1 Comparative Planetary Atmospheres1. Formation and evolution of planets and atmospheres Prof FW Taylor

2 Aim of Course To set the Earth in its context as a member of the Solar System To study the same physics in different atmospheres, and so gain further insight to the processes that control climate RECOMMENDED BOOKS: Lewis, J.S., The Physics and Chemistry of the Solar System, Academic Press, 1995. Chamberlain, J.W. and D.M. Hunten, Theory of Planetary Atmospheres, Academic Press, 2nd Edition 1987. Slides are on web (last year’s - this year’s at end of course). New textbook now in draft form.

3 Outline of Lectures (Weeks 1 & 2, Mondays and Fridays at 11am)Comparative Planetary Atmospheres Formation and evolution of the solar system, the planets, and their atmospheres The atmospheres of the terrestrial planets The atmospheres of the giant planets Measurements of temperature, composition and energy balance of planetary atmospheres Four lectures divided roughly like this

4 The Earth and the other planetsEarth is one of a family of 8 planets The planets form two groups: 4 small, rocky inner planets 4 large, gaseous outer planets

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7 Physical Properties of the PlanetsMercury Venus Earth Mars Jupiter Saturn Uranus Neptune Solar distance (AU) 0.387 0.723 1.0 1.524 5.2 9.5 19.2 30.1 Sidereal orbital period (years) 0.2408 0.6152 1.8808 11.9 29.5 84 165 Orbital Eccentricity 0.206 0.007 0.017 0.093 0.048 0.056 0.046 0.009 Equatorial radius Re (km) 2440.0 6052.0 6378.0 3397.0 71492 60268 25559 24764 Eq. radius relative to Earth. 0.38 0.95 0.53 11.2 9.4 4.0 3.9 Oblateness (Re-Rp)/Re 0.0065 0.065 0.098 0.023 Mass (1024 kg) 0.33 4.869 5.97 0.641 1898.8 568.5 86.625 102.78 Mass ( Relative to Earth) 0.055 0.816 0.107 318 95 14.5 17.1 Mean density (g/cm3) 5.43 5.24 5.515 3.94 1.33 0.70 1.27 1.76 Sidereal day 58.646d 243.02d 23h 56m 24h 37m 9h 55m 10h 39m 17h 14m 16h 6m Gravity at surface/1 bar (m/s2) 2.78 8.6 9.81 3.72 23.1 9.1 8.7 11.0 Escape velocity (km/s) 4.3 10.4 2.4 58.6 33.1 21.1 23.3 Obliquity 177.4° 23.5° 25.2° 3° 27° 98° 29° Go through definitions

8 Formation of Solar System & PlanetsCan we account for similarities and differences in the planets and their atmospheres? Is it Origin, Evolution, or Both? First need to understand how they were formed Contents of the Solar System Properties to be explained Formation of Protosolar Nebula Formation of Planets

9 Contents of the Solar System1 Star (99.95% of SS by mass) 8 Planets ~60 satellites & 1 (4?) ring system(s) ~6000 Asteroids (rocky) ~1 billion Comets, Oort cloud & Kuiper belt objects (icy) Gas & dust Magnetic fields & energetic particles Oort Cloud objects & Kuiper belt objects

10 Properties to be explainedPlanetary orbits are nearly all in the same plane (the ecliptic) Planetary orbits are roughly circular All of the planets orbit the Sun in the same sense Planets spin in the same sense as their orbits (not Venus, Uranus) Sun has almost all the mass, but planets (mostly Jupiter and Saturn) have most of the angular momentum of the solar system Rocky planets form a family close to the Sun; gas/ice giants form a family far from the Sun. Relative abundance of ‘heavy elements’ increases in the giant planets from 3 × solar for Jupiter to ~40 × solar for Neptune Space between planets increases with distance Solar system is transparent between mass concentrations Atmospheric compositions and isotopic ratios are different. Mars orbit non-circular by about 10% Bode’s Law or Titius-Bode Law Planets have different mixes of elements and isotopes

