar diffusion and escape with velocity of D i /2H i at the upper boundary. Column loss from the ion escape is cm 2 s 1 for H, cm 2 s 1 for C, and cm 2 s 1 for N. The ion escape rate is smaller than the escape rate for exothermic chemistry (De La Haye et al., 2007) by factors of 3.5 and 13 for C and N, respectively. The ion escape reduces ion densities at the upper boundary (1600 km) by 10 30%. Effect of ambipolar diffusion is much stronger: the ion scale heights are significantly smaller above 1400 km than those calculated without ambipolar diffusion, the electron densities (which are sums of the ion densities) are smaller than those without ion diffusion by factors of 3.7 at 1600 km and 1.36 at 1400 km. The approximation of photochemical equilibrium for ions in the models by Banaszkiewicz et al. (2000) and WA04 becomes valid below 1350 km.
21 246 V.A. Krasnopolsky / Icarus 201 (2009) Fig. 13. Ion composition of the daytime ionosphere above (a, b) and below (c) 600 km. The nighttime ionosphere (d) is calculated for the encounter T5 with strong precipitation of magnetospheric electrons. Table 7 Densities of the most abundant ions in the nighttime ionosphere. The INMS data at 1100 km are compared with the ion chemistry fitting by Vuitton et al. (2007), the model by Cravens et al. (2009), and two our models at 1060 km. a Ion HCNH + HC 3 NH + C 3 H + 3 C 2 H 3 CNH + CH 2 NH + 2 C 2 H + 5 C 3 H + 5 CH 3 CNH + C 5 H + 5 INMS VYM Cr Model Model Ion C 4 H + 3 C 4 H 3 NH + C 5 H + 7 C 4 H + 5 C 7 H + 7 NH + 4 C 6 H + 7 CH + 3 CH + 5 C 6 H + 5 F INMS VYM Cr b Model Model a Ion densities are in cm 3, INMS is the measured ion mass spectrum from Vuitton et al. (2007, Fig. 2), VYM is the ion chemistry fitting in Vuitton et al. (2007), andcr09 is Cravens et al. (2009). Model (see Appendix A) is with eddy diffusion from Hörst et al. (2008). F is the mean model-to-inms ratio (see text). b C 7 H + 7 and C 6H + 5 are not included, otherwise F = The ionospheric peak at 70 km (Fig. 11), which is formed by the cosmic rays, is narrow with a width of 80 km at half maximum. The calculated ion composition of this layer is shown in Fig. 13c. We do not consider ion clusters and negative ions (Molina-Cuberos et al., 1999; Borucki et al., 2006) because the cluster bond energies and electron affinities of negative ions are very low, and the visible and near infrared solar light may effectively dissociate them. However, the ion clusters and negative ions may be abundant in the nighttime. We also do not account for ionization of haze particles. According to Borucki and Whitten (2008), ionization potentials of these particles are rather low, and the ionization rate may significantly exceed that from the cosmic rays. Ionization by the cosmic rays is a source of chemistry in the lower atmosphere of Titan which is protected from the solar UV radiation by the haze absorption. Meteorites and the interplanetary dust are the sources of metals and some other atomic species with low ionization potentials. These species may be endmembers of the charge exchange chains or ionized directly by the far UV photons ( nm). Radiative recombination of atomic ions is slower than dissociative recombination of the molecular ions by five orders of magnitude, and the atomic ions may be rather abundant even at a very low production rate. According to Molina-Cuberos et al. (2008), the meteoric ions should appear near 700 km in Titan s atmosphere. Our model does not involve these ions Comparison with the INMS measurements at 1100 km Vuitton et al. (2007) and Cravens et al. (2009) published a mean ion mass spectrum during the egress of the T5 encounter. The INMS data collected at altitudes of 1027 to 1200 km were averaged and refer to a mean altitude of 1100 km, local time of 23:15, solar zenith angle of 127, and latitude of 74 N. Densities of the most abundant ions in this spectrum are given in Table 7. Vuitton et al. (2007) fitted the spectrum by a model of the ion chemistry that involved 150 ions and 1250 reactions. Densities of many neutral species at 1100 km were used as fitting parameters, and the obtained values are important to study the chemical composition of Titan s atmosphere. They are shown in Figs and considered
