PHOTOCHEMISTRY AND RADIATIVE TRANSFER STUDIES IN THE ATMOSPHERES OF JUPITER AND SATURN

Transcription

1 PHOTOCHEMISTRY AND RADIATIVE TRANSFER STUDIES IN THE ATMOSPHERES OF JUPITER AND SATURN CHRISTOPHER D. PARKINSON A thesis submitted to the Faculty of Graduate Studies in partial fulfillment of the requirements for the degree of Doctor of Philosophy Graduate Programme in Earth and Space Science York University North York, Ontario Summer 2002

2 Abstract This work is mainlhy an amalgam of work done over several years and different topics. There are three main areas of investigation: (a) Saturnian He 584 Å airglow intensity, (b) deuterium chemistry and emission in the Jovian Thermosphere, and (c) Jovian tropospheric deuterated species abundances, viz., CH 3 D (Parkinson et al., 1998; Parkinson et al., 2002a; and Parkinson et al., 2002a). Calculations of the Saturnian He 584 Å airglow intensity, using radiative transfer models with partial frequency redistribution and inhomogeneous atmospheric models, are presented. For reference conditions and an atmosphere consistent with the Voyager UVS occultation results, we require the eddy diffusion coefficients at the homopause, K h, to be greater than 10 9 cm 2 s 1 in order to fit the Ultraviolet Spectrometer measurements of Voyager 1 and 2 He 584 Å airglow measurements. These values of K h seem unreasonably high when compared to the earlier work of Sandel et al. (1982) and Atreya (1982). This suggests that either the values of one or more of the parameters of our model are not correct or that the measured UVS airglow is too bright and that there is a problem with calibration. Even so, we suggest that K h is likely to be greater than 10 8 cm 2 s 1 during the period of the Voyager encounters. Jupiter s atmosphere contains proto-solar abundances of H and D, and therefore may be studied in an attempt to derive the solar system D/H value. The solar system D/H problem can be approached in a number of ways: a) D and H Lyman α, Lyman-β, b) CH 3 D and CH 4, and c) HD and H 2 (from ISO/Galileo probe). First suggested by Ben Jaffel (private communication, 1999), this strategy is termed the i

3 global approach. Using this approach and utilising the same atmospheric model should ideally provide the same D/H ratio regardless of the technique used or the location in the atmosphere one is studying. The D/H problem has been extensively modelled utilising the global approach and the results of the analysis are reported here. Moreover, as a result of trying to determine the solar system D/H ratio, we are be better able to describe the structure, chemistry and dynamics of the Jovian atmosphere in a self-consistent manner. A detailed study is presented of the distribution of key deuterated species (viz., atomic D and HD) and the associated deuterium Lyman α airglow in the Jovian thermosphere. The reactions that appear to govern the abundances of these deuterated species are used in conjunction with C 2 -chemistry in a 1-D photochemical-diffusion model. While the D abundance is mainly sensitive to H densities and the vibrational temperature profile, the D vertical distribution also depends on other parameters such as eddy mixing and the uncertainty in the some of the reaction rates used. Different scenarios are considered by varying several parameters controlling the D distribution in the thermosphere. A radiative transfer model with coupling of the H and D Lyman α lines is employed to obtain line profiles and total intensities at disc centre for these scenarios. This allows a comparison of the impact of various parameters on the Jovian D Lyman α emission. An offshoot of these chemical processes in the Jovian thermosphere is the formation of CH 2 D, CH 3 D, and C 2 H 5 D, and other deuterated species. The source of these deuterated hydrocarbons and their abundance are also discussed. CH 3 D is an isotopic tracer in the deep Jovian atmosphere and susceptible to ii

4 transport and chemical effects. It is expected that the tropospheric ([D]/[H]) CH4 ratio calculated from data collected from various observations should be relatively invariable, yet previous determinations of this quantity in Jupiter have given results that may be inconsistent with experimental error bars. This suggests that there may be a problem with the interpretion of some of the observations, or that the apparent CH 3 D column abundance is variable. Reported here are the effects of varying important parameters over this pressure regime on a standard reference atmosphere, the discussion of how this impacts the CH 3 D mixing ratio, CH 3 D fractionation, the ([D]/[H]) CH4 and D/H (= ([D]/[H]) H2 ) ratios, and a comparison with the various CH 3 D and HD observations. We proceed by assuming that the in situ measurement of the D/H ratio obtained from the Galileo Probe Mass Spectometer (GPMS) is correct which is then used as a boundary condition in the deep atmosphere. The D/H value derived via the CH 3 D modelling from the deep atmosphere to the millibar pressure region reflects the GPMS D/H value initially assumed and makes our model self-consistent in the following sense: both the CH 3 D and HD reservoirs lead to the same D/H ratio when the CH 3 D chemistry and the atmospheric structure are described in the way presented in this study. Moreover, using this technique allows for the first time a way to explain the possible discrepancies in the ([D]/[H]) CH4 ratio observations since CH 3 D was first detected on Jupiter nearly 30 years ago by offering a plausible link between the CH 3 D observations and upper tropospheric dynamical processes. iii

5 Contents 1 Introduction The Planets Earth-Based and Planetary Missions Ground Based and Satellite Observations Planetary Missions Thesis Vertical Mixing and Photochemical Modelling of Atmospheres General Method of Solution Atmospheric Parameters Eddy Diffusion Coefficient, K(z) Temperature, T Boundary Conditions Photochemistry Solar Flux Photolysis and Reaction Rates iv

