EART164: PLANETARY ATMOSPHERESfnimmo/eart164/Week1.pdf · F.Nimmo EART164 Spring 11 Expectations...

F.Nimmo EART164 Spring 11

EART164: PLANETARY

ATMOSPHERES

Francis Nimmo


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)


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


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 2:00-3:10 in E&MS D236

• Office hours – MWF 3:15-4:15 (A219 E&MS) or by appointment (email: [email protected])

• Questions? - Yes please!

mailto:[email protected]


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.


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


Three classes of planetary bodies

“Rock”

1 ME

300 GPa

~6000 K

“Rock”+ice

~0.1 ME

~10 GPa

~1500 K

Ice + H,He

~15 ME

800 GPa

~8000 K

Mainly H,He

~300 ME

7000 GPa

~20,000 K

HD149026b

Other solar systems will

certainly contain planets

very different from ours

(super-Earths, mini-

Jupiters, iron planets . . .)

GJ876d


Venus Earth Mars Titan Jupiter Saturn Uranus Neptune

Solar constant (Wm-2) 2620 1380 594 15.6 50.5 14.9 3.7 1.5

Obliquity (o) 177 23.4 24.0 (27) 3.1 26.7 98 28.3

Orbital period (years) 0.62 1 1.88 (29.4) 11.9 29.4 84 165

Rotation period (hours) 5832 24 24.6 383 9.9 10.7 17.2 16.1

Bond albedo A 0.76 0.4 0.15 0.3 0.34 0.34 0.3 0.29

Molecular wt. m (g/mol) 43 29 43 29 2.2 2.1 2.6 2.6

Tsurface or T1bar (K) 730 288 220 95 165 134 76 72

Surface pressure (bar) 92 1 .007 1.47 n/a n/a n/a n/a

g (ms-2) 8.9 9.8 3.7 1.35 24.2 10.0 8.8 11.1

Teq (K) 229 245 217 83 113 84 60 48

Scale height H (km) 15 8.5 12 23 27 60 28 20

Radius (km) 6052 6370 3390 2575 71,500 60,300 25,000 24,800

Mass (1024 kg) 4.87 5.97 0.64 0.13 1900 568 87 102

Useful Data

Data mostly from Taylor, Appendix A


• SI in general but

• 1 bar = 105 Pa

• g/cc vs. kg/m3

• Per mol vs. per kg

Units!


Overview/Highlights


• 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)

Venus


Earth • Mostly N2,O2

• Moderate greenhouse

• Hydrological cycle & oceans

• Weathering buffer

• Moderate rotator

• Tilted (seasons)

• Hadley cell

• Milankovitch cycles

• Biological activity


Mars • Thin CO2 atmosphere

• Dust and polar caps important

• Massive climate change


• Tilted (seasons)

• Global dust storms

• Orbital forcing important (Milankovitch cycles)

• Mars Odyssey, Mars Express (ESA), Mars Exploration

Rovers, Mars Science Laboratory, MAVEN


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)


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


Titan • Moderate N2 atmosphere

• “Hydrological” cycle (methane)

• Subsurface replenishment


• Saturn tilted (seasons)

• Local clouds and storms

• Large atmospheric loss?

• Voyager, Cassini/Huygens


Thin Atmospheres


Exoplanets

Swain et al. 2008


1. How do planets form?


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


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)


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


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

Note sunspots

(roughly Earth-size)

1.4

mil

lion k

m


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

1cm chondrules

• They are also rich in volatile

elements

• Compositions are very

similar to Comet Halley,

also assumed to be ancient,

unaltered and volatile-rich


Meteorites vs. Photosphere

Basaltic Volcanism Terrestrial Planets, 1981

• 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


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):

Data from Lodders and Fegley, Planetary Scientist’s Companion, CUP, 1998

This is for all elements with relative abundances > 105 atoms.

Element H He C N O Ne Mg Si S Ar Fe

Log10 (No.

Atoms)

10.44 9.44 7.00 6.42 7.32 6.52 6.0 6.0 5.65 5.05 5.95

Condens.

Temp (K)

180 -- 78 120 -- -- 1340 1529 674 40 1337

• 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


Temperature and Condensation

Temperature profiles in a young (T

Tauri) stellar nebula, D’Alessio et al.,

A.J. 1998

Nebular conditions can be used to predict what components of

the solar nebula will be present as gases or solids:

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

Mid-plane

Photosphere

Earth

(~300K)

Saturn

(~50 K)

“Snow line” “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

“Snow line”


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?