22 Titan s photochemistry 247 in the previous sections. The calculated ion densities from Vuitton et al. (2007) are also given in Table 7. Cravens et al. (2009) calculated six versions of the ionospheric model for the adopted various neutral composition. Their preferable model at 1100 km is shown in Table 7 as well. The electron density profile in our model for the T5 conditions intersects the measured profile near 1060 km (Fig. 11), and we compare the ion densities from our model at this altitude with the observed densities in Table 7. We also show in this table the results of our second model from Appendix A. Some model densities fit better the observations and some worse. To estimate the overall quality of fitting, we involve a mean model-to-inms ratio that may be calculated as ( 1 n F = exp ln(ai /A i0 ) ), n i=1 where A i and A i0 are the ion densities in the models and measured by INMS, respectively. F = 1 for ideal fitting to the observed ion spectrum in this formulation, and, say, F = 2 means that a mean difference between the calculated and observed ion densities is a factor of 2. The calculated values of F forfourmodelsshow that the agreement of our models with the observed ion spectrum may be considered as very good: F = 1.74 and 2.01 for Vuitton et al. (2007) and our model, respectively. We cannot adjust abundances of neutral species in our model and therefore have much smaller degrees of freedom than Vuitton et al. (2007). The fitting in Vuitton et al. (2007) is better than that in our model for minor ions not shown in Table Discussion 5.1. Comparison with the observations We have already compared our model profiles with the observational data, and will discuss here some general features. There is a significant scatter between the observations by different instruments and techniques, and our profiles are sometimes close to the mean values. The UVIS occultation observations were analyzed assuming six major absorbers in the UVIS spectral range of 110 to 190 nm. Our model and the interpretation of the ion mass spectrum by Vuitton et al. (2007) involve many other abundant species, and some of them may contribute to the observed UV absorptions and result in significant corrections to the retrieved abundances. There are considerable differences in the slopes of the calculated profiles and those retrieved from the CIRS limb observations in the stratosphere. Formation of many species peaks near 500 km, and their sinks are below 200 km. These species form downward fluxes between 500 and 200 km. Flux Φ and mixing ratio f are related by Φ = Kn df dh, where n is the total number density. In our basic model such a species has a decrease in its mixing ratio from 500 km to the lower atmosphere (which is not supported by the CIRS limb observations) and a constant mixing ratio at 500 to 1000 km. If the species mass is near 28 (C 2 H x and HCN), then the constant mixing ratio is extended into the upper atmosphere. If the mass exceeds 28, then the mixing ratio diminishes in the upper atmosphere because of the diffusive separation. This is a qualitative explanation of the typical vertical profiles of photochemical products on Titan. Very low eddy diffusion below 75 km and a steep increase of K from 75 to 300 km by five orders of magnitude suggested by Fig. 14. Shares of the nitrile plus ion reactions in the productions of some hydrocarbons. Hörstetal.