6 2.3.3 J-values Enhancement of J-values by inclusion of Rayleigh scattering Hydrogen Distributions Hydrocarbons Jovian Thermospheric Deuterium and Deuterated Hydrocarbons Jovian Tropospheric CH 3 D Model Validation Radiative Transfer Introduction Method of Solution: Radiative Transfer Feautrier Technique H and D Lyman α He 584 Å He 584 Å Resonance Radiation in the Outer Planets Introduction Jupiter Saturn Uranus Neptune Deuterium Chemistry and Emission in the Jovian Thermosphere Introduction v

7 5.2 Description of Models General Photochemical Model Chemistry Results and Discussion Conclusions Deuterium Abundance from HD and CH 3 D Reservoirs in the Atmosphere of Jupiter Introduction Results and Discussion Tropospheric CH 3 D Conclusions 159 A Two Stream Approximation for Calculation of J-values 164 B Jovian Hot Region Lower Bound Justification 168 vi

8 List of Figures 2.1 Nicolet diagram shows the basic interdependency between these primary hydrocarbon species (reproduced from Atreya, 1986) Eddy diffusion profiles used to calculate isotopic enrichment values and CH 3 D mixing ratios The standard temperature profile adopted for Jupiter is based on a smoothed version of the results of Sieff et al. (1997) and Lécluse et al. (1996) Temperature profiles adopted for the mesosphere and thermosphere of Jupiter and thermosphere of Saturn. (a) The Romani et al. (1996) temperature curve is based upon Voyager IRIS observations of Jupiter s southern midlatitudes (profile A). The other temperature profile is taken from Sieff et al. (1997) (profile B). See text for details (b) various vibrational temperature profiles corresponding to T v = nt, where n = 1 to 4. See text for details (c) same as (a) except bulge region shown included (d) standard reference thermospheric Saturnian temperature profile.. 60 vii

9 2.5 Enhanced J-values for the test case: H flux = cm 2 s 1 ; D flux = cm 2 s 1 ; K h = cm 2 s 1 ; K(z) = K h *(n(z 0 )/n(z)) Rayleigh cross section (cm 2 ) as a function of wavelength (Å) C 2 H 2, C 2 H 4, and C 2 H 6 mixing ratios vs. pressure where H flux = cm 2 s 1 ; D flux = cm 2 s 1 ; K h = cm 2 s 1 ; K(z) = K h *(n(z 0 )/n(z)) (a) Gladstone et al. (1996) rates with Gladstone T profile; and (b) Romani et al. (1993) rates and Romani (1996) case c reaction with Romani temperature profile Single scattering albedo vs. wavelength Optical depth vs. wavelength Enhanced J-values for diurnally averaged, optically thin case: H flux = cm 2 s 1 ; D flux = cm 2 s 1 ; K h = cm 2 s 1 ; K(z) = K h *(n(z 0 )/n(z)) Enhanced J-values for C 2 H 2 : H flux = cm 2 s 1 ; D flux = cm 2 s 1 ; K h = cm 2 s 1 ; K(z) = K h *(n(z 0 )/n(z)) Enhanced J-values for all species using Rayleigh scattering and two stream approximation: H flux = cm 2 s 1 ; D flux = cm 2 s 1 ; K h = cm 2 s 1 ; K(z) = K h *(n(z 0 )/n(z)) Enhanced J-values comparison case: H flux = cm 2 s 1 ; D flux = cm 2 s 1 ; K h = cm 2 s 1 ; K(z) = K h *(n(z 0 )/n(z)) viii

10 2.14 C 2 H 2, C 2 H 4, and C 2 H 6 mixing ratios vs. pressure using enhanced J- value calculation with Λ=1 where H flux = cm 2 s 1 ; D flux = cm 2 s 1 ; K h = cm 2 s 1 ; K(z) = K h *(n(z 0 )/n(z)) Romani et al. (1993) rates and Romani (1996) case c reaction with Romani temperature profile Mixing ratio vs. pressure; molecular diffusion not included Mixing ratio vs. pressure; molecular diffusion included The model atmospheres of some of the more relevant species considered, viz., H 2, CH 4, CH 3, CH 2 D, CH 3 D,, HD, H and D. In panel (a) A neutral temperature profile corresponding to Seiff et al. (1997) (viz., T v = T) was used, while for panel (b) the standard reference temperature profile with T v = 3T was used. Panel (c) is a close-up of panel (b) for the altitude range 400 to 700 km., corresponding to the scattering region. Panel (d) is the same as (c) except it includes a modified H + CH 3 CH 4 reaction rate suggested by Lee et al. (2000), Moses et al. (2000) as described in text. Panel (e) similar to (c) except comparing D number density profiles for T v = 3T for Seiff et al. (1997) and Romani et al. (1996) neutral temperature profiles ix

11 5.2 Various profiles resulting from calculations utilising vibrational temperature profiles corresponding to T v = nt, where n = 1, 2, 2.5, 3 and 4. (a) D profiles, and (b) CH 3 profiles, and (c) HD profiles, and (d) CH 3 D profiles (e) C 2 H 5 D profiles. Since H is constrained, the H and CH 4 profiles are not sensitive to T v. Hence, only one profile for these species is plotted in Figure 5.2c and d Density profiles for various K h and T v = 3T (a) D (A), HD (B), and CH 4 (C), and (b) D (A) and CH 3 (B), and (c) D (A) and CH 3 D (B), and (d) same as (a) except T v = 4T. The labels 1, 2, 3, and 4 refer to K h =5 10 5, 10 6, , and cm 2 s 1, respectively Column H and D as a function K h ; for the standard reference temperature profile and several values of T v : (a) H column density versus Eddy diffusion coefficient K, and (b) D column density versus Eddy diffusion coefficient K. All cases are for columns above τ CH4 = x

