atmospheres of cool stars
DESCRIPTION
Atmospheres of Cool Stars. Radiative Equilibrium Models Extended Atmospheres Heating Theories. Radiative Equilibrium Models. Gustafson et al. (2005): MARCS code difficult because of UV “line haze” (millions of b-b transitions of Fe I in 300-400 nm range, and Fe II in 200-300 nm range) - PowerPoint PPT PresentationTRANSCRIPT
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Atmospheres of Cool Stars
Radiative Equilibrium ModelsExtended Atmospheres
Heating Theories
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Radiative Equilibrium Models
• Gustafson et al. (2005): MARCS codedifficult because of UV “line haze”(millions of b-b transitions of Fe I in 300-400 nm range, and Fe II in 200-300 nm range)
• Convection important at depth
• Metallicity and line blanketing causes surface cooling and back warming
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Dwarfs Giants
• Solid line = Solar abundances[Fe/H]=0
• Dashed line =Metal poor[Fe/H]=-2
• Dot-dashed =Kurucz LTE-RE model
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Semi-empirical Models Based on Observations of Iλ(μ=1,τ=1)
• Solar spectrum shows non-thermal components at very long and short wavelengths that indicate importance of other energy transport mechanisms
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Major b-b and b-f transitions for solar opacity changes
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• Determine central specific intensity across spectrum
• Get brightness temperaturefrom Planck curve for Iλ
• From opacity get optical depth on standard depth scale
• T(h) for h=height
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Optical:photosphere
EUV: higher
X-ray: higher yet
Reality: Structured and Heated
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Extended Atmospheres
• Photosphere• Chromosphere• Transition region• Corona• Wind
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Corona
• Observed during solar eclipses or by coronagraph (electron scattering in optical)
• Nearly symmetric at sunspot maximum, equatorially elongated at sunspot minimum
• Structure seen in X-rays (no X-ray emission from cooler, lower layers)
• Coronal lines identified by Grotrian, Edlén (1939): Fe XIV 5303, Fe X 6374, Ca XV 5694
• High ionization level and X-rays indicate T~106 K
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X-ray image of Sun’s hot coronal gas
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Chromosphere
• Named for bright colors (“flash spectrum”) observed just before and after total eclipse
• H Balmer, Fe II, Cr II, Si II lines present: indicates T = 6000 – 10000 K
• Lines from chromosphere appear in UV (em. for λ<1700 Å; absorption for λ>1700 Å)
• Large continuous opacity in UV, but lines have even higher opacity: appear in emission when temperature increases with height
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Transition Region
• Seen in high energy transitions which generally require large energies (usually in lines with λ<2000 Å)
• Examples in solar spectrum:Si IV 1400, C IV 1550 (resonance or ground state transitions)
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Stellar Observations
• Chromospheric and transition region lines seen in UV spectra of many F, G, K-type stars (International Ultraviolet Explorer)
• O I 1304, C I 1657, Mg II 2800• Ca II 3968, 3933 (H, K) lines observed as
emission in center of broad absorption (related to sunspot number in Sun; useful for starspots and rotation in other stars)
• Emission declines with age (~rotation)
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Chromospheres in H-R Diagram• Emission lines
appear in stars found cooler than Cepheid instability strip
• Red edge of strip formed by onset of significant convection that dampens pulsations
• Suggests heating is related to mechanical motions in convection
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Coronae in H-R Diagram• Upper luminosity
limit for stars with transition region lines and X-ray coronal emission
• Heating not effective in supergiants (but mass loss seen)
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Theory of Atmospheric Heating
• Increase in temperature cannot be due to radiative or thermal processes
• Need heating by mechanical or magneto-electrical processes
T T T ateff e ff4 4 43
4
2
3
1
20
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Acoustic Heating
• Large turbulent velocities in solar granulation are sources of acoustic (sound) waves
• Lightman (1951), Proudman (1952) show that energy flux associated with waves is
where v = turbulent velocity and cs = speed of sound
F v cac s 8 5/
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Acoustic Heating
• Acoustic waves travel upwards with energy flux = (energy density) x (propagation speed)= ½ ρ v2 cs
• If they do not lose energy, then speed must increase as density decreases→ form shock waves that transfer energy into the surrounding gas
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Wave Heating
• In presence of magnetic fields, sound and shock waves are modified into magneto-hydrodynamic (MHD) waves of different kinds
• Damping (energy loss) of acoustic modes depends on wave period:ex. 5 minute oscillations of Sun in chromosphere with T = 10000 K yields a damping length of λ = 1500 km
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Wave Heating
• Change in shock flux with height is
• Energy deposited (dissipated into heat) at height h is
where
dF
dhFmechm ech
1
F h F emech mh
0/
F hmech /
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Wave Heating
• Energy also injected by Alfvén waves (through Joule heating caused by current through a resistive medium)
• Observations show spatial correlation between sites of enhanced chromospheric emission and magnetic flux tube structures emerging from surface: magnetic processes cause much of energy dissipation
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Balance Heating and Cooling
• Energy loss by radiation through H b-f recombination in Lyman continuum (λ < 912 Å) and collisional excitation of H
• In chromosphere, H mainly ionized, primary source of electrons
• for H recombination for H collisional excitation
• Similar relations exist for other ions
E n f T
n n f Trad e
e H H
2 ( )
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Radiative Loss Function
• Below T = 15000 K, f(T) is a steep function of T because of increasing H ionization
• Above T = 15000 K, H mostly ionized so it no longer contributes much to cooling
• He ionization becomes a cooling source for T > 20000 K:
• Above T=105 K, most abundant species are totally ionized → slow decline in f(T)
f TT
10
300002 022
, .
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Radiative Loss Function
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Energy Balance
• In lower transition region (hot) Pg = 2 Pe
Electron density:
Radiation loss rate:
(almost independent of T since f(T)~T 2.0)• Set T(h) by Einput = Erad
• Suppose Einput = Fmech(h) / λ
nP
kTe
g2
E f TP
kTrad
g
2
22
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(1) Einput = constantT(h) increases with h
Increasing height in outer atmosphere
Each line down corresponds to a 12% drop in Pg or a 26% drop in Pg
2
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(2) Einput declines slowly with hT(h) still increases with h
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(3) Einput declines quickly with h T(h) may not increase with h
No T increase for damping length λ and pressure scale height H ifλ < H/2
H is large in supergiants so heating in outer atmosphere does not occur.
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Temperature Relation for Dwarfs• Suppose λ >> H in lower transition region
so that Fmech(h) / λ ≈ constant
f T n B T n BTP
kconst
d T
dh
d P
dh
d T
dh
d P
dh
d P
dh HgP
g m
kT
d T
dh
g m
kT
A
T
dT
dhA
e e
g
g
g g H
H
2 2 2
2
2
2 2 0
2
2
1
2
2
( ) .
ln ln
ln ln,
ln
ln
Constant temperature gradient
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Heating in the Outer Layers
• T>105 K, rad. losses cannot match heating
• T increases until loss by conductive flux downwards takes over (+ wind, rad. loss)
• Conductive flux (from hot to cool regions by faster speeds of hotter particles)
• Find T(h) from (η ~ 10-6 c.g.s.)
| | /F TdT
dhF h F F hc m ech wind rad 5 2
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