(2008)help to fit the CIRS limb observations but result in some other problems. Our model with this eddy diffusion is considered in Appendix A Coupling of hydrocarbon, nitrile, and ion chemistries The role of the nitrile and ion reactions in the production and loss of hydrocarbons looks moderate in the discussion in Section 4.1. For example, these reactions give 24% of the C 2 H 2 loss, 9% of both production and loss of C 2 H 4, 26% of the C 2 H 6 loss, and 50% of the benzene production. The nitrile and ion reactions are very important in the balances of hydrocarbons in the upper atmosphere and below 100 km (Fig. 14) and significantly affect the abundances of hydrocarbons, especially at high altitudes. Therefore, hydrocarbon, nitrile, and ion chemistries are strongly coupled on Titan, and attempts to calculated them separately (e.g., in some previous models and in all models of ionospheric composition) may result in significant error Polycyclic aromatic hydrocarbons and heavy ions in the upper atmosphere TheINMSmeasurementsnear1000km(Waite et al., 2007) revealed the C 6 H 6 mixing ratio of and its scale height close to that for diffusion. Our model is in good agreement with these results. Waite et al. (2007) argue that the abundant benzene may form PAH near 1000 km, and this is an important source of tholins in the upper atmosphere. They see a confirmation of this hypothesis in the low resolution mass spectrum of ions observed with the Cassini Plasma Spectrometer (CAPS). They conclude that the ultimate fate of benzene dissociation is the production of higher-order PAHs. However, while the mixing ratio of looks significant, the corresponding number density at 1000 km is very low, 7000 cm 3, and only fast reactions between radicals and ion reactions may be essential at these densities. Fascella et al. (2004) made quantum chemistry calculations of key reactions involved in the formation of naphthalene C 10 H 8 and indene C 9 H 8 that are the simplest bicyclic hydrocarbons and may be initial blocks for PAHs. They found that the fastest initial reaction is C 6 H 6 + C 4 H 5 C 10 H 10 + H, k = e 3220/T cm 3 s 1, for which the column rate is only 0.02 cm 2 s 1 in our model. Another reaction that may initiate PAH is C 6 H 6 + C 6 H 5 C 12 H 10 + H, k = e 2010/T cm 3 s 1
23 248 V.A. Krasnopolsky / Icarus 201 (2009) (Fahr and Stein, 1989) with a column rate of 712 cm 2 s 1 which is extremely low as well. Therefore, the neutral chemistry is not responsible for formation of PAHs in the upper atmosphere. The ion chemistry may be more promising. Recombination of C 6 H + 7 and C 7H + 7 gives 50% of the column production of C 6H 6 and almost full production in the upper atmosphere. It results in the benzene mixing ratio of at 950 km (see Section and Fig. 8), in excellent agreement with the INMS data (Waite et al., 2007). We have found two reactions in the compilation by McEwan and Anicich (2007) that form bicyclic hydrocarbon ions: C 7 H C 2H 4 C 9 H hν, C 7 H C 3H 4 C 10 H hν. Proton affinity of naphthalene is high (8.32 ev, Hunter and Lias, 1998), and we adopted reactions of formation of C 10 H + 9 and recombination of C 9 H + 11,C 10H + 9, and C 10H The sum of these ions has peaks of 380 cm 3 at 780 km and 850 cm 3 at 70 km (Figs. 13a and 13c). The CAPS ion mass spectrum gives 300 cm 3 for ions with mass exceeding 115 at 1000 km (Waite et al., 2007). C 9 H + 11, C 10H + 9, and C 10H + 11 may be initial blocks for heavier ions, and the ion densities in the model are close to those observed except for the difference in altitude (800 and 1000 km). The CAPS spectrum was observed at SZA = 90 while our model is for 60, and this adds to the difference between the model and the CAPS spectrum. This suggests that other processes could be involved as well Negative ions The initial CAPS spectrum of negative ions (Waite et al., 2007) gave the number density of 10 4 cm 3 for negative ions with mass μ > 80. Later the observations were analyzed by Coates et al. (2007); they give 100 cm 3 for these ions at 950 km and SZA = 105. Coates et al. (2007) consider dissociative electron attachment as a mechanism for formation of negative ions: AB + e A + B. Evidently this process may proceed if electron