12 5.5 Calculated D (A) and H (B) profiles corresponding to varying the H column density by changing input H flux, Φ H at the model atmosphere upper boundary. The labels 1, 2, 3, 4, and 5 refer to Φ H =1 10 9, , , , and cm 2 s 1, respectively Calculated D (A), HD (B), and CH 3 D (C) profiles corresponding to changes in HD mixing ratio at the model atmosphere lower boundary, f HD. The labels 1, 2, and 3 refer to f HD = , , and , respectively Chemical and Diffusion time constants for D and HD, τ D, τ HD, τ diff (a)h and D Lyman α intensity profiles for several solar zenith angles with the same viewing angle (i.e. SZA = viewing angle) for the standard reference atmosphere, (b) same as (a) except close up of D Lyman α intensity profiles at 0.33Å offset from H at Å, (c) close up of D Lyman α intensity profiles at 0.33Å offset from H at Å for subsolar case for various T v = nt D Lyman α subsolar intensities as a function of vibrational temperature Chemical and diffusion time constants for the standard temperature profile. The diffusion time constant curves are for the eddy diffusion profiles given in xi

13 6.2 CH 3 D mixing ratio profiles as a function of pressure for (a) various eddy diffusion profiles considered, and (b) upper and lower bounds for reaction rate k r for standard reference case Isotopic enrichment factors as a function of pressure for various eddy diffusion profiles considered where (a) focuses on the upper atmosphere and (b) focuses on the lower atmosphere Calculated standard reference and observed D/H ratios. Measurements included with error bars: A, Beer and Taylor, (1973); B, Gautier and Owen, (1986); C, Carlson et al., (1993); D, Encrenaz et al., (1996); E, Lellouch et al., (1997); F, Mahaffy et al., (1998). Also shown is the protosolar D/H ratio, G (Geiss and Reeves, 1993) These results have been normalised to our standard reference values for f(z) and the CH 4 mixing ratio (cf. Table 3.4) Calculated isotopic enrichment factor, f(z), as a function of the eddy diffusion coefficient, K(z) xii

14 ACKNOWLEDGEMENTS It is a rare treat to be able to have the opportunity to acknowledge for the public record those who help you on a long and difficult journey. Many are there at the beginning to wish you well along your way, but few stay the course and help you see a thing through to the end, especially when the going gets tough. Thank you to all my special friends for remaining my friends during a time when I have been very, very busy and I didn t always have the time to reciprocate properly! I trust that I will improve on that count now that I am done! I am writing this before my defense and so have cautious optimism regarding my survival of that process! Smiling, I would like to thank all the members of my examining committee who put their signatures at the front of this tome: Drs. Diane Michelangeli and Vincent Tao from York University s Department of Earth and Atmospheric Science I have only just recently met. Thank you both for taking interest in my work and special thanks to Diane for making herself available during her holiday in France: John Hancock s signature could not have been more valuable to me, I am sure! Also, I would like to thank my external examiner, Dr. Sushil Atreya, not only for lending his considerable expertise in reviewing my work, but also for taking time out of a VERY busy schedule to be here. Some of you I know well, viz., Drs. Keith Aldridge, John Caldwell, and Ralph Nicholls. Keith and Ralph, I have known you both since my early days at York. Thank you for your advice, friendship and interest in my progress over the years. John, I can say the same about you although we have known each other for less time, but I especially want to thank you for your unwavering xiii

15 belief in the quality of the work I have been doing and your faith in my ability to do it. Thank you also for visiting all my talks at the DPS without fail. Finally, I would like to thank Dr. Jack McConnell, my Ph.D. supervisor, whose patience, mentoring, and forbearance saved my academic career from turning into a disaster. It is a debt I can t repay, but thank you for all the opportunities and really making me work for it. I am a much better scientist today because of that. Some other colleagues deserve special mention and thanks. My time in Paris with Dr. Lotfi Ben Jaffel of the Institut d Astrophysique de Paris was very productive and enlightening. Thank you for all your help, fruitful discussions and mentoring I was there on my Bourse Postdoctorale Chateaubriand. My friend, Dr. Reneé Prangé I thank for your hospitality, many synergistic discussions, and opportunities extended. I always have appreciated your enthusiasm for my work. I look forward to working together more with both of you in the future. I really love working in the field: I have had enthusiam for the planets and their atmospheres since I was four when I first realised that the solar system is my backyard and cool things are going on there. For my daughter and son, Emily and Luke, I hope that you can both realise your aspirations and dreams, whatever they are and remember that passion and love are the only real reasons for doing anything. I love you lots, you two. This feeling is well reflected for this work in the following quotation: Enough for me the mystery of the eternity of life, and the inkling of the marvellous structure of reality, together with the single-hearted xiv