11 Formation of the Solar NebulaThe universe contains a lot of ‘loose’ material, mostly hydrogen and helium, but also some (~2%) heavier elements from old stars Any local accumulation of mass will tend to grow, forming a region of locally enhanced density (a ‘molecular cloud’) in space About 4,600 million years ago a molecular cloud collapsed under its own weight to form the Sun and planets, including Earth First sums for gravitational collapse worked out by Sir James Jeans One of the oldest problems in science - contributions by Kepler, Newton, Halley, Laplace, Faraday………… Age of SS comes from radiometric dating of meteorites Born: 11 Sept 1877 in Ormskirk, Lancashire Died : 16 Sept 1946 in Dorking, Surrey

12 Existence of snow line

13 Formation of Protosolar NebulaOrder of magnitude calculation - detailed models exist of course

14 Molecular Clouds Molecular clouds are observed in the UniverseTypical values for temperature and density are: T = 20K, r ~ 1010 H atoms (10-17 kg) m-3 gives MJ ≈ 1031kg ≈ 10 solar masses. R ≈ 40,000 AU or about 1 light year. Molecular Cloud Barnard 68 (VLT)

15 Formation of Circumstellar Disc

16 Planet Formation Gravity is not important until aggregates reach planetesimal size (definition of a planetestimal)

17 Circumstellar Discs Emission lines OIII, H alpha and NII

18 Loss of mass and angular momentum in circumstellar discs: The T-tauri phase of the SunProblems: MJ is >> present mass of Solar System. Distribution of angular momentum between Sun and planets. Answer: Soon after the Sun began to fuse hydrogen it entered its ‘T-tauri’ phase with ~ 3 x current luminosity and a very dense, high speed solar wind. Mass loss of 10-8 MSun/year over 107 years. Planets formed before the T-tauri phase. Remaining solar nebula was swept away. Angular momentum carried away by the solar wind, ‘despinning’ the Sun. How to (1) get rid of excess mass (ii) redistribute angular momentum (iii) form planets. Angular momentum removal involved Sun’s magnetic field.

19 Formation of Planets I: Encounter TheoriesPropounded by Jeans (and Maxwell) but not generally accepted today Basic idea Planets pulled off sun by collision with a comet or by gravitational attraction of passing star. Explains Common direction of planets' orbital motion Planets' nearly circular and coplanar orbits Disadvantages Hot gas could not have condensed into planets Probability of a near encounter in our region of the Galaxy is vanishingly small, less than one in many millions Propounded by Jeans (and Maxwell)

20 Formation of Planets II: Condensation and Collapse of Solar Nebula into Planetesimals** a planetesimal is a chunk of matter large enough to affect others through gravitation (~a few km across). The theory explains: The varying composition and size of planets: Inner planets formed where the nebula was hotter and so only rocky materials could condense. Further from the Sun, volatiles especially water could also condense. Planetisimals in the outer Solar System grew quickly and could accrete hydrogen and helium before T-Tauri phase removed remaining gas. Jupiter and Saturn grew more quickly than Uranus and Neptune so the latter are more ice-rich. The regular motions of the planets and moons: all revolve in the nearly same plane, in nearly circular orbits, in same direction the Sun rotates. Small bodies (asteroids and comets, KBOs, moons and rings) Exceptions to the trends (Earth’s Moon, axial tilts, eccentricities) Problem: recent discovery of ‘Hot Jupiters’ around other stars A planetesimal is a chunk of matter large enough to affect others through gravitation. ~ A few km across.

21 Formation of Terrestrial Planet AtmospheresDid the atmosphere: form with the planet out of the solar nebula? outgas later from the interior? accumulate from the solar wind? arrive later as icy meteorites and comets? Obtain clues from the relative abundances and isotopic ratios of the noble gases, allowing that some of these are of radiogenic origin. For Venus, Earth and Mars it is found that: the ratio of 20Ne to 36Ar is similar on all 3 planets, but different in the Sun: argues against (1) and (3) primordial argon decreases by several orders of magnitude from Venus to Earth and from Earth to Mars. Argues against (4). This leaves (2). Plus, outgassing is still observed (e.g. volcanoes).