affinity of A exceeds dissociation energy of AB, and this does not work for the most of species on Titan. This process is possible for high-energy electrons; however, negative ions are very minor products and typically not observed in the interactions of high-energy electrons. Direct electron attachment to complex molecules as PAHs is more effective and has a rate coefficient of cm 3 s 1 for anthracene (Moustefaoui et al., 1998). Coates et al. (2007) consider the conditions at 950 km and SZA = 105 as relative darkness eliminating electron photodetachment processes. However, these conditions correspond to the sunlight passing through the atmosphere at the tangential altitude of 830 km. Using our model of the chemical composition, the atmosphere is transparent in this geometry for λ>150 nm and opaque for λ<100 nm. Electron affinities of PAHs are ev, and photodetachment should be very effective in the visible light. Therefore, there is no photoionization at 950 km and SZA = 105 but electron photodetachment proceeds. Photodetachment rates in the Earth s atmosphere are in the range of 10 3 to 1 s 1 (Schunk and Nagy, 2000). A mean of this range scaled to Titan s heliocentric distance is s 1. Using electron density of 1000 cm 3 and the negative ion density of 100 cm 3 for μ > 80, one can evaluate the PAH density from a balance between electron attachment and photodetachment: PAH + e PAH, PAH + hν PAH + e. Fig. 15. Formation of haze on Titan. Recombination of heavy ions and polymerization of C x H y N have secondary peaks at km. Processes of condensation and precipitation occur below 100 km and are not shown. Reaction rates are scaled by (1 + h/r 0 ) 2. Then [PAH]= = cm 3. The atmospheric density is cm 3 at 950 km, and the calculated PAH mixing ratio is 10 6 at this height. It may be compared with the benzene mixing ratio of at 950 km in the INMS observations and our model and looks reasonable. Therefore, the observed heavy negative ions require the PAH density which is about a quarter of the benzene density at 950 km. However, other reactions of negative ions are not ruled out Production of haze Sources of the haze on Titan are shown in Fig. 15. Recombination of the heavy ions has two peaks at 800 and 70 km with productions of 150 and 80 g cm 2 Byr 1, respectively. Polymer formation in the reactions peaks at 500 km and is equal to 1.8 kg cm 2 Byr 1. Reaction 207 C 6 H + C 4 H 2 gives 80% of this value. Polymerization with nitriles also gives 1.8 kg cm 2 Byr 1 with the bulk contribution of C 3 N + C 4 H 2 (95%). Its secondary peakat100kmproduces30gcm 2 Byr 1. A C:H:N ratio in the haze from polymerization and recombination of the heavy ions is 1:0.33:0.07. Condensation and precipitation of hydrocarbons and nitriles below 100 km (Table 6) adds 3.2 kg cm 2 Byr 1 to the total haze production that is equal to 7 kg cm 2 Byr 1. As in the previous models, we do not consider sublimation of species near the surface where the atmosphere is warmer than at the tropopause. Therefore it is better to consider this value as an upper limit. A C:H:N ratio in the condensate is 1:2.5:0.05. The C:H:N ratio in the total aerosol flux at the surface is 1:1.2:0.06 in the numbers of atoms. This ratio is in contrast with results of the Huygens Aerosol Collector and Pyrolyser (ACP) which gave N > C in Titan s aerosol (Israël et al., 2005). Though N 2 is much more abundant than CH 4 and other hydrocarbons on Titan, the N 2 molecule is highly inert and may be broken only by the solar EUV photons (λ <100 nm) and the energetic particles. The hydrocarbon chemistry on Titan is initiated by the solar far UV photons ( nm). A FUV/EUV photon flux ratio is 100, almost all products of the hydrocarbon chemistry except H and H 2 condense and precipitate as the haze particles, and hydrocarbons must be the major component of these particles. Clusters including N 2 and scavenging of HCN are the only ways to get haze particles with high N/C ratio.
