EART164: PLANETARY ATMOSPHERES

1. EART164: PLANETARY ATMOSPHERES Francis Nimmo

2. Course Overview • How do we know about the gas envelopes of planetary bodies? Their structure, dynamics, composition and evolution. • Techniques to answer these questions • Remote sensing (mostly) • In situ sampling • Modelling • Case studies – examples from this Solar System (and exoplanets)

3. Course Outline • Week 1 – Introduction, overview, basics • Week 2 – Energy balance, temperature • Week 3 – Composition and chemistry • Week 4 – Clouds and dust • Week 5 – Radiative Transfer; Midterm • Week 6 – Dynamics 1 • Week 7 – Dynamics 2 • Week 8 – Exoplanets • Week 9 – Climate change & Evolution • Week 10 –Recap; Final

4. Logistics • Website: http://www.es.ucsc.edu/~fnimmo/eart164 • Set text –F.W. Taylor, Planetary Atmospheres (2010) • Another good reference (higher level) is Lissauer & DePater, Planetary Sciences 2nd ed. (2010), Chs. 3&4 • Prerequisites – some knowledge of calculus expected • Grading – based on weekly homeworks (~40%), midterm (~20%), final (~40%). • Homeworks due by 5pm on Monday (10% penalty per day) • Location/Timing –MWF 2:00-3:10in E&MS D236 • Office hours – MWF 3:15-4:15 (A219 E&MS) or by appointment (email: fnimmo@es.ucsc.edu) • Questions? - Yes please!

5. Expectations • I’m going to assume some knowledge of calculus • Homework typically consists of 3 questions • If it’s taking you more than 1 hour per question on average, you’ve got a problem – come and see me • Midterm/finals consist of short (compulsory) and long (pick from a list) questions • Showing up and asking questions are usually routes to a good grade • Plagiarism – see website for policy.

6. This Week • Introductory stuff • Overview/Highlights (Taylor Ch. 1) • How do planets form? (Taylor Ch. 2) • Where do atmospheres come from? • What observational constraints do we have on atmospheric properties? (Taylor Ch. 3) • Introduction to atmospheric structure

7. Other solar systems will certainly contain planets very different from ours (super-Earths, mini-Jupiters, iron planets . . .) GJ876d HD149026b Three classes of planetary bodies “Rock” 1 ME 300 GPa ~6000 K Ice + H,He ~15 ME 800 GPa ~8000 K Mainly H,He ~300 ME 7000 GPa ~20,000 K “Rock”+ice ~0.1 ME ~10 GPa ~1500 K

8. Useful Data Data mostly from Taylor, Appendix A

9. Units! • SI in general but • 1 bar = 105 Pa • g/cc vs. kg/m3 • Per mol vs. per kg

10. Overview/Highlights

11. Venus • Thick CO2 atmosphere • Hot (“runaway greenhouse”) • Cloud-covered • Lost a lot of water • Slow rotator (retrograde), not tilted • Fast winds (“superrotation”) • Sulphur cycle (active volcanism) • Pioneer Venus, Venera & Vega probes (USSR), Magellan, Venus Express (ESA)

12. Earth • Mostly N2,O2 • Moderate greenhouse • Hydrological cycle & oceans • Weathering buffer • Moderate rotator • Tilted (seasons) • Hadley cell • Milankovitch cycles • Biological activity

13. Mars • Thin CO2 atmosphere • Dust and polar caps important • Massive climate change • Moderate rotator • Tilted (seasons) • Global dust storms • Orbital forcing important (Milankovitch cycles) • Mars Odyssey, Mars Express (ESA), Mars Exploration Rovers, Mars Science Laboratory, MAVEN

14. Jupiter & Saturn • Thick H/He atmospheres • ~10 Earth mass rock/ice cores • Internal energy sources • Rapid rotators • Saturn is tilted • Banded winds + storms • Multiple cloud layers • Voyagers, Cassini, Galileo, Juno (we hope)

15. Uranus & Neptune • Thin (relatively) H/He atmos. • Massive rock/ice cores • Rapid rotators • Banded winds + storms • Multiple cloud layers • Uranus is tilted (seasons) • Poorly understood • Voyagers

16. Titan • Moderate N2 atmosphere • “Hydrological” cycle (methane) • Subsurface replenishment • Moderate rotator • Saturn tilted (seasons) • Local clouds and storms • Large atmospheric loss? • Voyager, Cassini/Huygens

17. Thin Atmospheres

18. Exoplanets Swain et al. 2008

19. 1. How do planets form?

