Download cranmer_nessc_oct2008

Survey
yes no Was this document useful for you?
   Thank you for your participation!

* Your assessment is very important for improving the work of artificial intelligence, which forms the content of this project

Document related concepts

Astronomical spectroscopy wikipedia, lookup

Formation and evolution of the Solar System wikipedia, lookup

CoRoT wikipedia, lookup

Star formation wikipedia, lookup

Standard solar model wikipedia, lookup

Corona wikipedia, lookup

Advanced Composition Explorer wikipedia, lookup

Stellar kinematics wikipedia, lookup

Theoretical astronomy wikipedia, lookup

R136a1 wikipedia, lookup

Cygnus X-1 wikipedia, lookup

Transcript
Solar/Stellar Winds:
An Overview
(Untapped gold mines for
space physicists?)
Steven R. Cranmer
Harvard-Smithsonian Center for Astrophysics
Outline
The Stellar Zoo:
• The Sun
• Young low-mass stars (T Tauri, FU Ori, etc.)
• Evolved low-mass stars (giants, supergiants, Miras)
• High-mass stars (Wolf-Rayet, O, B, A, LBVs)
Universal Physical Processes:
• Pulsations / waves / shocks / turbulence
• Magnetic fields: dynamos, reconnection between open & closed regions
• “Radiation hydrodynamics:” forcing, heating/cooling, kinetic effects?
Solar/Stellar Winds: An Overview
S. R. Cranmer, NESSC
October 27, 2008
Stellar winds on the HR Diagram
106
no coronae?
I
"cool"
dense
radiatively
driven winds
104
(slow?)
winds
"warm"
hybrid
winds
102
III
Be stars
"hot"
solar-type
winds
V
1
10–2
Sun

flare
stars
O
30,000
Solar/Stellar Winds: An Overview
B
A
10,000
F G
6,000
K
M
3,000
S. R. Cranmer, NESSC
October 27, 2008
Driving a stellar wind
• Gravity must be counteracted above the photosphere (not below)
by some continuously operating outward force . . .
 Gas pressure: needs T ~ 106 K (“coronal heating”)
F = ma
 Radiation pressure: possibly important when L* > 100 L
• ion opacity? (Teff ~> 15,000 K)
• free electron (Thomson) opacity? (goes as 1/r2 ; needs to be supplemented)
• dust opacity? (Teff <~ 3,500 K)
 Wave pressure: can produce outward acceleration in a time-averaged sense
 Magnetic buoyancy: plasmoids can be “pinched” like melon seeds and
carry along some of the surrounding material . . .
Solar/Stellar Winds: An Overview
S. R. Cranmer, NESSC
October 27, 2008
The solar wind: very brief history
• Mariner 2 (1962): first direct confirmation of continuous supersonic solar wind,
validating Parker’s (1958) model of a gas-pressure driven wind.
• Helios probed in to 0.3 AU, Voyager continues past 100+ AU.
• Ulysses (1990s) left the ecliptic; provided 3D view of the
wind’s connection to the Sun’s magnetic geometry.
• SOHO gave us new views of “source regions” of solar wind
and the physical processes that accelerate it . . .
Solar/Stellar Winds: An Overview
S. R. Cranmer, NESSC
October 27, 2008
The coronal heating problem
• We still don’t understand the physical processes responsible for heating up the
coronal plasma.
A lot (not all!) of the heating occurs in a narrow “shell.”
• Most suggested ideas involve 3 general steps:
1. Churning convective motions that tangle up
magnetic fields on the surface.
2. Energy is “stored” (above the photosphere) in
the magnetic field.
3. Energy is released as heat, either via particleparticle or wave-particle “collisions.”
Heating
Solar wind acceleration
Solar/Stellar Winds: An Overview
S. R. Cranmer, NESSC
October 27, 2008
The solar wind acceleration debate
• What determines how much energy and
momentum goes into the solar wind?
Waves & turbulence input from below?
vs.
Reconnection & mass input from loops?
Solar/Stellar Winds: An Overview
S. R. Cranmer, NESSC
October 27, 2008
The solar wind acceleration debate
• What determines how much energy and
momentum goes into the solar wind?
Waves & turbulence input from below?
vs.
Reconnection & mass input from loops?
• Cranmer et al. (2007) explored
the wave/turbulence paradigm
with self-consistent 1D models
of individual open flux tubes.
• Boundary conditions imposed
only at the photosphere (no
arbitrary “heating functions”).
• Wind acceleration determined by a combination of
magnetic flux-tube geometry, gradual Alfvén-wave
reflection, and outward wave pressure.
Solar/Stellar Winds: An Overview
S. R. Cranmer, NESSC
October 27, 2008
Stellar winds on the HR Diagram
106
no coronae?
I
"cool"
dense
radiatively
driven winds
104
(slow?)
winds
"warm"
hybrid
winds
102
III
Be stars
"hot"
solar-type
winds
V
1
10–2
Sun