24 Titan s photochemistry 249 Rannou et al. (2003) estimated the required haze production rate as 200 g cm 2 Byr 1 at 400 km. This rate is significantly smaller than that in our model. Rate coefficients of the reactions of polymerization are poorly known and may be adjusted to fit the required haze production rate. However, detailed modeling of the haze is beyond our objectives. WA04 did not estimate the haze production rate in their model. LCV do not consider contribution of condensation in the haze production. According to LCV, the nitrogen component in the haze production is 220 g cm 2 Byr 1, that is, 46% of the total haze production rate in LCV. The nitrogen precipitation rate from polymerization is 260 g cm 2 Byr 1 in our model, similar to that in LCV. LCV involve 73 reactions with unidentified products that remove reactants from the atmosphere and are similar in this aspect to the reactions of polymerization. However, effects of these reactions are not considered in the production of haze in LCV Lifetimes and photochemical balance of the atmosphere The escape rates of H 2 and CH 4 in our model (Table 6) support the calculation of hydrodynamic escape by Strobel (2008) and the conclusions by Yelle et al. (2008, 2006). Escape of H is 160 g cm 2 Byr 1, escape of CH x, NH x, OH, and H 2 O is 62 g cm 2 Byr 1, and ion escape is 2.4 g cm 2 Byr 1.Totalescape is 4.4 kg cm 2 Byr 1, that is, near the haze production without condensation. Here we do not consider nonthermal escape processes (De La Haye et al., 2007) and adopt the assumption of Strobel (2008) that only species with low mass (μ 18 in our model) may escape. Actually this problem needs further study. Total loss of the atmosphere is 11 kg cm 2 Byr 1 with a C:H:N ratio of 1:4.0:0.05, that is, CH 4 :N 2 = 40:1. The total losses of C and H are in the exact proportion for methane. Formation of nitriles is the only irreversible sink of N 2 in our model. Lifetime of N 2 relative to this process is 25 Byr, much longer than the age of the Solar System. Another long-living species is CO with a chemical lifetime of 0.17 Byr. If both N 2 and CO are involved in hydrodynamic escape, their lifetimes may be significantly shorter. The haze production rate of 0.2 kg cm 2 Byr 1 from Rannou et al. (2003) may be considered as a lower limit and compared to the global-mean depth of the liquid and solid organic sediments which is 3 m, that is, 0.3 kg cm 2 (Lorenz et al., 2008). This amount is delivered by the haze particles in less than 1.5 billion years. If methane is a significant part of the observed lakes and seas, then the time for recycling of the precipitating product at the surface is even smaller. Lifetimes of CH 4 and H 2 are 21 and 0.8 Myr in the atmosphere, and lifetimes of some species are shown in Table 6. Even if the precipitating hydrocarbons and nitriles are reprocessed in the solid body of Titan to return N 2 and CH 4,thetotal loss of methane by its escape and escape of H and H 2 (reduced by factors of 4 and 2 to get the loss of methane) for 4.6 Byr is equivalent to a global ocean of 0.5 km deep. Water and ammonia may be additional sources of hydrogen (and nitrogen in the case of ammonia). This problem is beyond the scope of this paper. 6. Conclusions Unlike the previous models, our coupled model of Titan s atmosphere and ionosphere involves ambipolar diffusion, ion escape, hydrodynamic escape of light species and calculates densities of H 2 and CO near the surface that were assigned in some previous models. The model includes new species observed or indirectly derived from the Cassini observations. The model results are in reasonable agreement with the Cassini data and other observations. However, the calculated profiles of species in the stratosphere are systematically steeper that those measured by the CIRS on Titan s limb. Hydrocarbon, nitrile, and ion chemistries are strongly coupled on Titan. Therefore the approach in some previous models, when at first hydrocarbons, then nitriles, and finally ions were calculated, may result in significant error. The same refers to models of ionospheric composition that neglect effects of the ion chemistry on the neutral atmosphere. The ionosphere includes layers at , km and above 900 km with a peak at 1060 km, and a narrow layer at 80 km due to cosmic ray ionization. The model does not directly involve negative, metal and cluster