modeling the structure of hot star disks
DESCRIPTION
Modeling the Structure of Hot Star Disks. Jon E. Bjorkman Ritter Observatory. Systems with Disks. Infall + Rotation Young Stellar Objects (T Tauri, Herbig Ae/Be) Mass Transfer Binaries Active Galactic Nuclei (Black Hole Accretion Disks) Outflow + Rotation - PowerPoint PPT PresentationTRANSCRIPT
Modeling the Structure of Hot Star Disks
Jon E. Bjorkman
Ritter Observatory
Systems with Disks
• Infall + Rotation– Young Stellar Objects (T Tauri, Herbig Ae/Be)– Mass Transfer Binaries– Active Galactic Nuclei (Black Hole Accretion Disks)
• Outflow + Rotation– AGBs (bipolar planetary nebulae)– LBVs (e.g., Eta Carinae)– Oe/Be, B[e]
• Rapidly rotating (Vrot = 350 km s-1)• Hot stars (T = 20000K)• Ideal laboratory for studying disks
General Wind Considerations
• Radial Momentum Equation
• Radial Motion
• Be disk line profiles– Widths and symmetry =>
v
dv
dr= -
GM
r2+
vf2
r+ grad
vr ~
Vesc (radiation-driven)
= a (Keplerian)
ÏÌÔÔ
ÓÔÔ
vr £ a (Dachs, Hanuschik, …)
Disk probably is Keplerian
General Wind Considerations
• Azimuthal Motion
vf µ
r - 1 (vr ? vf ; angular momentum-conserving)
r (magnetically dominated; solid body rotation)
r - 1/ 2 (vr = vf ; Keplerian)
Ï
Ì
ÔÔÔÔÔ
Ó
ÔÔÔÔÔ
Rotating (in/out) Flows
• Fluid Equations (cylindrical coords: )
• Equation of State
€
ϖ,z,φ
€
P = a2ρ
€
a =kT
μmH
(isothermal sound speed)
Keplerian Disks
• Fluid Equations
• Vertical scale height
(Keplerian orbit)
(Scale height)
(Hydrostatic)
€
(vϖ << vφ;v z = 0)
€
fϖ
€
fz
€
Fgrav
€
T = 15000K
€
H /ϖ = 0.04
€
Δθ =2.5°
Disk Variability
• Viscosity
• Viscous Diffusion Timescale
– too large, unless a ~ 0.1–1
n = aaH
tn = v 2 / n
=Vcrit
aa2v R
= 20yr 0.01
a
Ê
ËÁÁÁ
ˆ
¯˜̃˜̃
v
R
(eddy viscosity)
Viscous Decretion -Disks
vf =Vcrit R / v
vv =&M
2pv S
S =&MVcritR
1/ 2
3paa2v 3/ 2
Rmax
v- 1
È
Î
ÍÍÍ
˘
˚
˙˙˙
r =S
2pHe- 0.5(z/H )2
H = (a / vf )v
(surface density)
(scale height)
(Keplerian orbit)
(hydrostatic)
(continuity eq.)
