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Rotational Line BroadeningGray Chapter 18
Geometry and Doppler ShiftProfile as a Convolution
Rotational Broadening FunctionObserved Stellar Rotation
Other Profile Shaping Processes
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Doppler Shift of Surface Element
• Assume spherical star with rigid body rotation• Velocity at any point on visible hemisphere is
3
€
v = Ω × R
=
x^
y^
z^
0 Ωsini Ωcosi
x y z
=
zΩsini − yΩcosi( ) x^
+ xΩcosi( ) y^
+ −xΩsini( ) z^
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Doppler Shift of Surface Element
• z component corresponds to radial velocity• Defined as positive for motion directed away
from us (opposite of sense in diagram)• Radial velocity is
• Doppler shift is
4€
vR = xΩsin i
€
Δλ =λ0
cvR =
λ 0
cxΩsini( )
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Radial velocity depends only on x position.Largest at limb, x=R.
v = equatorial rotational velocity,v sin i = projected rotational velocity
€
ΔλL =λ 0
cRΩsini( ) =
λ 0
cv sin i( )
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Flux Profile
• Observed flux is (R/D)2 Fν where
• Angular element for surface element dA
• Projected element
• Expression for flux
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€
Fν = Iν cosθ dω∫
€
dω = dAR2
€
dx dy = dA cosθ
€
Fν =IνR2 dx dy∫∫
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Assumption: profile independent of position on visible hemisphere
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€
Fλ = H(λ − Δλ )Ic dx dy /R2∫∫
= H(λ − Δλ ) Icdy
Rd
Δλ
Δλ L
⎛
⎝ ⎜
⎞
⎠ ⎟
−y1
+y1
∫−R
+R
∫
y1 = R2 − x 2( )
1/ 2= R 1 −
Δλ
Δλ L
⎛
⎝ ⎜
⎞
⎠ ⎟
2 ⎡
⎣ ⎢ ⎢
⎤
⎦ ⎥ ⎥
1/ 2
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Express as a Convolution
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G(Δλ ) =1
Δλ L
Ic−y1
+y1
∫ dy /R
Ic cosθ dω∫for Δλ ≤ Δλ L
0 for Δλ > Δλ L
⎧
⎨ ⎪ ⎪
⎩ ⎪ ⎪
Fλ
Fc
= H(λ − Δλ ) G(−R
+R
∫ Δλ ) = H(λ − Δλ ) G(−∞
+∞
∫ Δλ )
= H(λ )∗G(λ )
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G(λ) for a Linear Limb Darkening Law
• Denominator of G
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IcIc
0 =1 −ε +ε cosθ
€
Ic cosθ dω =0
π / 2
∫∫ Ic cosθ sinθ dθ dφ μ = cosθ( )0
2π
∫
= −1
0
∫ Icμ dμ dφ = 2π Ic0
1
∫0
2π
∫ μ dμ
= 2π Ic0 (1−ε )μ +εμ 2 dμ =
0
1
∫ 2π Ic0 1−ε
2+ε
3
⎡ ⎣ ⎢
⎤ ⎦ ⎥
= π Ic0 1−
ε
3
⎡ ⎣ ⎢
⎤ ⎦ ⎥
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G(λ) for a Linear Limb Darkening Law
• Numerator of G
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€
IcIc
0 =1 −ε +ε cosθ
€
Icdy
R= 2Ic
0 2Ic0 1−ε( ) +ε cosθ[ ]
0
y1
∫ dy
R−y1
+y1
∫
= 2Ic0 1−ε( )
y1
R+ 2εIc
0 cosθdy
R0
y1
∫
= 2Ic0 1−ε( ) 1−
Δλ
Δλ L
⎛
⎝ ⎜
⎞
⎠ ⎟
2 ⎡
⎣ ⎢ ⎢
⎤
⎦ ⎥ ⎥
1/ 2
+ 2εIc0 1
R2 R2 − x 2 − y 2 dy0
y1
∫
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G(λ) for a Linear Limb Darkening Law
• Analytical solution for second term in numerator
• Second term is
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€
IcIc
0 =1 −ε +ε cosθ
€
2εIc0 1
2
1
R2 y(R2 − x 2 − y 2)1/ 2 + (R2 − x 2)arcsiny
R2 − x 2
⎡
⎣ ⎢
⎤
⎦ ⎥0
y1
=εIc
0
R2 (R2 − x 2)π
2
⎡ ⎣ ⎢
⎤ ⎦ ⎥ y1 = R2 − x 2
( )
=π
2εIc
0 1−Δλ
Δλ L
⎛
⎝ ⎜
⎞
⎠ ⎟
2 ⎡
⎣ ⎢ ⎢
⎤
⎦ ⎥ ⎥
€
A2 − y 2( )
1/ 2dy =
1
2y A2 − y 2( )
1/ 2+ A2 arcsin
y
A
⎡
⎣ ⎢
⎤
⎦ ⎥∫
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G(λ) for a Linear Limb Darkening Law
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€
IcIc
0 =1 −ε +ε cosθ
€
∴G Δλ( ) =
2 1 −ε( ) 1−Δλ