ions, and ionization of the haze particles. The model is extended to bicyclic hydrocarbons (up to C 12 H 10 for neutrals and C 10 H + 11 for ions) but does not involve PAH chemistry. The model estimates of heavy positive and negative ions are in reasonable agreement with the Cassini results. Formation of the haze occurs in a few layers with a total rate of 4 7 kg cm 2 Byr 1, and nitrogen is 6% of this value. Therefore, the model does not support the low C/N ratio observed by the Huygens ACP in the haze. The production looks too high and may be reduced by proper reduction of the rate coefficients of polymerization. Using the data on hydrodynamic escape of methane and H 2 in our model, the loss of methane is 8.5 kg cm 2 Byr 1 and equivalent to a methane ocean 0.5 km deep for the age of the Solar System. This may be compared with the current global-mean depth of the organic sediments of 3 m. Acknowledgments I am grateful to the reviewers for some helpful comments. Appendix A. Photochemical model with eddy diffusion from Hörst et al. (2008) LCV recognized that a low value of eddy diffusion below 100 km and a steep increase above this height may facilitate fitting to the CIRS limb observations (Vinatier et al., 2007; Teanby et al., 2007) that show low gradients of the hydrocarbon mixing ratios at km. Hörst et al. (2008) suggest eddy diffusion increasing as n 2 from K = 400 cm 2 s 1 below 70 km to cm 2 s 1 above 300 km. Their version of eddy diffusion is shown in Fig. 3. We have calculated a model for this eddy diffusion which is significantly different from our basic model. Main results of this model are shown in Figs. A1 A6 which are analogs to Figs. 7 10, 13, and 15. Advantages and disadvantages of this model and LCV relative to our basic model are summarized in Table A1. The extreme behavior Table A1 Comparison of the model with eddy diffusion from Hörst et al. (2008) and LCV with the basic model. Model K H 2 CH 4 C 2 H 2 C 2 H 4 C 2 H 6 C 3 H 4 C 3 H 6 C 3 H 8 C 4 H 2 C 6 H 6 HCN CH 3 CN HC 3 N M LCV Model C 2 H 3 CN C 2 H 5 CN C 5 H 5 N C 2 N 2 C 4 H 3 N HC 5 N NH 3 H 2 O CO 2 CO Ions Haze M LCV Plus, zero, and minus mean better, similar, and worse, respectively, agreement with the observational data than that for the basic model. Nothing is shown if data are not available. The comparison is made separately for the lower and upper atmosphere if the observational data are available. M is the model with eddy diffusion from Hörst et al. (2008).
25 250 V.A. Krasnopolsky / Icarus 201 (2009) Fig. A1. CH 4,H 2,andthemajorC 2 H x hydrocarbons in the model with eddy diffusion from Hörst et al. (2008). Fig. A2. C x > 2 H y hydrocarbons in the model with eddy diffusion from Hörst et al. (2008). of eddy diffusion is difficult to explain in terms of the atmospheric dynamics; therefore it is a disadvantage of the model. The profiles of C 2 H 2,C 2 H 4, and C 2 H 6 fit much better the CIRS limb observations, and three pluses are put in Table A1. However, the ethane profile disagrees with the INMS observations in the upper atmosphere, and this is a minus. Fittings to the observations of the C 3 H 4 profiles in both models are of comparable quality, while fitting of C 3 H 6 is worse and that of C 3 H 8 is much worse in the current model. C 4 H 2 and especially C 6 H 6 disagree with the CIRS nadir observations in this model. Fittings of some nitriles to the observations are also worse in the current model. To fit the observed CO mixing ratio of ppm,
26 Titan s photochemistry 251 Fig. A3. Nitriles in the model with eddy diffusion from Hörst et al. (2008). Fig. A4. Some NH, oxygen species, and radicals in the model with eddy diffusion from Hörst et al. (2008). the required flux of O + is equal to cm 2 s 1 at 1100 km and cm 2 s 1 scaled to the surface. Electron density profiles in this model are similar to those in the basic model. Mean difference between the observed INMS spectrum of ions in the nighttime ionosphere at 1100 km (Vuitton et al., 2007) and the calculated ion densities (Table 7) is a factor of 2.10 for this model, similar to 2.01 for the basic model. (This factor is 1.74 for the model by Vuitton et al. (2007).) Production of the haze by polymerization and recombination of heavy ions is smaller in this model and LCV than that in the basic model by a factor of 2.7 and agrees better with the estimate by Rannou et al. (2003).