20. Solar System Formation - Overview • Some event (e.g. supernova) triggers gravitational collapse of a cloud (nebula) of dust and gas • As the nebula collapses, it forms a spinning disk (due to conservation of angular momentum) • The collapse releases gravitational energy, which heats the centre • The central hot portion forms a star • The outer, cooler particles suffer repeated collisions, building planet-sized bodies from dust grains (accretion) • Young stellar activity (T-Tauri phase) blows off any remaining gas and leaves an embryonic solar system • These argument suggest that the planets and the Sun should all have (more or less) the same composition

21. Sequence of events • 1. Nebular disk formation • 2. Initial coagulation (~10km, ~105 yrs) • 3. Orderly growth (to Moon size, ~106 yrs) • 4. Runaway growth (to Mars size, ~107 yrs), gas blowoff • 5. Late-stage collisions (~107-8 yrs)

22. What is the nebular composition? • Why do we care? It will control what the planets (and their initial atmospheres) are made of! • How do we know? • Composition of the Sun (photosphere) • Primitive meteorites (see below) • (Remote sensing of other solar systems - not yet very useful) • An important result is that the solar photosphere and the primitive meteorites give very similar answers: this gives us confidence that our estimates of nebular composition are correct

23. Solar photosphere • Visible surface of the Sun • Assumed to represent the bulk solar composition (is this a good assumption?) • Compositions are obtained by spectroscopy • Only source of information on the most volatile elements (which are depleted in meteorites): H,C,N,O 1.4 million km Note sunspots (roughly Earth-size)

24. Primitive Meteorites • Meteorites fall to Earth and can be analyzed • Radiometric dating techniques suggest that they formed during solar system formation (4.55 Gyr B.P.) • Carbonaceous (CI) chondrites contain chondrules and do not appear to have been significantly altered • They are also rich in volatile elements • Compositions are very similar to Comet Halley, also assumed to be ancient, unaltered and volatile-rich chondrules 1cm

25. Meteorites vs. Photosphere • This plot shows the striking similarity between meteoritic and photospheric compositions • Note that volatiles (N,C,O) are enriched in photosphere relative to meteorites • We can use this information to obtain a best-guess nebular composition Basaltic Volcanism Terrestrial Planets, 1981

26. Nebular Composition • Based on solar photosphere and chondrite compositions, we can come up with a best-guess at the nebular composition (here relative to 106 Si atoms): • Blue are volatile, red are refractory • We would expect planetary atmospheres to consist primarily of H, He, C,N,O,Ne, Ar and their compounds Data from Lodders and Fegley, Planetary Scientist’s Companion, CUP, 1998 This is for all elements with relative abundances > 105 atoms.

27. Temperature and Condensation Nebular conditions can be used to predict what components of the solar nebula will be present as gases or solids: Mid-plane Photosphere “Snow line” “Snow line” Saturn (~50 K) Earth (~300K) Condensation behaviour of most abundant elements of solar nebula e.g. C is stable as CO above 1000K, CH4 above 60K, and then condenses to CH4.6H2O. From Lissauer and DePater, Planetary Sciences Temperature profiles in a young (T Tauri) stellar nebula, D’Alessio et al., A.J. 1998

28. “Snow line” • Beyond the “snow line” (~180 K), water ice condenses • Ice is ~10 times more abundant (by mass) than rock in the solar nebula • So it is much easier to build big planets beyond the snow line • Gas giants need a big solid core to start accumulating H or He (see next slide) • Close-in exoplanets almost certainly formed beyond the snow line and then migrated

29. Gas/ice giant formation • Once a solid planet gets to ~10 Earth masses, its gravity is large enough to trap H2 and He present in the local nebula • J,S,U and N all have cores made of “high-Z” elements (rock+ice) • J,S have thick H/He envelopes; U,N have thin H/He envelopes • So the cores of J&S probably grew early enough to trap nebular H/He before it dissipated. U&N were too slow. Why?

30. Migration (hot Jupiters) • If the gas disk is still present, planets will migrate inwards • This migration can be very rapid (~104-105 yrs) • Migration stops where the disk stops (e.g. due to stellar magnetic fields) • This is why there are so many “hot Jupiters” • But it apparently didn’t happen in our solar system planet Gas disk (with density waves)

31. “Hot” population Present day Planetesimals transiently pushed out by Neptune 2:1 resonance “Cold” population J N U S Neptune stops at original edge 3:2 Neptune resonance (Pluto) 2:1 Neptune resonance See Gomes, Icarus 2003 and Levison & Morbidelli Nature 2003 Nice Model Early in solar system Ejected planetesimals (Oort cloud) “Hot” population Initial edge of planetesimal swarm J N S U 18 AU 30 AU 48 AU

32. Planet Formation - Summary • Initial nebular composition is well-known • Planetary volatile abudance depends (mostly) on where the planet formed (temperature) • Timing of planet growth relative to nebular blowoff also important • The planets may have moved during or after the formation phase

33. 2.Where do atmospheres come from?