flare
stars
O
30,000
Solar/Stellar Winds: An Overview
B
A
10,000
F G
6,000
K
M
3,000
S. R. Cranmer, NESSC
October 27, 2008
Pre-Main-Sequence accretion phases
Feigelson &
Montmerle
(1999)
M. Burton
(UNSW)
Solar/Stellar Winds: An Overview
S. R. Cranmer, NESSC
October 27, 2008
T Tauri stars: accreting young Suns
• T Tauri stars exhibit signatures of disk accretion (outer parts), “magnetospheric
accretion streams” & X-ray corona (inner parts), and various (polar?) outflows.
• Nearly every observational diagnostic varies in time, sometimes with stellar
rotation, but often more irregularly.
(Romanova et
al. 2007)
(Matt & Pudritz 2005, 2008)
Solar/Stellar Winds: An Overview
S. R. Cranmer, NESSC
October 27, 2008
Accretion-driven T Tauri winds
• Recent work has extended the solar wave/turbulence models to outer atmospheres
of young, solar-type stars (Cranmer 2008, arXiv:0808.2250).
• The impact of inhomogeneous “clumps” on the stellar surface generates MHD
waves that propagate horizontally (like solar Moreton/EIT waves?).
• These “extra” waves input orders of magnitude more energy into a turbulent MHD
cascade, and can give rise to stellar winds with dM/dt up to 106 times solar!
Solar/Stellar Winds: An Overview
S. R. Cranmer, NESSC
October 27, 2008
Stellar winds on the HR Diagram
106
no coronae?
I
"cool"
dense
radiatively
driven winds
104
(slow?)
winds
"warm"
hybrid
winds
102
III
Be stars
"hot"
solar-type
winds
V
1
10–2
Sun

flare
stars
O
30,000
Solar/Stellar Winds: An Overview
B
A
10,000
F G
6,000
K
M
3,000
S. R. Cranmer, NESSC
October 27, 2008
Cool-star mass loss rates
Solar/Stellar Winds: An Overview
S. R. Cranmer, NESSC
October 27, 2008
Cool-star dimensional analysis . . .
• Stellar wind power:
• Reimers (1975, 1977) proposed a semi-empirical scaling:
• Schröder & Cuntz (2005) investigated an explanation via convective turbulence
generating atmospheric waves . . .
• Funny things happen during rapid evolutionary stages!
(e.g., Willson 2000, Ann. Rev.)
Solar/Stellar Winds: An Overview
S. R. Cranmer, NESSC
October 27, 2008
Cool-star mass loss rates
Schröder & Cuntz (2005)
scaling for lum. classes I, III, V
Solar/Stellar Winds: An Overview
S. R. Cranmer, NESSC
October 27, 2008
How can massive winds be “cold?”
• The extended solar corona is so low-density, the conservation of internal energy is
essentially a balance between local heating, downward conduction, and upward
adiabatic losses.
• When the outer atmosphere becomes massive enough, though, radiative cooling
[~ρ2 Λ(T)] becomes more efficient throughout the wind:
• The high-density wind
becomes an extended
chromosphere (supported by
wave pressure??).
• For this case, Holzer et al.
(1983) showed the energy
equation is ~irrelevant in
determining mass flux! A
simple analytic model (of the
momentum equation) suffices.
Solar/Stellar Winds: An Overview
S. R. Cranmer, NESSC
October 27, 2008
Cool stars: Mira supergiants
• How are the presumably cool (“corona-free”)
winds of red giants and supergiants actually
accelerated?
• How do these winds affect the shapes of the
planetary nebulae that are formed at the end
of stellar evolution?
• High-luminosity: radiative driving... of dust?
• Shock-heated “calorispheres?” (Willson 2000)
Solar/Stellar Winds: An Overview
S. R. Cranmer, NESSC
October 27, 2008
Stellar winds on the HR Diagram
106
no coronae?
I
"cool"
dense
radiatively
driven winds
104
(slow?)
winds
"warm"
hybrid
winds
102
III
Be stars
"hot"
solar-type
winds
V
1
10–2
Sun