16 endeavour to comprehend a portion, be it never so tiny, of the reason that manifests itself in nature. Albert Einstein, [56] There are other passions and loves too without which humanity would have utterly failed by now. I am so glad I have lived to find this out: thank you for making me realise that, Dearest Heart...I wrote this poem to say so...xoxoxo Eudaemonism To hear your voice, sweet music to my soul Your kiss on my brow Like a gentle breeze on heather. Your lips on my body our arms around each other as tightly as our souls are entwined. Knowing your touch, making my breath and pulse surge for all the love in my heart for you. Your choosing to be mine, accepting me as your very own, and my giving myself to you. xv

17 To look in your eyes and tell you I love you, feeling the universe shifting all around our nerve endings exposed to the stars. Reaching out and touching you, companion of my heart, as close as hope s desire. Loving you more than anything feeling the bonds weaving themselves together, pulling us tight, heart to heart. Trusting you, feeling secure, being kind to each other in ways that only we could ever know. Knowing with absolute certainty that you will be there, And realising I have never been happier. nuff said...i am out of here! xvi

18 Chapter 1 Introduction Throughout recorded history humankind has pondered on the origin of the universe, tried to understand how it evolved into its present state, and speculated on what will finally become if it. While the questions Where did the universe come from? How and why did it begin? Will it come to an end, and if so, how? (Steven Hawking, 1988) certainly are of interest to us all, understanding what the universe is now and how it got that way is of equal interest. In a little corner of the Milky Way galaxy, our solar system slowly evolved into its present state after the universe came into being. Most of the matter in the solar system formed our Sun, the remainder going into the objects that are within its gravitational influence: the terrestrial and giant gas planets, Pluto, planetary moons, the asteroid belt, the comets, the Kuiper belt and Oort cloud. The subject of this thesis is the study of interactions between the important atmospheric constituents of two of the four gas giants orbiting the Sun, viz., Jupiter and Saturn, and the modelling of the scattering of solar radiation of some species of interest. 1

19 1.1 The Planets The terrestrial planets (Mercury, Venus, Earth and Mars) are characterised by the planet being solid and having atmospheres that are relatively thin (tens or hundreds of km) or virtually no atmosphere at all. Their atmospheric compositions vary widely and they have a composition that is not representative of solar values. The gas giants (Jupiter, Saturn, Uranus, and Neptune) are much greater in size and have atmospheres that are much thicker (tens of thousands of kilometres) than their terrestrial counterparts. All but one of the planets in our solar system have an atmosphere and moreover, some moons of these planets also have atmospheres. Jupiter and Saturn have no solid surface whereas Uranus and Neptune are thought to have ice and rocky cores. Planetary atmospheres may be divided into different regions by temperature structure: these regions have somewhat different dynamical and compositional properties. On the Earth, the lowest layer is called the troposphere which goes from the surface up to a mean height of about 15 km. In this region, the temperature decreases with height up to the tropopause, located at around 15 km. Vertical mixing in this region occurs on the order of weeks, with the exception of thunderstorms and other severe weather phenomena. The stratosphere lies above the troposphere going from about 15 km to 50 km and is a region of slower vertical transport. Here the temperature increases with height up to the stratopause owing to the absorption of solar UV radiation by ozone. Above the stratopause is the mesosphere which is characterised by a decrease in temperature from 50 to 85 km and more rapid vertical transport 2

20 (months rather than years). The region separating the mesosphere from the thermosphere above is called the mesopause. The thermosphere is where the temperature again increases with height up to about 300 km. Analogues to these regions occur in the majority of the planets, with the exception that often the stratosphere and mesosphere merge into one region as there is no absorber of solar UV, equivalent to ozone, in the stratosphere. Another way to characterise the atmosphere is in terms of its vertical mixing properties. Two distinct regions obtain and the boundary between them is called the homopause. Fluid motions on planetary scales right down to metres (or even smaller) cause the lower layer to be dominated by mechanical mixing which can often be approximated by use of an eddy diffusion coefficient, K. This results in long lived species being well mixed in this part of the atmosphere. The distribution of species in the uppermost layer is dominated by molecular diffusion. The characteristic feature of this upper region is the tendency towards separation of species into layers with each species having its own scale height: heavier species are lower with lighter species higher up. The density of the whole atmosphere decreases with height. The homopause is loosely defined as the location where K = D. At sufficiently high altitudes where the density is smaller than 10 8 cm 3, the molecules are in freefall trajectories where collisions are very infrequent. This region is called the exosphere and molecules can escape and re-enter the atmosphere. Upward-moving particles can escape from the atmosphere if their kinetic energy exceeds the planetary gravitational potential. The vertical variation of pressure in a planetary atmosphere is a feature which 3

21 reflects the temperature structure. Above 1 bar, the variation of pressure, p, with height, z, is given by p(z) = p(z o ) e ( z zo dz/h) where p(z o ) is the pressure at some reference level, z o (e.g. Chamberlain and Hunten, 1987). H = kt/mg is called the (pressure) scale height, where k is Boltzmann s constant, m is the mean mass of an atmospheric molecule (kg) and g is the acceleration due to gravity. At the Earth s surface, H is about 8 km while on Jupiter H 25 km at T = 150K. Thus there is an exponential decrease of atmospheric pressure with height (and using the perfect gas law, p = nkt, also an exponential variation of number density, n). Other very basic, but important, aspects of a planetary atmosphere are the composition and the way by which the temperature structure is maintained by energy input from the Sun, and for the outer planets, by heating from the planet s interior. As planets are spheroidal, the solar radiation heating the planet is spatially inhomogeneous with most heat being deposited (on an annual basis) at the sub-solar region (called the tropics on the Earth; Northern and Southern Equatorial Belts on Jupiter). This differential heating necessarily generates temperature gradients and pressure gradients. The pressure gradient force induces motions which then act to redistribute the energy to cooler regions. How this occurs in detail is different for each planet: its composition, rotation rate, and the degree of internal heating, etc. In the case of Jupiter, the internal heating is roughly on the order of the solar radiation component. The gas giants are similar in that their atmospheres consist largely of molecular hydrogen and helium with trace amounts of methane. However, their deep interiors may be quite different which is indicative of conditions when they condensed out of 4