22 Formation of Outer Planet AtmospheresUnlike the terrestrial planets, the gas giants were too massive, cold, and distant from the Sun to have lost their original atmospheres If so, the giant planets are made up of primitive material from the solar nebula Both Jupiter and the Sun are ~85% hydrogen, ~15% helium The H2/He ratios on all four planets resemble that of the Sun Heavier element abundances and noble gas ratios are difficult to measure and interpret because of condensation, chemistry, interior processes etc. They remain controversial research topics that are slowly yielding a detailed picture of the evolution of the Solar System.

23 Processes affecting the evolution of atmospheres to their present stateThermal escape to space* Condensation, e.g. on permanent polar caps or as permafrost below the surface Dissolve in oceans & subsequent removal, e.g. carbonate formation removes CO2 on Earth Regolith absorption/chemical combination, e.g O2 ➔ rust Hydrodynamic escape (lighter atoms move heavier ones) Solar wind erosion (especially if no mag. field, Venus & Mars) Impact erosion* (incoming mass blasts gases into space) Sources (e.g. comets, volcanism) (Earth's CO2 would be removed in ~10,000 years if not replaced.)

24 Thermal Escape: Jeans' Formula (1)Exobase is ~100 km on Earth.

25 Thermal Escape: Jeans' Formula (2)For derivation see Chamberlain and Hunten (1987).

26 Characteristic time for reduction of exospheric density by Jeans escape for Earth’s atmosphereScale Height (km) 1460 730 365 91 Most probable speed (km s-1) 4.96 3.51 2.48 1.24 Mean expansion velocity (km s-1) 7.32 × 10-2 8.71 × 10-4 9.94 × 10-8 7.14 × 10-32 Exospheric escape time (s) 2.0 × 104 8.38 × 105 3.67 × 109 1.28 × 1033 (Assumes exobase is at 500km altitude with temperature = 1480K)

27 Characteristic Jeans escape times for different gases on several planets.Moon Mercury Mars Venus Jupiter T (K) 300 600 365 700 155 Re (km) 1738 2439 3590 6255 69500 g (ms-2) 1.62 3.76 3.32 8.27 26.2 te(H) (s) 3.55 × 103 3.32 × 103 1.39 × 104 5.71 × 105 5.14×10617 te(He) (s) 2.03 × 104 1.40 × 105 2.66 × 108 2.85 × 1016 1.18×102455 te(O) (s) 2.25 × 109 7.37 × 1013 1.04 ×1028 7.87 × 1061 1.03×109820 te(Ar) (s) 3.29 × 1020 2.57 × 1032 1.97 × 1068 6.20×10153 6.61× te(Kr) (s) 3.53 × 1041 9.09 × 1066 4.45×10142 4.67×10322 3.72× t

28 Impact Escape • The inner planets suffered a period of heavy bombardment and the energy imparted can drive off some of the atmosphere. • This is most important for Mars and in that case even extends to driving off solid material (the SNC meteorites). If Me is the mass of gas driven off by an impactor of radius R then: where Ma = mass of atmosphere per unit area, ve = (2GM/r)1/2 is the escape velocity, and vs is the impact velocity. N.B. molecular mass is not a factor so no fractionation of species. 1/2 M ve**2 = mMG/r (r = radius of planet. Above, R = radius of impactor)

29 Comparative Planetary AtmospheresII. Terrestrial Planets

30 Mercury Diameter 1.4 times MoonMuch denser than Moon: 5.43 vs g cm-3 Temperature range 70 to 700 K Thin atmosphere: surface pressure ~10-15 bar Icy polar deposits in shaded craters

31 Venus Solid body resembles EarthSmall inclination and eccentricity – no seasons Complete cloud cover of mainly 75%H2SO4.25%H2O. No liquid water & very little vapour Surface temperature ~ 730 K Net insolation < Earth! Equilibrium temperature ~ 240 K 500K greenhouse effect (Earth ~ 30K) Very thick CO2 atmosphere km-atm of CO2 (Earth: 10-3) Surface pressure 92 bars.