flare
stars
O
30,000
Solar/Stellar Winds: An Overview
B
A
10,000
F G
6,000
K
M
3,000
S. R. Cranmer, NESSC
October 27, 2008
Hot star winds: radiative driving
• Castor, Abbott, & Klein (1975) worked out how a hot
star’s radiation field can accelerate a time-steady wind,
even if its “Eddington factor” is << 1.
• Bound electron resonances have higher cross-sections
than free electrons (i.e., spectral lines dominate the
opacity)
• In the accelerating wind, these narrow opacity sources
become Doppler shifted with respect to the star’s
photospheric spectrum.
• Acceleration depends on velocity & velocity
gradient! (Turns F=ma on its head?
Nonlinear feedback...)
• Wind momentum (~mass loss rate times
terminal speed) scales as a “universal”
power law function of stellar luminosity.
Solar/Stellar Winds: An Overview
Kudritzki et al.
S. R. Cranmer, NESSC
October 27, 2008
Hot star winds: pulsations & waves
• O, B, A type stars have small convective cores, but not a
subsurface convection zone (i.e., no corona).
 They exhibit “g-mode” pulsations, with velocity amplitudes
often exceeding the sound speed at the photosphere!
 Can the pulsational energy “leak” out of the star into the
stellar wind? It seems to be doing it in at least one case
(BW Vul); probably many others.
 Some stars rotate so rapidly they become oblate, with some
models predicting convection and turbulence at equators.
?
MPG courtesy John Telting
Massa 1994
Solar/Stellar Winds: An Overview
S. R. Cranmer, NESSC
October 27, 2008
Hot star winds: magnetic channeling?
• O,B-star winds are highly supersonic (v > 1000
km/s) from very close to the stellar surface.
• Ud-Doula & Owocki (2002, 2008) modeled the
interaction between this kind of wind and a dipole
surface field.
• Collisions of parcels
• Fast rotation: a rigid
“magnetosphere”
develops (seen in
Ap/Bp stars?)
faster rotation
lead to stochastic
upflows/downflows.
stronger field
Solar/Stellar Winds: An Overview
S. R. Cranmer, NESSC
October 27, 2008
Conclusions
A goal of this round of NESSC:
• Positive feedback and collaboration between
the solar physics, space physics, plasma
physics, and astrophysics communities.
• What do you find interesting?
Topics for discussion (later?):
• “You do this with telescopes?!”
(observational diagnostics of stellar mass loss)
click
• When is a temperature not the temperature?
(preferential ion heating & kinetic effects)
click
• Accretion disks as torque wrenches (T Tauri)
click
• Excretion disks as torque wrenches!! (Be stars)
click
Solar/Stellar Winds: An Overview
S. R. Cranmer, NESSC
October 27, 2008
Extra slides . . .
Solar/Stellar Winds: An Overview
S. R. Cranmer, NESSC
October 27, 2008
Diagnostics of stellar mass loss
• Optical/UV spectroscopy: simple blueshifts or full
“P Cygni” profiles
• IR continuum: circumstellar dust causes SED excess
• Molecular lines (mm, sub-mm): CO, OH maser
• Radio: free-free emission from (partially
ionized?) components of the wind
(Bernat 1976)
• Continuum methods need V from
another diagnostic to get mass loss rate.
•
wind
star
• Clumping?
Solar/Stellar Winds: An Overview
(van den Oord &
Doyle 1997)
S. R. Cranmer, NESSC
October 27, 2008
Multi-line spectroscopy
• 1990s: more self-consistent treatments of radiative transfer AND better data
(GHRS, FUSE, high-spectral-res ground-based) led to better stellar wind
diagnostic techniques!
• A nice example: He I 10830 Å for TW Hya (pole-on T Tauri star) . . .
Dupree
et al.
(2006)