22 the protosolar nebula. For instance, Uranus and Neptune both have an enhanced D/H ratio, which is thought to be due to their deuterated icy cores, and this ratio is greater than that of either Jupiter or Saturn. They also have a larger aliquot of methane (10 times solar) than Jupiter or Saturn: this larger amount of methane, which absorbs red light, results in the characteristic blue colour observed in the Uranian and Posedian atmospheres. Jupiter is 5.2 A.U. (1 A.U. is the astronomical unit distance of the Earth to the Sun) from the Sun and the largest of the planets: it is much larger than the Earth being 1320 times greater by volume and 320 times greater by mass. Its composition is mainly H 2 with a near solar abundance of helium of 13% by volume. Jupiter exhibits light and dark bands that have been named zones and belts, respectively. The belts and zones circle the planet at a roughly constant latitude: the belts have a disturbed, chaotic appearance whereas the zones are more uniform. Jupiter s dominant large-scale weather patterns (having dimensions of 10,000 km) are zonal jets and long-lived ovals. The jets have been flowing east and west and have constant speeds of up to 180 m s 1 for at least 100 years (Ingersoll et al., 1981; Limaye, 1986; Vasavada et al., 1998). The zonal jets are located on the boundaries between the belts and zones, with the westward jets on the poleward edges of the belts and the eastward jets on the poleward edge of the zones (Ingersoll et al., 2000). These jets appear to get their energy from small-scale eddies, which impart eastward momentum into the eastward jets and westward momentum into the westward jets (Ingersoll et al., 1981). The large ovals roll between the jets in an anticyclonic direction (clockwise in the northern hemisphere and counterclockwise in the southern hemisphere), 5

23 very often assimilating smaller anticyclonic eddies (Mac Low and Ingersoll, 1986). Ingersoll et al. (2000) suggest that the eddies, which ultimately drive the jets and ovals, obtain their energy from moist convection. Moist convection similar to large clusters of thunderstorms cells on the Earth is a dominant factor in converting heat flow (from a poorly understood internal heat source or sunlight) into kinetic energy (Gierasch et al., 2000). This is consistent with observations of lightning on Jupiter and would explain the anticyclonic rotation and drifting toward the poles of the eddies: moreover, this suggests patterns of upwelling and downwelling remeniscent of large-scale axisymmetric circulation in the Earth s atmosphere. Jovian cloud tops are at approximately 500 mb consisting mainly of ammonia ice crystals and haze. The main cloud formations are located at higher pressures, down to about 6 bar, with variable compostion higher up, but between 3 and 6 bar water is the only predicted condensate (Weiderschilling and Lewis, 1973; Atreya and Romani, 1985; Edgington et al., 1998). The CH 4 mixing ratio is on the order of 10 3 in a reducing atmosphere. Photodissociation of CH 4 higher in the atmosphere leads to the formation of higher order hydrocarbons such as acetylene, ethylene, ethane and aerosols which can act as tracers in the statosphere from about 100 mb to 0.01 mb. In the thermosphere (located above 1 µbar) temperatures of 1000 K or more are observed compared to what simple calculations might suggest, viz., K. This phenomena is not well understood, but high thermospheric temperatures seem characteristic of the giant planets. The strong Jovian magnetic field is able to restrain the solar wind and supports a giant magnetosphere containing high energy particles that precipitate into 6

24 polar regions causing aurorae which have been the subject of much study in recent years. Jupiter s atmosphere, is also bombarded by exogenic material (viz., Io torus, Shoemaker-Levy cometary material, rings) which can affect regions at or above the tropopause by altering ionospheric structure and the chemical composition. At 9.5 A.U., Saturn has similar composition and dynamics to that of Jupiter. It also has a very pronounced ring system, which is a possible exogenic source of H 2 O and OH, which could affect the ionospheric structure. Like Jupiter, Saturn has a pronounced internal heating source. Saturn s thermosphere is at least 420K, with a possible range of 400 to 800K. It is much hotter than anticipated and has aurorae, although not as pronounced as Jupiter s. Ionospheric densities are much smaller than expected on Jupiter and Saturn, and effects due to dynamics and vibrationally excited H 2 could account for its low electron densities. In addition, on Saturn there is also evidence to suggest that exogenic material (viz., H 2 O or OH) from the rings and satellites may also be important since, via charge exchange reactions, they transform the slowly recombining protons into rapidly combining molecular ions. Like the Earth, Saturn s obliquity results in seasonal variations however, the time scale is much longer since a Saturnian year is approximately 29 Earth years! One of Saturn s moons, Titan, has an interesting atmosphere worth noting. Voyager radio and UVS occultation observations show Titan to have an atmosphere of about 1.5 bar of N 2, some H 2 O and CO 2, methane, and a surface temperature of about 90K. In addition, solar UV radiation can break the methane down as previously described to create more complex hydrocarbons: these might condense onto the 7