32 Earth Water in all three phases Widespread water clouds70% liquid H2O coverage N2 – O2 atmosphere Surface pressure 1 bar Mean surface temperature 288 K Life is part of climate

33 Mars Thin CO2 atmosphere Thin CO2 and H2O clouds.Surface Pressure ~7 mb (variable) Surface temperature 218 K (very variable)

34 Radiative equilibrium temperature of a PlanetApplying the Stefan-Boltzmann law we obtain for the total radiant power of the Sun: ESun = 4πσ RS2 (TS)4 (W) where σ is equal to x 10-8 W m-2 K-4. The solar constant S, is then given by S = σ TS4 (RS/DS)2 ( = kW m-2 for DS = 1AU) where DS is the planet’s distance to the sun in AU. For equilibrium EE = 4πσRE2(TE)4 = (1-A)SπRE2 (W), whence the radiative equilibrium temperature TE, given the albedo A.

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36 Calculating Model Vertical Temperature ProfilesAssumptions No absorption of sunlight in the atmosphere, only at the surface, i.e. atmosphere optically thin at short wavelengths (< 4 µm). Most of the energy emission to space takes place from a narrow range of altitudes near the tropopause Troposphere is optically thick, stratosphere is optically thin, in the mid and far infrared (> 4 µm)

37 Tropospheric Lapse RateConsider a ‘parcel’ of dry air at pressure p and temperature T rising adiabatically, i.e. so that p is that of its surroundings while T may change. From the first law of thermodynamics for one mole: dQ = CvdT + pdV = 0 And from the perfect gas law pdV + Vdp = RdT = (Cp – Cv)dT Using the hydrostatic equation, for balance between rising and falling CpdT = Vdp = -V gdz = -Mgdz Hence where Γ is called the adiabatic lapse rate.

38 Stratospheric TemperatureThe observed lapse above the tropopause, where convection stops, tends to zero (i.e. constant temperature with height) Each layer is heated by radiation from the optically thick atmosphere below, and cooled by radiating to space, to the same degree; to first order height is no longer important. The stratospheric temperature TS may be estimated by treating the region as if it were a single slab of gas which is optically thin at all wavelengths, rather than just on average, which is the real situation. TS is related to the effective radiative temperature TE by the expression for the energy balance of the stratosphere, treated as a slab of emissivity and absorptivity ε: εσ (TE)4 = 2εσ (TS)4 whence TS = TE/21/4

39 A Simple Greenhouse model (Earth)Transparent λ < 4µm Opaque λ > 4µm T0 Need a further fixed point to make profile unique Represent the atmosphere as a homogeneous slab at temperature Ta overlying the surface at temperature T0 The atmosphere is completely transparent to solar radiation (at λ<4µm) and completely opaque to thermal radiation (λ>4µm). At the surface, σ(T0)4 = 2σ (255)4 so T0 = x 255 = 303 K.

40 Radiative equilibrium of terrestrial planetsVenus Earth Mars Distance from Sun (Mkm) 108 150 228 Solar irradiation (kW/m2) 2.62 1.38 0.60 Bolometric (Bond) albedo, A 0.76 0.3 0.15 Radiative equilibrium temperature, TE (K) 239 255 226 Stratospheric temperature (K) 201 214 190 Tropospheric Lapse Rate (K km-1) 10.7 9.8 4.5 Surface temperature model/measured (K) 285/730 303/288 255/235

41 Multilayer radiative equilibrium model for Venus

42 The Temperature profile on VenusRadiometric temperature = cloud top temperature Cloud top pressure ~ 0.1 bar Surface pressure ~ 100 bar Pressure scale height H = RT/Mg, ≈ Earth so ≈ 7 km Lapse rate = g/cp ≈ 10 K km-1 Cloud height z ≈ ln (0.1/100)H ≈ 50 km Surface temperature ≈ *10 = 740 K DD