Solar/Stellar Winds: An Overview
S. R. Cranmer, NESSC
October 27, 2008
The solar wind mass loss rate
• The sphere-averaged “M” isn’t usually considered by solar physicists.
• Wang (1998, CS10) used empirical relationships between B-field, wind speed,
and density to reconstruct M over two solar cycles.
ACE (in ecliptic)
Solar/Stellar Winds: An Overview
S. R. Cranmer, NESSC
October 27, 2008
Sun’s mass loss history
• Did liquid water exist on Earth 4 Gyr ago? If “standard” solar models are correct,
a strong greenhouse effect was needed.
• Sackmann & Boothroyd (2003) argued that a more massive (~1.07 M) young Sun
could have been luminous enough to solve this problem, but it would have needed
strong early mass loss . . .
Sackmann & Boothroyd
(2003)
M ~ LX1.3
M ~ LX1.0
M ~ LX0.4
M ~ LX0.1
Solar/Stellar Winds: An Overview
S. R. Cranmer, NESSC
October 27, 2008
Multi-fluid collisionless effects?
Polar coronal hole
model
Solar/Stellar Winds: An Overview
S. R. Cranmer, NESSC
October 27, 2008
Multi-fluid collisionless effects?
O+5
O+6
protons
electrons
(thermal core only)
Solar/Stellar Winds: An Overview
S. R. Cranmer, NESSC
October 27, 2008
Departures from thermal equilibrium
• UVCS/SOHO observations rekindled theoretical efforts to understand collisionless
heating and acceleration effects in the extended corona.
• Ion cyclotron waves (10–10,000 Hz)
suggested as a “natural” energy source that
can be tapped to preferentially heat &
accelerate heavy ions.
MHD turbulence
Solar/Stellar Winds: An Overview
something
else?
Alfven wave’s
oscillating
E and B fields
ion’s Larmor
motion around
radial B-field
cyclotron resonancelike phenomena
S. R. Cranmer, NESSC
October 27, 2008
Angular momentum removal
• Many accreting T Tauri stars are slowly
rotating, despite the fact that disk accretion
adds angular momentum to the star (e.g.,
CTTS
(accreting)
Bouvier et al. 1993; Edwards et al. 1993).
• How is angular momentum carried away?
 By field lines that thread the disk?
(“disk locking”) This would
imply that magnetic reconnections
(and X-rays!) scale with accretion
rate. No...
 By CME-like ejections from the
tangled field in the disk?
 By a stellar wind! Matt & Pudritz
(2005, 2008) say Mwind ~ 0.1 Macc
can do the job.
Solar/Stellar Winds: An Overview
WTTS
(non-accreting)
Rot. Period (days)
L. Hartmann,
lecture notes
S. R. Cranmer, NESSC
October 27, 2008
What kind of outflow is it really?
• YSOs (Class I & II) show jets that remain
collimated far away (AU → pc!) from the
central star.
• However, EUV emission lines and He I
10830 P Cygni profiles indicate the
blueshifted outflow is close to the star.
• Stellar winds & disk winds may co-exist.
(Ferreira et al. 2006)
Solar/Stellar Winds: An Overview
S. R. Cranmer, NESSC
October 27, 2008
Be stars: “decretion disks”
• Be stars are non-supergiant B-type stars
with emission in hydrogen Balmer lines.
• Be stars are rapid rotators, but are not
rotating at “critical” / “breakup:”
Vrot  (0.5 to 0.9) Vcrit
• How does angular momentum get
added to the circumstellar gas ?
Hints:
• Many (all?) Be stars undergo nonradial pulsations (NRPs).
• Rivinius et al. (1998, 2001) found correlations between emission-line “outbursts”
and constructive interference (“beating”) between NRP periods.
• Ando (1986) & Saio (1994) suggested that some NRPs could transfer angular
momentum outwards: similar to wave “radiation pressure” (Cranmer, in prep.)
Solar/Stellar Winds: An Overview
S. R. Cranmer, NESSC
October 27, 2008