25 surface due to the very cold atmospheric temperatures. Uranus is at 19.2 A.U. and is composed mainly of molecular H 2 with a relatively high abundance of methane, reflecting formation processes in the more distant and colder regions of the protosolar nebula. The upper troposphere temperature is about 50K, so cold that CH 4 can condense and form clouds. Uranus obliquity is about 98 degrees causing the rotation axis to be almost in the equatorial plane, resulting in the polar regions receiving direct solar radiation for long periods of time. Observations by Voyager revealed a very stagnant middle atmosphere with relatively small vertical mixing in the thermosphere compared with the other gas giants. As for all the other outer planets, Uranus thermosphere is much hotter than anticipated. Temperatures near the exosphere are order of 700 K (compared to less than 100K estimated from calculations). Information about the ionosphere is sparse since so few ionospheric profiles have been measured. It is known that the magnetic field is about 2/3 the strength of the Earth s dipole and is both off-set from the planetary centre by about 0.3 of a planetary radius and is tilted about 59 degrees from the planet s rotation axis. Neptune, 30.1 A.U. from the Sun, is quite similar to Uranus in size and rotation rate and composition. With more methane than Uranus and a very cold upper troposphere, the methane can freeze and form clouds. An atmospheric feature of particular note that has disappeared since it was first discovered by Voyager, was the dark oval similar to Jupiter s great red spot. Vertical mixing in the Posedian thermosphere is more pronounced than on Uranus. An analysis of its chemical composition above the tropopause, as revealed by Voyager and Earth-bound observations, indicates its 8

26 atmosphere is much more turbulent than that of Uranus and is much more like the Earth s troposphere. It does get sufficiently cold in its stratosphere that some of the methane breakdown products, such as ethane and acetylene will freeze out and form ice crystals. Similar to Uranus, information on the ionosphere is limited, but it is known that the thermosphere exhibits an unexpectedly high temperature of order of 800 K. Neptune s obliquity is similar to Earth, so seasonal variations may be expected. However, a Posedian year is approximately 165 Earth years: hence any seasonal variation may be barely perceptible even if one were to observe over several years. Enigmatic Pluto, on average the furthest planet from the sun with a maximum distance of nearly 40 A.U., has recently just returned to outside the orbit of Neptune. It s atmosphere is thought to be mainly molecular N 2 and CH 4. At the time of this writing, Pluto is starting to get farther from the Sun: this will make the atmosphere cooler and hence, its atmosphere will start to condense. 1.2 Earth-Based and Planetary Missions By devising clever experiments and building sophisticated instruments, data are collected which can then be compared with calculations from theoretical modelling. This synergy between measurements and modelling furthers our understanding of the processes that govern the various behaviour of planetary atmospheres. 9

27 1.2.1 Ground Based and Satellite Observations There is perhaps a tendency to think that all our recent information on the planets results from planetary fly-bys, orbiters and landers. However, the planets have been explored from Earth since 1610 A.D. when Galileo first made observations of Jupiter using the telescope. Improvements to the telescope and better instrumentation to include wavelengths outside of the visible spectral region have facilitated new discoveries regarding our distant neighbours. Stars have also been used as an occultation tools to probe the atmospheres of the planets: acting as a light source the star is occulted by the planet as it passes behind, as viewed from the Earth. This can yield valuable information on the temperature structure of the mesospheres of the planets. This technique was in use in pre-spacecraft days and this method was how the rings of Uranus were first discovered in 1977, prior to the arrival of Voyager at Uranus. Most of the information obtained from Earth has been by use of passive remote sensing methods. However, the power of radars have been turned to our nearest neighbours to reveal surface information and rotation periods from radar returns from their surfaces. The Earth s atmosphere is opaque at many wavelengths except in the visible and near-ir with small windows in the mid-ir. Additionally, the atmosphere itself is not homogeneous and each little air parcel refracts visible light and distorts what is seen by the instrument. Ergo, it is desireable to get above the surface to mountain tops reduces the impact of the lower atmosphere. Better yet to go into space past 10

28 the atmosphere where the full range of the electromagnetic spectrum is available. This was first done using rockets to get above the Earth s atmosphere and for a few minutes observe at shorter wavelengths. Observations of far-uv and extreme-uv of Jupiter, Venus were obtained. However, rockets only offer tantalisingly brief glimpses of the planets. Space-based observatories have been present from the early seventies, lead by the astronomical community, mainly for the observation of stars and galaxies and the interstellar medium. However, planetary science has also benefited from their technology. One of the first observatories was the International Ultraviolet Explorer (IUE) launched in January, 1978 and which lasted more than 18 years until September, It was a combined mission between NASA, ESA and SERC (UK). It yielded much new information by extending the wavelength region accessible from Earth: it had instruments spanning the wavelength range 115 nm to 335 nm. The IUE made many planetary and cometary observations. One of the most interesting planetary observations was a long time series of Jovian aurorae. More recently the Hubble Space Telescope (HST) with its better spatial resolution and improved instrumentation have also probed the atmospheres of the planets and satellites. Its launch was delayed by the Shuttle tragedy in 1986, but it finally made it to orbit using the shuttle in April, Its prime targets are extra-solar system, nevertheless, it has made many exciting observations of the planets. For example, it has yielded information about aurorae on Jupiter, changing seasons on Saturn, and most recently, some of the most detailed Mars pictures ever taken from Earth. Its imaging spectrograph spans the wavelengths from 115 nm to 1000 nm (0.115 to