Power Law Approximations
• Keplerian Decretion Disk
• Flaring
b = 98
a = 198
(flat passive disk; T µ r - 3/ 4)
b = 54
a = 114
(flared passive disk; T µ r - 1/ 2)
b = 32
a = 72
(isothermal disk; T = const)
r = r0(R* / v )a exp -z
H(v )
Ê
ËÁÁÁ
ˆ
¯˜̃˜̃
2È
Î
ÍÍÍ
˘
˚
˙˙˙
H = H0(v / R*)b
Isothermal Keplerian Disk Density
= 1.5
= 3.5
€
ρ(ϖ ) = ρ0 (R* / ϖ )α exp −12
zH (ϖ ) ⎛ ⎝ ⎜
⎞ ⎠ ⎟2 ⎡
⎣ ⎢ ⎢
⎤
⎦ ⎥ ⎥
H (ϖ ) = H0 (ϖ / R* )β
Monte Carlo Radiation Transfer
• Divide stellar luminosity into equal energy packets
• Pick random starting location and direction• Transport packet to random interaction location
• Randomly scatter or absorb photon packet• When photon escapes, place in observation bin
(frequency and direction)
€
Eγ = L Δt / Nγ
€
τ =−lnξ (ξ is a random number)
REPEAT 106-109 times
MC Radiative Equilibrium• Sum energy absorbed by each cell• Radiative equilibrium gives temperature
• When photon is absorbed, reemit at new frequency, depending on T
• Problem: Don’t know T a priori• Solution: Change T each time a photon is
absorbed and correct previous frequency distribution
avoids iteration
€
Eabs = Eemit
nabsEγ = 4πmiκ PB(Ti )
Temperature Correction
Bjorkman & Wood 2001
Frequency Distribution:
€
dP
dν= jν (T + ΔT ) − jν (T )
=κ ν ΔTdBν
dT
Model of B[e] Star
Bjorkman 1998
Disk Temperature
Bjorkman 1998
B[e] SED
Bjorkman 1998
T Tauri Envelope Absorption
Disk Temperature
Snow LineWater Ice
Methane Ice
Effect of Disk on Temperature• Inner edge of disk
– heats up to optically thin radiative equilibrium temperature
• At large radii– outer disk is shielded by inner disk– temperatures lowered at disk mid-plane
• Permits dust formation in outer disk
• Requires a different opacity source at smaller radii
NLTE Monte Carlo RT• Gas opacity depends on:
– temperature– degree of ionization – level populations
• During Monte Carlo simulation:– sample radiative rates
• Radiative Equilibrium– Whenever photon is absorbed, re-emit it
• After Monte Carlo simulation:– solve rate equations– update level populations and gas temperature– update disk density (solve hydrostatic equilibrium)
determined by radiation field
Disk Temperature
Carciofi & Bjorkman 2004
Disk Density
Carciofi & Bjorkman 2004
NLTE Level Populations
Carciofi & Bjorkman 2004
SED and Polarization
Carciofi & Bjorkman 2004
IR Excess
Carciofi & Bjorkman 2004
LTE Line-Blanketed Polarization
Observed Observed
MC Simulation MC Simulation
Flux Polarization
Bjorkman & Carciofi 2003
Inner Disk:• NLTE Hydrogen • Flared Keplerian• h0 = 0.07, = 1.5
• R* < r < Rdust
Outer Disk:• Dust • Flared Keplerian• h0 = 0.017, = 1.25
• Rdust < r < 10000 R*
HAeBe Model
Summary• Viscous Timescale:
– 20 years ( = 0.01)– probably a bit too long (but may be larger)
• NLTE Modeling of Keplerian Disk– Fully self-consistent 3-D model
• determines radiative equilibrium temperature• vertical hydrostatic equilibrium• steady state disk surface density• Single parameter: (and inclination angle i)
– Reproduces Polarization and SED– Temperature
• Inner disk: falls as r -3/4 (like thin blackbody)• Outer disk: isothermal
&M
&M = (a / 0.01)¥ 10- 11Me yr- 1
Acknowledgments
• Rotating winds and bipolar nebulae– NASA NAGW-3248
• Ionization and temperature structure– NSF AST-9819928– NSF AST-0307686
• Geometry and evolution of low mass star formation– NASA NAG5-8794
• Collaborators: A. Carciofi, K.Wood, B.Whitney, K. Bjorkman, J.Cassinelli, A.Frank, M.Wolff
• UT Students: B. Abbott, I. Mihaylov, J. Thomas• REU Students: A. Moorhead, A. Gault
Be Star H Profile
Carciofi and Bjorkman 2003
i = 82º
Polarization vs IR Excess
Coté & Waters 1987
P ~ sin2i
Edge-on
Pole-on
Gault, Bjorkman & Bjorkman 2002
MC Polarization vs IR Excess
Disk Orientation: Inclination
QuickTime™ and aTIFF (Uncompressed) decompressor
are needed to see this picture.
Quirrenbach et al. 1996
Interferometric sin2i
Pol
arim
etri
c s
in2 i