Stellar winds - Universität zu Kölnszecsi/Publications/Talks/... · 2019. 3. 11. · Basic ideas,...
Transcript of Stellar winds - Universität zu Kölnszecsi/Publications/Talks/... · 2019. 3. 11. · Basic ideas,...
Stellar winds
Dorottya Szécsi
PhD Meeting of Stellar students12th November 2013
Outline
• Basic ideas, M and v∞• Radiation field of stars with different Teff• Physical processes in the atmosphere• Line driven winds• The bi-stability jump• Vink et al. 2000• The famous P Cygni profile• WR stars, cool stars• What is actually implemented in BEC
Basic ideas, M and v∞Stellar wind: continuous outflow; stars emit radiation AND particles=’wind’
Sun:• flows driven by gas pressure gradients (1958)
• Alfvén wave driven winds (1965)
• magnetic rotator theory (1967)
• radiation driven winds (1970)
Two main parameters to derive from wind theories (and spectra): M and v∞• kinetic energy that the wind deposits into the ISM per unit time (→momentum
and energy transfer):
dEkindt
=12Mv2∞ (1)
• velocity law: v(r) distribution of the velocity of the wind with radial distance –approximation:
v(r) ' v0 + (v∞ − v0)(1− R∗
r
)β(2)
v0: initial velocity at the photosphereβ: steepness of the law
(v∞ = v(r→∞))
Radiation field of stars with different Teff
B(ν,Teff ) =2hc2
ν3
ehν
kBTeff − 1
(3)
0
5e-06
1e-05
1.5e-05
2e-05
2.5e-05
0 2e+14 4e+14 6e+14 8e+14 1e+15
Sp
ectr
al ra
dia
nce
[e
rg/c
m2
]
Frequency [1/s]
(2*h/c**2)*(x**3)/(exp((h/k)*x/3000)-1)(2*h/c**2)*(x**3)/(exp((h/k)*x/4000)-1)(2*h/c**2)*(x**3)/(exp((h/k)*x/5000)-1)
0
0.0001
0.0002
0.0003
0.0004
0.0005
0.0006
0.0007
0 5e+14 1e+15 1.5e+15 2e+15 2.5e+15 3e+15 3.5e+15 4e+15
Sp
ectr
al ra
dia
nce
[e
rg/s
/Hz/c
m2
/sr]
Frequency [Hz]
300040005000
1000015000
0
0.005
0.01
0.015
0.02
0.025
0.03
0.035
0 2e+15 4e+15 6e+15 8e+15 1e+16 1.2e+16 1.4e+16
Sp
ectr
al ra
dia
nce
[e
rg/s
/Hz/c
m2
/sr]
Frequency [Hz]
300040005000
10000150003500055000
0
0.001
0.002
0.003
0.004
0.005
0.006
0.007
0.008
0 2 4 6 8 10 12 14 16 18 20
Flu
x [
erg
/cm
2/s
/Hz/s
r]
Frequency [1015
Hz]
Teff = 55000 K log L/L⊙ = 6.1 M = 80.3 M⊙
FLyFW
= 7.6e+05 L⊙FLy
BB = 7.4e+05 L⊙
FHeIFW
= 3.2e+05 L⊙
FHeIBB
= 2.6e+05 L⊙
Lyman continuum
HeI co
ntin
uum
FASTWIND SPECTRUM
BLACK BODY
Radiation field of stars with different Teff
B(ν,Teff ) =2hc2
ν3
ehν
kBTeff − 1
(3)
0
5e-06
1e-05
1.5e-05
2e-05
2.5e-05
0 2e+14 4e+14 6e+14 8e+14 1e+15
Sp
ectr
al ra
dia
nce
[e
rg/c
m2
]
Frequency [1/s]
(2*h/c**2)*(x**3)/(exp((h/k)*x/3000)-1)(2*h/c**2)*(x**3)/(exp((h/k)*x/4000)-1)(2*h/c**2)*(x**3)/(exp((h/k)*x/5000)-1)
0
0.0001
0.0002
0.0003
0.0004
0.0005
0.0006
0.0007
0 5e+14 1e+15 1.5e+15 2e+15 2.5e+15 3e+15 3.5e+15 4e+15
Sp
ectr
al ra
dia
nce
[e
rg/s
/Hz/c
m2
/sr]
Frequency [Hz]
300040005000
1000015000
0
0.005
0.01
0.015
0.02
0.025
0.03
0.035
0 2e+15 4e+15 6e+15 8e+15 1e+16 1.2e+16 1.4e+16
Sp
ectr
al ra
dia
nce
[e
rg/s
/Hz/c
m2
/sr]
Frequency [Hz]
300040005000
10000150003500055000
0
0.001
0.002
0.003
0.004
0.005
0.006
0.007
0.008
0 2 4 6 8 10 12 14 16 18 20
Flu
x [erg
/cm
2/s
/Hz/s
r]
Frequency [1015
Hz]
Teff = 55000 K log L/L⊙ = 6.1 M = 80.3 M⊙
FLyFW
= 7.6e+05 L⊙FLy
BB = 7.4e+05 L⊙
FHeIFW
= 3.2e+05 L⊙
FHeIBB
= 2.6e+05 L⊙
Lyman continuum
HeI co
ntin
uum
FASTWIND SPECTRUM
BLACK BODY
Physical processes in the atmosphere
• line scattering (ground state: resonance scattering→ resonance lines)
• recombination→ emission (e.g. Hα emission)
• collision→ emission
• masering
Spectral lines: photospheric lines , wind lines (Doppler shifted, absorption+emission)
Resonance lines:• most sensitive indicators of mass loss from hot stars
• scattering:
• photon emitted by the photosphere is absorbed by an atom→• photo-excitation of a ground-state electron→• new photon is emitted, but in another direction
• observable, because: abundant elements; oscillator strength of atomic resonancetransitions is large
• examples: CIV, NV, SiIV (O-B stars), CII (B-A stars), MgII (B-M stars) – in UV!