29 microns). The near-ir camera and multi-object spectrometer sense wavelengths from 0.8 to 2.5 microns. The Hopkins Ultraviolet Telescope (HUT) is a 90 cm telescope that samples the wavelength range 82 to 185 nm and has been used for Venusian observations. At shorter wavelengths, the Extreme Ultaviolet Explorer (EUVE) telescope with a wavelength grasp of 7 nm to 76 nm was launched June, 1992 and was operational until January, The ROSAT (Rontgen Satellite, June, 1990-February, 1999), a combined project with Germany, NASA and the UK is an X-ray observatory that was also used to observe Jupiter during the Shoemaker-Levy 9 impact. The Infrared Space Observatory (ISO) was a ESA satellite observatory for infared astrophysics. It was launched in 1995 and operated until its on-board supply of liquid helium was exhausted in April Its suite of four instruments provided imaging, photometry and spectroscopy at wavelengths from 2.5 to 240 micrometres. FUSE (Far Ultraviolet Spectroscopic Explorer) was launched in June, 1999 using high resolution spectroscopy in the wavelength range 90 to 120 nm. This was a joint mission between NASA and its partners Centre Nationale d Etudes Spatiale (CNES) and the Canadian Space Agency (CSA). The mission is operated out of John Hopkins University in Baltimore Planetary Missions An important means of exploring the solar system, for the past three decades, has been by spacecraft, either as fly-bys, orbiters, probes or some combination. In the last 30 to 40 years, our understanding of planetary science has been transformed by the 12

30 information sent back to Earth by these spacecraft as a result of the massive efforts of teams of engineers and scientists. Because of the weight and power constraints required to get objects into space, it is a necessity that the instrumentation used is small, light and require a minimum of power to operate. As a result, missions beyond Mars generally do not use solar power. For example, Voyager used a radioisotope thermoelectric generator to provide electrical power. Furthermore, the more distant from the Earth the more problems there are with sending data back to Earth. Generally two types of spacecraft have been used by NASA. The simpler but hardy Pioneer class, which is a continuously spinning spacecraft, has often been used. Another main class of spacecraft is the Mariner class, which flys with 3-axes stabilised but usually has a scan platform. This spacecraft is much more expensive than the Pioneer class probes. NASA often changes the name once the spacecraft has been successfully launched and thus the Jupiter-Saturn mission with two Mariner class spacecraft became the Voyager mission. The Pioneer 10 mission to Jupiter which was launched in 1972 and followed by the Pioneer 11 mission to Jupiter and Saturn in April These missions were designed to see how feasible it was to pass through the asteroid belt and, if successful, to probe the expected (with knowledge gained from earlier Radio observations) high-energy particle laden magnetosphere of Jupiter and the environs of the rings of Saturn. The missions were highly successful: Jupiter s magnetosphere was found to contain very high energy particles which could penetrate thin-shelled instruments and thicker protective casings were put around the electronics of the Mariner/Voyager 13

31 instruments. The Voyager mission with two spacecraft, originally began as a Mariner Grand Tour mission designed to visit all of the outer planets, Jupiter, Saturn, Uranus, Neptune (i.e. the giant planets) and Pluto. However, cut backs lead to de-scoping of the mission so that just Jupiter and Saturn were the designated targets with fly-bys of Uranus and Neptune as hopeful options. The Voyager mission turned out to be one of NASA s most successful missions even though there were a number of serious problems when the prime receivers failed, backup receivers faltered, and the scan platform on Voyager 2 stuck. Exceeding expectations, between them, Voyager 1 and Voyager 2 visited all the giant planets. Voyager 1 went to Jupiter in 1979 and was bombarded by the particles in Jupiter s magnetosphere and many instruments had to be turned off to avoid electrical problems. Spectacular pictures of the atmosphere and the satellites were taken by the cameras. Active volcanos were discovered on the Jovian moon Io, ejecting sulphur and oxygen to form an ionised plasma around the planet. The Galilean satellites presented themselves as four quite different worlds. Voyager 1 also discovered lightning on Jupiter and an ethereal ring of micron sized satellite debris. Galileo was launched from the Space Shuttle Atlantis in October 1989 and reached Jupiter in It consisted of an orbiter and probe. The probe entered Jupiter s atmosphere on December, 1995, and returned data from pressures to about 22 bar, deep within the cloud deck. In spite of a problem antenna and tape recorder, the orbiter revealed many details of Jovian storms, active volcanos on Io, and possible oceans on Europa and Callisto. 14

32 Cassini-Huygens is a joint NASA/ESA/ASI (Italian Space Agency) mission to Saturn and Titan. It was launched in October, 1997 and is due for orbit insertion at Saturn in July, It consists of a Saturn orbiter and the Huygens probe for Titan. 1.3 Thesis The goal of this thesis is to explore the photochemistry and radiative transfer properties of certain key species in the atmospheres of Jupiter and Saturn. The tools that are employed to yield important information about fundamental parameters that are indicators of the state of the atmosphere are examined in detail in Chapter 2 and Chapter 3. In the former, the emphasis is on describing the photochemical modelling of outer planetary atmospheres whereas, in the latter, the radiative transfer modelling of the atmosphere is the main topic. The model calculations, done utilising these tools, are discussed in the remaining chapters. Where applicable, the modelling efforts are related to relevant observations to further elucidate our understanding of these similar, yet different, outer planetary atmospheres. Chapter 4 contains a discussion of He 584 Å resonance radiation of Jupiter and Saturn with a brief extension to include Uranus and Neptune. This chapter mainly summarises the Saturnian work published by Parkinson et al. (1998), but also highlights contributions made to the Jovian case (Vervack et al, 1995). Chapter 5 and 6 include direct reproductions of the Jovian thermospheric D Lyman α publication (Parkinson et al., 2002a) and tropospheric CH 3 D publication (Parkinson 15