43 Model temperature profiles

44 Terrestrial Planet temperature profilesProfiles based on measurements match simple theory ozone heating on Earth dissociation & ionisation heats thermospheres Cool thermospheres on Mars/Venus - more CO2

45 Radiative-Convective Model Temperature ProfileHEIGHT Optical depth changes TEMPERATURE

46 Radiative-Convective Model Temperature ProfileHEIGHT Albedo changes TEMPERATURE

47 Radiative-Convective Model Temperature ProfileHEIGHT Optical depth changes Albedo changes TEMPERATURE

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49 Composition of Terrestrial Planet AtmospheresVenus Earth Mars Carbon dioxide 0.96 0.003 0.95 Nitrogen 0.035 0.770 0.027 Oxygen ~0 0.21 0.0013 Water vapour ~ ~0.01 ~0.003 Sulphur dioxide 150 ppm 0.2 ppb Carbon monoxide 0.12 ppm 0.007 Surface Pressure 92 bar 1 bar 0.007 bar Clouds H2SO4 H2O H2O, CO2

50 Key Questions: (1) CO2 Terrestrial planet atmospheres are secondary, i.e. outgassed from the interior Suppose all three atmospheres were about the same at first, as indicated (with complications) by chemically inactive gases e.g. N2 The primordial gases were primarily CO2 and H2O, plus smaller amounts of nitrogen (possibly as ammonia), argon &c, SO2 Most of Mars’s CO2 lost by impact, perhaps a few very large impacts Venus lost its liquid H2O and kept its CO2 in the atmosphere Earth’s CO2 was dissolved in oceans and became carbonate rocks

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52 Key Questions: (2) N2 primordial gases were primarily CO2 and H2O, plus smaller amounts of nitrogen, argon, SO2 etc Nitrogen is inert so either still there (Earth, Venus) or physically removed from planet (Mars) Most likely loss mechanism on Mars is impact Isotopic ratios suggest ~99% of Martian atmosphere lost by non-thermal escape

53 Key Questions: (3) H2O Water - all 3 planets probably very wet to start with Oceans remove CO2 to produce carbonates on Earth Did this happen on Mars? Water (as H and O) escapes from Venus after dissociating Water on Mars froze (?) Surface pressure on Mars is close to the triple point of water

54 VEx SOIR: D/H on Venus is 240+25 times that on Earth

55 Key Questions: (4) SO2 Commonest S compound, volcanic in originSulphates on Mars surface confirm early volcanic phase Absence of sulphurous gases in Martian atmosphere now indicates low current volcanism On Earth sulphur is released from fossil fuels (anthropogenic) and affects global warming/cooling Atmospheric sulphur (mostly as SO2 & H2SO4) on Venus is >> Earth, by a factor 100,000 in stratosphere, and more than 10 million in troposphere. Atmospheric sulphur on Venus >> equilibrium value with surface Implies Venus very active volcanically in last ~106 years

56 SO2 abundance at 40 mbar, ppbupper limits YEAR X - VIRTIS 35 km

57 ← Venera 9 obtained the first image from the surface of Venus in 1975.↓ View of a plain in Phoebe Regio from Venera 13 on March 1, 1982

58 Maat Mons from Magellan Radar Image ca. 1991Egyptian goddess of truth and justice 8 km above mean radius (Maxwell = 12 km, Everest = 8.85 km) Maat Mons from Magellan Radar Image ca. 1991

59 Venus has more volcanoes than any other planet in the solar systemVenus has more volcanoes than any other planet in the solar system. Over 1600 major volcanoes or volcanic features are known (see map), and there are many, many more smaller volcanoes. No one has yet counted them all, but the total number may be over 100,000 or even over 1,000,000. [Oregon State University]

60 Key Questions: (5) CO CO2 is rapidly converted to CO by solar UV on all three planets CO2 + hν → CO + O, λ< 224 nm OH recycles CO on Earth H2O + hν → OH + O, λ< 190 nm. 2OH + CO → CO2 + H2O Mars probably similar to Earth Venus: (i) depletion in clouds due to OH, (ii) removal by S compounds in lower atmosphere and at surface, e.g. FeS2 + 2H2O + CO → FeO + 2H2S + CO2