• 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

Migration (hot Jupiters)

planet Gas disk (with

density waves)


Nice Model

Initial edge of

planetesimal

swarm

30 AU

Ejected planetesimals (Oort cloud)

48 AU 18 AU

J S U N

“Hot” population

Early in solar system

2:1 Neptune

resonance

J S U N

Neptune

stops at

original edge

Planetesimals transiently pushed

out by Neptune 2:1 resonance

“Hot” population

“Cold”

population

See Gomes, Icarus 2003 and Levison & Morbidelli Nature 2003

3:2 Neptune

resonance

(Pluto)

Present day


• 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

Planet Formation - Summary


2.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)

Where do atmospheres come from?


• 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)

Where do atmospheres go to?


3. Observational constraints

(see Taylor ch.3)


• “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

Radiometry (Spectroscopy)


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


• Atmospheric absorption

of light/radio waves

provides information on

composition, pressure

and temperature

Occultations

observer

observer

• Good for probing thin atmospheres (e.g. Pluto, Enceladus)

Hansen et

al. 2006


In situ sampling

LeBreton et al. Nature 2005

• Galileo probe (Jupiter)

• Huygens (Titan)

• Venera/Vega

probes/balloons

• Viking landers

• Cassini INMS

• 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)


4. Atmospheric structure


Temperature structure of

stratosphere in reality can

be more complicated

because of

photochemistry (e.g.

ozone)

Typical structure

T

z

troposphere

stratosphere

tropopause

Lower atmosphere

consists of a thick part

(troposphere) where

convection dominates,

and a thinner part above

(stratosphere) where

radiation dominates


Ideal Gas Equation

m

RTP

What is density of air at Earth’s surface?

What is the column mass of Earth’s atmosphere? (kg/m2)

P=pressure, =density, R=gas constant,

T=temperature (in K), m=molar mass (in kg)

m m

Venus (CO2) 0.04 Jupiter (H,He) 0.0022

Earth (N2,O2) 0.03 Saturn (H,He) 0.0021

Mars (CO2) 0.04 Uranus (H,He) 0.0026

Titan (N2) 0.03 Neptune (H,He) 0.0026


Atmospheric Structure (1)

• Atmosphere is hydrostatic:

• Gas law gives us:

• Combining these two (and neglecting latent heat):

)()( zgzdzdP

RT

gP

dz

dP m

Here R is the gas constant, m is the mass of one mole, and

RT/gm is 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

m

RTP


• 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!

Exobase and mean free path

l prmol

2

2 mol

mol

r

m

pl

What’s the mean free path

at the surface of the Earth?

rmol is typically 1 Angstrom=10-10 m


• Snow line

• Migration

• Troposphere/stratosphere

• Primary/secondary/tertiary atmosphere

• Emission/absorption

• Occultation

• Scale height

• Hydrostatic equilibrium

• Exobase

• Mean free path

Key concepts

First homework

due next

Monday!


End of lecture


An Artist’s Impression

The young Sun

gas/dust

nebula

solid planetesimals


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


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


Cartoon of Nebular Processes

• Scale height increases radially (why?)

• Temperatures decrease radially – consequence of lower

irradiation, and lower surface density and optical depth

leading to more efficient cooling

Polar jets

Stellar magnetic field

(sweeps innermost disk clear,

reduces stellar spin rate)

Disk cools by radiation

Dust grains Infalling

material

Nebula disk

(dust/gas)

Hot,

high

Cold,

low


Planetary Compositions • Which elements actually condense will depend on the

local nebular conditions (temperature)

• E.g. volatile species will only be stable beyond a “snow

line”. This is why the inner planets are rock-rich and the

outer planets gas- and ice-rich

• The compounds formed from the elements will be

determined by temperature (see next slide)

• The rates at which reactions occur are also governed by

temperature. In the outer solar system, reaction rates

may be so slow that the equilibrium condensation

compounds are not produced

• The mass of a planet determines the mass and

composition of its atmosphere

EART164: PLANETARY ATMOSPHERESfnimmo/eart164/Week1.pdf · F.Nimmo EART164 Spring 11 Expectations...

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