33 et al., 2002b), respectively. These two papers represent work done at York University and the Institut d Astrophysique de Paris to resolve questions about deuterium and deuterated hydrocarbons in the thermosphere and troposphere of Jupiter. In the former case, a brief review of H Lyman α resonance radiation is also included and and in the latter case, unpublished work on thermospheric deuterated ethane is also included. As with all theses, not all topics can be covered completely, leaving some questions unanswered or not completely resolved. Work to date concerning He 584 Å resonance radiation on Saturn has still failed to determine with certainty a definitive range of eddy diffusion values. Even with new estimates of the helium mixing ratio (Gautier and Owen, 2000), all we can say is that it would appear that K h cm 2 s 1 and thus still remains incompatible with that deduced from the solar and stellar occultation experiment analysis of Smith et al. (1983) and are more in line with the estimates of Sandel et al. (1982) and Atreya (1982). Attempting to estimate the D/H ratio in the Jovian thermosphere proved relatively fruitless but yielded interesting surprises regarding thermospheric physics. The tropospheric studies also yielded interesting speculations regarding the dynamics of this region which needs to be further explored. Future work would be greatly enhanced by obtaining a better estimate of the vibrational temperature profile in this part of the atmosphere. Pressure induced absorption by of H 2 in the Jovian troposphere and its possible effect on the effective temperature versus latitude has not been considered. The strong zonal and meridional mixing that may result conceivably could affect the isotropic 16

34 mixing ratio as described in Chapter 6. 17

35 Chapter 2 Vertical Mixing and Photochemical Modelling of Atmospheres Most of the planets in our solar system have a substantial atmosphere, with the exception of Mercury and Pluto which have very tenuous atmospheres. Several moons in the solar system also have atmospheres. An atmospheric gas can be made up of a variety of chemical species that were distributed unevenly at the time of the formation of the solar system. A basic relationship between fundamental quantities governing the gas distribution of chemical species is the ideal gas law, p = n(z)kt, which becomes increasingly less valid for pressures greater than 1 bar, after which van der Waals equation of state should be used. 18

36 2.1 General Method of Solution The vertical distribution of a minor constituent in a planetary atmosphere is governed by the 1-dimensional continuity equation for each species, i, where the vertical flux, φ i, can be approximated by n i t + φ i z = P i L i n i (2.1) φ i = φ K i + φ D i. (2.2) The eddy flux, φ K i, φ K i = K( n i z + ( T H av T z )n i) (2.3) represents the vertical flux that parameterizes macroscopic motions, such as the large scale circulation and gravity waves, and φ D i φ D i = D i ( n i z + (1 + α i) T T z + n i ) (2.4) H i is the vertical flux carried by molecular diffusion. The species number density is given by n i, P i is the chemical production rate (cm 3 s 1 ) and L i is the loss frequency (sec 1 ) at altitude z and time t (e.g., Chamberlain and Hunten, 1987). D i and K = K(z) are, respectively, the molecular and vertical eddy diffusion coefficients. The molecular diffusion coefficients, D i, are taken from Mason and Marerro (1970), Atreya (1986), and Cravens (1987) where applicable using the formula D i = b i n bg = AT s n bg where b is the binary collision parameter (expressed in terms of the coefficients A and s) and the subscript bg denotes background. H i and H av are respectively the constituent 19

37 and background atmospheric pressure scale heights, i.e., H i = kt M i g and H av = kt M avg where M i and M av are respectively the molecular weights of the constituent and the atmosphere. In these calculations we have neglected the effects of the thermal diffusion factor, α i, as its inclusion contributed less than 1% to a given species column in test runs. Eddy mixing tends to homogenize the atmosphere such that, where there are no effects due to chemistry, all species would be distributed according to the mean atmospheric pressure scale height. Molecular diffusion tends to separate constituents by their individual molecular weights. The atmospheric level at which the molecular diffusion coefficient is equal to the eddy diffusion coefficient is defined as the homopause for the i th constituent. Above this altitude, molecular diffusion dominates and the time constant for reaching diffusive equilibrium is given by τ D = H2 av D i (Chapman and Cowling, 1970; Colegrove et al., 1966). Below the homopause, eddy diffusion dominates and the long lived species are mixed, and the mixing time constant is analogously expressed as τ K = H2 av K. Equation (2.1) is solved using a finite central difference approximation for the vertical derivatives and the species densities are solved semi-implicitly in time using a simple tridiagonal solver. For these applications we have assumed a steady state exists and so have driven the solution so that n i t 0. The details of the model are given by Parkinson et al. (2002a) and in this chapter. Extensions to the original model (J. McConnell, private communication 1985) now include: additional species; updated reactions; reaction rates and quantum yields from the literature; improved calculation of J-values to account for a Rayleigh scattering atmosphere with a Lam- 20