61 Key Questions: (6) Surface Pressure ScenariosMars’ atmospheric density and hence temperature may have been reduced by erosion , while the source due to volcanism also declined, resulting in net loss. Eventually, first H2O and then CO2 could form ice caps at the poles, reducing the mean surface pressure and temperature still further. In the warmer equatorial regions, liquid water persisted for a time and removed CO2 as carbonates. It may be significant that Mars’ atmospheric surface pressure has stabilized close to the triple point of water (6.1 mb), since then carbonate formation would cease. On Venus something similar may occur via the Urey reaction (see next) On Titan the surface temperature & pressure are close to the triple point of methane (see later)

62 Surface Pressure on VenusPerhaps CO2 is in equilibrium with common minerals on surface (Urey, 1952) - The dominant reaction is expected to be CaCO3 + SiO2 ⇔ CaSiO3 + CO2 Use thermodynamic data (e.g. Adamik and Draper, 1963) for the temperature dependence of the pressure of CO2 in the form log P = ΔH/RT +ΔS/R + A + BT + CT-2 (P= pressure, T= temperature, ΔH, ΔS enthalpy and entropy changes) Compare this to the radiative-convective equilibrium model for the surface pressure as a function of surface temperature Harold C. Urey: Nobel Prize in Chemistry Nobel Lecture, February 14, 1935: Some Thermodynamic Properties of Hydrogen and Deuterium

63 Contact between reactants neededVenus: Radiative-convective equilibrium and geochemical equilibrium possible together at 92 bars and 735K! BUT…. Very simple model Contact between reactants needed Equilibrium is unstable for Urey reaction alone At 1000 K, P(atm) = 323 bar and P(Urey) = 3,100 bar Problems (1) existence/stability of equilibrium (2) access to minerals without water (3) other reactions eg Fe (4) volcanoes - SO2 is not in equilibrium

64 Climate change models for VenusEarthlike albedo A=0.3 Present Venus p0 = 92 bar T0 = 735 K Earthlike surface pressure and cloud p0 = 1 bar T0 = 375 K Surface 400 K

65 Climate Change on Mars Geological evidence that Martian atmosphere was much thicker in the distant past Mars was warm and wet What changed & why? Mars is the only planet except Earth where we have detailed evidence for major climate change

66 Evidence for liquid water - fluvial features (Nanedi Valles)

67 Global altitude map of surface by laser altimeter on orbiterEvidence for liquid water - coastlines of paleo-ocean? Global altitude map of surface by laser altimeter on orbiter

68 Olympus Mons Mariner ValleyWhat changed? 1. Massive volcanic and tectonic features suggest early geophysical activity, now dormant Olympus Mons 27 km high (c.f. 11km for Everest) Mariner Valley 100km wide, 10km deep, 4800 km long

69 What changed? 2. Mars’ residual magnetic field measured from orbit shows it once had a global field, but not now

70 What changed. 3. Collisions stripped away >95% of atmosphereWhat changed? 3. Collisions stripped away >95% of atmosphere. Solid debris reached Earth. Very large impact features still visible on surface e.g. Hellas.

71 What changed? 4. Milankovitch Cycles are large for Marslarge eccentricity and obliquity, together with precessions, cause variations in solar forcing layers everywhere indicate period climate change in response

72 Mars model temperature profilesMariner 9 Calculated temperature profiles for the atmosphere of Mars by Gierasch and Goody (1972), at two local times of day, 0600 (morning) and 1800 (evening). The left-hand pair of profiles was calculated assuming an atmosphere without suspended dust; the right-hand pair includes the radiative effects of a model dust profile. The shaded area encloses the range of measured temperature profiles obtained by Mariner 9. The temperature profiles from the present, simple model (thick lines) show the global mean for present conditions, (a) without and (b) with suspended dust in the atmosphere, and (c) possible past conditions where the surface pressure p0 was higher, and the temperature was warm enough for liquid water to be present. Gierasch and Goody (1972)