27 252 V.A. Krasnopolsky / Icarus 201 (2009) Fig. A5. Ion composition of the daytime ionosphere (a, b, c) in the model with eddy diffusion from Hörst et al. (2008). The nighttime ionosphere for the encounter T5 during strong precipitation of magnetospheric electrons is shown in (d). Fig. A6. Formation of the haze by recombination of heavy ions and polymerization of hydrocarbons and nitriles. Overall, this model has some advantages and shortcomings relative to the basic model and may be of some interest. References Adachi, H., Basco, N., James, D.G.L., A quantitative study of alkyl radical reactions by kinetic spectroscopy. III. Absorption spectrum and rate constants of mutual interaction for the ethyl radical. Int. J. Chem. Kinet. 11, Adachi, H., Basco, N., James, D.G.L., The acetyl radicals CH 3 CO and CD 3 CO studied by flash photolysis and kinetic spectroscopy. Int. J. Chem. Kinet. 13, Adam, L., Hack, W., Zhu, H., Qu, Z.W., Schinke, R., Experimental and theoretical investigation of the reaction NH(X 3 ) + H( 2 S) N( 4 S) + H 2 (X 1 + g ).J.Chem. Phys. 122, Allen, C.W., Astrophysical Quantities. The Athlone Press, London. Ashfold, M.N.R., Fullstone, M.A., Hancock, G., Ketley, G.W., Singlet methylene kinetics: Direct measurements of removal rates of ã 1 A 2 and b 1 B 1 CH 2 and CD 2. Chem. Phys. 55, Atreya, S.K., Atmospheres and Ionospheres of the Outer Planets and Their Satellites. Springer-Verlag, New York. Au, J.W., Cooper, G., Burton, G.R., Olney, T.N., Brion, C.E., The valence shell photoabsorption of the linear alkanes, C n H 2n+2 (n = 1 8), absolute oscillator strengths (7 220 ev). Chem. Phys. 173, Banaszkiewicz, M., Lara, L.M., Rodrigo, R., Lopez-Moreno, J.J., Molina-Cuberos, G.J., A coupled model of Titan s atmosphere and ionosphere. Icarus 147, Bartels, M., Edelbuttel-Einhaus, J., Hoyermann, K., The detection of CH 3 CO, C 2 H 5, and CH 3 CHO by rempi/mass spectrometry and the application to the study of the reactions H + CH 3 CO and O + CH 3 CO. Symp. Int. Combust. Proc. 23, Baulch, D.L., and 10 colleagues, Evaluated kinetic data for combustion modeling. J. Phys. Chem. Ref. Data 21, Baulch, D.L., and 11 colleagues, Evaluated kinetic data for combustion modelling. J. Phys. Chem. Ref. Data 23 (Suppl. I), Becker, K.H., Engelhardt, B., Wiesen, P., Rate constants for CH(X 2 π) reactions at low total pressures. Chem. Phys. Lett. 154, Benilan, Y., Bruston, P., Raulin, F., Cossart-Magos, C., Guillemin, J.C., Mid-UV spectroscopy of propynenitrile at low temperature: Consequences on expected results from observations of Titan s atmosphere. J. Geophys. Res. 99, Benilan, Y., Bruston, P., Raulin, F., Courtin, R., Guillemin, J.C., Absolute absorption coefficient of C 6 H 2 in the mid-uv range at low temperature implications for the interpretation of Titan atmospheric spectra. Planet. Space Sci. 43, Benilan, Y., Andrieux, D., Khlifi, M., Bruston, P., Raulin, F., Guillemin, J.C., Cossart- Magos, C., Temperature dependence of HC 3 N, C 6 H 2, and C 4 N 2 mid-uv absorption coefficients: Application to the interpretation of Titan s atmospheric spectra. Astrophys. Space Sci. 236, Berteloite, C., le Picard, S.D., Birza, P., Gazeau, M.C., Canosa, A., Benilan, Y., Sims, I.R., Low temperature ( K) kinetics study of the reactions of the CH 4 radical with various hydrocarbons observed in Titan s atmosphere. Icarus 194, Bohland, T., Temps, F., Wagner, H.G., The contribution of intersystem crossing andreactionintheremovalofch 2 (ã 1 A 1 ) by hydrocarbons studied with the LMR. Ber. Bunsen-Ges. Phys. Chem. 89, Borucki, W.J., Whitten, R.C., Influence of high abundances of aerosols on the electrical conductivity of the Titan atmosphere. Planet. Space Sci. 56,
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