34. Where do atmospheres come from? • Primary – directly accreted from nebula • Secondary – outgassed from planet • Tertiary – derived from comets, asteroids and/or solar wind • We’ll discuss more later in the quarter. Examples: • Does Earth’s hydrosphere come from comets or asteroids? (D/H ratio) • How much outgassing has there been on Earth, Venus, Mars, Titan? (40Ar) • Did the gas giants acquire a solar composition? (C/H, H/He)

35. Where do atmospheres go to? • Again, we’ll discuss more later, but there are several processes which can remove atmospheres • Loss to space • Thermal processes (Jeans escape) • Hydrodynamic escape • Sputtering & photodissociation • Impacts • Loss to surface/interior • Chemical reactions (e.g. carbonate formation) • “Ingassing” (e.g. plate tectonics) • Freeze-out (Mars, Pluto)

36. 3. Observational constraints(see Taylor ch.3)

37. Radiometry (Spectroscopy) • “Near” infra-red: 0.7-5 mm, reflected sunlight • “Thermal” IR: 5-1000 mm, emission from atmosphere • Absorption/emission tells us what species are present, and where in the atmosphere they are • Background spectrum (~black body) tells us about temperature structure of atmosphere

38. Radiometry (cont’d) • We can see to different depths within an atmosphere by using different wavelengths • By looking at emission from the limb, we can probe the vertical temperature and pressure structure Haze layer

39. Occultations observer • Atmospheric absorption of light/radio waves provides information on composition, pressure and temperature observer • Good for probing thin atmospheres (e.g. Pluto, Enceladus) Hansen et al. 2006

40. In situ sampling • Galileo probe (Jupiter) • Huygens (Titan) • Venera/Vega probes/balloons • Viking landers • Cassini INMS LeBreton et al. Nature 2005 • Very useful! Ground truth for pressure, wind, temperature etc. Sensitive to trace gases (GCMS). • Generally limited duration (e.g. Venus) • Point measurement – what happens if you land in an anomalous region? (Galileo probe)

41. 4. Atmospheric structure

42. Typical structure z stratosphere Temperature structure of stratosphere in reality can be more complicated because of photochemistry (e.g. ozone) tropopause troposphere T Lower atmosphere consists of a thick part (troposphere) where convection dominates, and a thinner part above (stratosphere) where radiation dominates

43. Ideal Gas Equation P=pressure, r=density, R=gas constant, T=temperature (in K), m=molar mass (in kg) What is density of air at Earth’s surface? What is the column mass of Earth’s atmosphere? (kg/m2)

44. Atmospheric Structure (1) • Atmosphere is hydrostatic: • Gas law gives us: • Combining these two (and neglecting latent heat): Here R is the gas constant, m is the mass of one mole, and RT/gmis the pressure scale height of the (isothermal) atmosphere (~10 km) which tells you how rapidly pressure decreases with height e.g. what is the pressure at the top of Mt Everest? Most scale heights are in the range 10-30 km

45. Exobase and mean free path • The exobase is the place where the mean free path of molecules exceeds scale height. This is where molecules can start to escape efficiently (if travelling fast enough) • You can think of the exobase as the effective “top” of the atmosphere • For planets with thin atmospheres, the exobase may be at the surface! What’s the mean free path at the surface of the Earth? prmol2 l rmol is typically 1 Angstrom=10-10 m

46. Key concepts • Snow line • Migration • Troposphere/stratosphere • Primary/secondary/tertiary atmosphere • Emission/absorption • Occultation • Scale height • Hydrostatic equilibrium • Exobase • Mean free path First homework due next Monday!

47. End of lecture

48. An Artist’s Impression gas/dust nebula The young Sun solid planetesimals

49. Observations (1) • Early stages of solar system formation can be imaged directly – dust disks have large surface area, radiate effectively in the infra-red • Unfortunately, once planets form, the IR signal disappears, so until very recently we couldn’t detect planets (now we know of ~400) • Timescale of clearing of nebula (~1-10 Myr) is known because young stellar ages are easy to determine from mass/luminosity relationship. This is a Hubble image of a young solar system. You can see the vertical green plasma jet which is guided by the star’s magnetic field. The white zones are gas and dust, being illuminated from inside by the young star. The dark central zone is where the dust is so optically thick that the light is not being transmitted. Thick disk

50. Observations (2) • We can use the present-day observed planetary masses and compositions to reconstruct how much mass was there initially – the minimum mass solar nebula • This gives us a constraint on the initial nebula conditions e.g. how rapidly did its density fall off with distance? • The picture gets more complicated if the planets have moved . . . • The observed change in planetary compositions with distance gives us another clue – silicates and iron close to the Sun, volatile elements more common further out