Δλ L
⎛
⎝ ⎜
⎞
⎠ ⎟
2 ⎡
⎣ ⎢ ⎢
⎤
⎦ ⎥ ⎥
1/ 2
+π
2ε 1 −
Δλ
Δλ L
⎛
⎝ ⎜
⎞
⎠ ⎟
2 ⎡
⎣ ⎢ ⎢
⎤
⎦ ⎥ ⎥
πΔλ L 1 −ε
3
⎛
⎝ ⎜
⎞
⎠ ⎟
= c1 1 −Δλ
Δλ L
⎛
⎝ ⎜
⎞
⎠ ⎟
2 ⎡
⎣ ⎢ ⎢
⎤
⎦ ⎥ ⎥
1/ 2
+ c2 1 −Δλ
Δλ L
⎛
⎝ ⎜
⎞
⎠ ⎟
2 ⎡
⎣ ⎢ ⎢
⎤
⎦ ⎥ ⎥
ellipse
parabola
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Grey atmosphere case: ε = 0.6
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v sin i = 20 km s-1
v sin i = 4.6 km s-1
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Measurement of Rotation
• Use intrinsically narrow lines• Possible to calibrate half width with v sin i, but
this will become invalid in very fast rotators that become oblate and gravity darkened
• Gray shows that G(Δλ) has a distinctive appearance in the Fourier domain, so that zeros of FT are related to v sin i
• Rotation period can be determined for stars with spots and/or active chromospheres by measuring transit times
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Rotation in Main Sequence Stars
• massive stars rotate quickly with rapid decline in F-stars(convection begins)
• low mass stars have early, rapid spin down, followed by weak breaking due to magnetism and winds (gyrochronology)
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L = M R v
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Angular Momentum – Mass Relation• Equilibrium with gravity = centripetal acceleration
• Angular momentum for uniform density
• In terms of angular speed and density
• Density varies slowly along main sequence19
€
GM
R2 =v 2
R⇒
GM
R3 =v 2
R2 =ω 2 ⇒ ω ∝ ρ
€
L = MRv = MR2ω L = Iω = k 2MR2ω( )
€
R3 =GM
ω 2 ⇒ R =GM
ω 2
⎛
⎝ ⎜
⎞
⎠ ⎟
1/ 3
L ∝ M M 2 / 3ω −4 / 3ω = M 5 / 3ω −1/ 3 ∝ M 5 / 3ρ −1/ 6
€
L ∝ M 5 / 3
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Rotation in Evolved Stars
• conserve angular momentum, so as R increases, v decreases
• Magnetic breaking continues (as long as magnetic field exists)
• Tides in close binary systems lead to synchronous rotation
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Fastest Rotators• Critical rotation
• Closest to critical in the B stars where we find Be stars (with disks)
• Spun up by Roche lobe overflow from former mass donor in some cases (ϕ Persei)
21
€
vcrit =GM
R= 437
M /Msun
R /Rsun
⎛
⎝ ⎜
⎞
⎠ ⎟
1/ 2
km s−1
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Other Processes That Shape Lines• Macroturbulence and granulation
http://astro.uwo.ca/~dfgray/Granulation.html
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Star Spots
24
Vogt & Penrod 1983, ApJ, 275, 661
HR 3831Kochukhov et al. 2004, A&A, 424, 935http://www.astro.uu.se/~oleg/research.html
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Stellar Pulsationhttp://staff.not.iac.es/~jht/science/
25Vogt & Penrod 1983, ApJ, 275, 661
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Stellar Winds
• Atoms scatter starlight to create P Cygni shaped profiles
• Observed in stars with strong winds (O stars, supergiants)
• UV resonance lines (ground state transitions)
26http://www.daviddarling.info/encyclopedia/P/P_Cygni_profile.html
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FUSE spectra (Walborn et al. 2002, ApJS, 141,443)
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To really know a star ... get a spectrum
• “If a picture is worth a thousand words, then a spectrum is worth a thousand pictures.”(Prof. Ed Jenkins, Princeton University)
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