Line driven winds
Luminous hot stars: radiation is emitted mainly in UV+
many absorption lines here (= elements in the atmosphere with potential transitions)+
Doppler effect: otherwise all radiation would be absorbed in the inner atmosphere...but outer parts are moving away (redshifted) :)
Momentum and energy transfer from radiation to the gas:• a bit complicated to calculate all Doppler-shifted transfers of all wavelength...
• and gravity may play a role... and radiation field may be not perfect Black Body...
• and not every gas-component scatters – they must bring everything with them(Coulomb coupling) to reach a steady outflow of plasma...
Sobolev approximation:treats the ’line’ (=small frequency range of the scattering) as a δ-function and theoptical depth as a step-function.
CAK theory (Castor, Abbott, Klein):a realistic estimate of the radiative force due to lines: computation of the degree ofionization and excitation of large number of energy levels for many different elements
The bi-stability jump
Effective escape velocity (at the stellar surface):
vesc =
√2(1− Γe)
GMR
(4)
The ratio of v∞/vesc depends on Teff :
(Markova & Puls 2008)
21 000 K . Teff :C, N, O highly ionized lines in theLymann continuum
10 000 K . Teff :metal lines in the Lymann andBalmer continuum
Teff . 10 000 K:continuum driven wind (?) dust,molecules (?)
Vink et al. 2000
Model atmospheres of stars with different L∗, M∗, Teff , v∞/vesc• covering the OB-range of the HRD
• multiple scattering!
• method: total loss of radiative energy is linked to the gain of momentum of thematerial outflow
• model atmosphere→ radiative acceleration and M are calculated
• M as a function of the basic parameters are fitted (= "mass loss recipe"), e.g. for12 500 K < Teff < 22 500 K:
logM = −6.69 + 2.21logL∗
105 − 1.34logM∗30
+ 1.07logTeff
20000− 1.6log
v∞/vesc2.0
(5)
• Vink et al. 2001: metallicity dependence (+0.85logZ)
• behaviour around the bi-stability jumps
• the calculated M match the observations
The famous P Cygni profile
violet shifted absorptionlines are observable evenat low column density;
also a symmetricemission component is
observable coming fromthe ’halo’ region of thestar (recombination)
M, v∞ and β canbe accurately
derived
WR stars, cool stars
Hα line (λ=6562Å→ optical spectra of OB supergiants):• emission line, high M needed
• widely used to derive parameters
WR stars (Hamann et al. 1995)
• high density strong winds→ spectrum is dominated by emission lines
• emission comes from recombination ∼ ρ2 ⇒ emission comes from lower layers
• problem with radius determination: normally, R∗ = R(τ = 2/3) at thephotosphere (where the continuum comes from); but WR wind is optically thick!→ the ’surface’ of the star is somewhere in the wind
• measured radius is different in different wavelength bands (τ = τ(λ)) :)
• mass loss rate for hot stars: derive M from radio continuum measurements + v∞from P Cygnu profile in UV
Cool stars (Nieuwenhuizen & de Jager 1990)
• molecular emission lines, e.g. CO, SiO, OH
• additionally: dust!
• dusty winds emit radiation in IR and mm bands, but the energy distribution ofdust is typical
What is actually implemented in BEC
• m.dat: MTU=0, MTU=20, MTU=51 etc.
• BEC/CODE/dev/maslos.f
• "N": Nieuwenhuizen & de Jager (1990)
• "V": Vink et al. (2000-2001)
• "T": interpolated
• "W": 0.1 × Hamann et al. (1995)
• "O": ad hoc expression by Sung-Chul(?), see Yoon et al. (2006)
5.8
5.9
6
6.1
6.2
6.3
6.4
6.5
4.7 4.8 4.9 5 5.1 5.2 5.3 5.4 5.5
77-350
77-400
0
"N"
"V"
"T"
"W"
"O"
Thank you for your
attention!