Transcript T. Thompson, Jan 2007

Feedback in Starburst Galaxies
Todd Thompson
Princeton University
with Eliot Quataert, Norm Murray, & Eli Waxman
Outline
• Goal: A model for the global structure of
starbursts.
• Why starbursts? The physical conditions.
• Radiation pressure feedback.
• Magnetic fields, cosmic rays, & -rays.
Systematics of Star Formation
• Schmidt Law:
Ý* 7g / 5

• ``Star-forming” galaxies:
– Extended, few-kpc scales.
– ~ billion year timescales.
Starbursts
• ``Starburst” galaxies:
– Compact, 100’s pc scales.
– 1-100 million year timescales.
Star-forming galaxies
• Pressure: P ~  G g2
Kennicutt (1998)
Regulation & Feedback in Galaxies
• Low star formation efficiency:
Suggests feedback and/or
regulation over a broad range
of conditions.
Starbursts
• Q~1 observed in disks.
(Martin & Kennicutt 2001)
• Stellar processes (?): Stellar
winds, radiation, supernovae,
HII regions, etc.
Star-forming galaxies
• Non-stellar processes (?): MRI.
(Sellwood & Balbus 99;Piontek & Ostriker 04)
Kennicutt (1998)
Why Starbursts?
IRAS 19297-0406
M82
M51
NOAO
NGC 253
Arp 220
Backgrounds & Starbursts
Dole et al. (2006)
Why Starbursts?
• Starbursts & U/LIRGs
– lie on the same scaling relations with normal galaxies.
– constitute a large fraction of the IR background, the star-formation
rate density at high z (also, -ray & MeV/TeV  backgrounds).
– may be a key phase in the growth of super-massive black holes &
spheroids.
– are connected physically to super-star clusters, starburst cores.
– have turbulent velocities v > 10 km/s.
• What do we want to know?
– Constituents: radiation, gas/dust, magnetic fields, and cosmic rays.
– The origin and systematics of the scaling relations of galaxies.
The Physical Conditions
Arp 220 (d ~ 80 Mpc):
• Two counter-rotating cores, ~100pc.
• Circumbinary disk R~300pc.
300 pc
•
•
•
•
•
•
•
gas ~ 5 g cm-2
n ~ 103-104 cm-3
Mgas ~ 109 -1010 M
v ~ 100 km s-1
LFIR ~ 21012 L
LX ~ 3109 L
tdyn ~ 106 n4-1/2 yr
Solomon, Sakamoto
Beswick 2006; Mundell et al; Lonsdale et al
Pressures
• Accounting:
What processes regulate Star Formation
in ULIRGs?
• The standard lore: Energy injection by supernovae, stellar
winds, HII regions (e.g., McKee & Ostriker ‘77). However, in a
dense ISM, radiative losses are large: E  n-1/4.
• Another Option: Radiation Pressure:
– Starburst photons absorbed & scattered by dust: UV ~ 100’s cm2/g.
– Dust is collisionally coupled to gas:  ~ 0.01 pc a0.1 n3-1.
– Starbursts: optically thick to re-radiated IR : IR ~ gasIR > 1.
– Radiative diffusion: efficient coupling to cold, dusty component,
most of the mass.
Scoville (2003)
(2005)
Thompson, Quataert, & Murray
Radiation Pressure Supported
Starbursts
• Radiative flux:
• Radiative diffusion:
• Radiation pressure:
• Obtain Eddington-limited starbursts:
Some Predictions
• The “Schmidt”-law for
optically-thick starbursts:
7/5
Ý
* g

g
Ý
*
Higher  implies more
pressure support, which
implies a lower star
formation rate & efficiency.


Kennicutt (1998)
The Rosseland Mean Opacity
• Sublimation: Tsub ~ 1000 K.
• Dust dominates T < 1000 K.
• At T < 200 K — in the
Rayleigh limit — = 0T2.
• Overall normalization is
dependent on metallicity
and the dust-to-gas ratio.
Semenov et al. (2003)
Some Predictions
• The “Schmidt”-law:
• When = 0T2:
no dependence on anything, but 0.
A Characteristic Flux?
ULIRGs are compact.
Intrinsic size?
Appeal to radio size,
hoping that the radio
reliably traces the star
formation.
Data from Condon et al. (1991)
Evidence for a Characteristic Flux?
Davies et al. (2006)
Why Radiation Always Wins
• Schmidt law:
• Flux:
• Radiation pressure:
Ý* 7g / 5

• Hydrostatic pressure:

• Critical surface density:
Magnetic Fields & Cosmic Rays
The FIR-Radio Correlation
How do CR electrons cool?
Radio synchrotron from CR e-’s
accelerated by SNe.
FIR traces star formation,
massive stars, SNe.
“Calorimeter” theory:
synchrotron cooling timescale
shorter than the escape time:
Starbursts
Star-Forming Galaxies
Yun et al. (2001)
tsynch < < tescape
(Völk‘89; generally unaccepted)
galaxy = CR beam dump
Magnetic Fields & Cosmic Rays
• In the Milky Way, B~5-10G and
• In starburst galaxies, how do we estimate B?
– “Minimum energy” (UB~UCR; Burbidge 1956): (~5-10G in MW).
Depends on the ratio [p/e] and on the injected CR spectral index.
– Magnetic energy density in equipartition with total hydrostatic
pressure: (~5-10G in MW)
Magnetic Fields
Conclusion:
Magnetic fields in
star-forming
galaxies are both
minimum energy &
equipartition.
and
Magnetic Fields
Conclusion:
Either
the minimum energy
estimate is wrong,
or
magnetic fields are
dynamically weak in
starburst galaxies.
Thompson et al. (2006)
Bmin Must Underestimate the True Field
UBmin/Uph measures
the importance of
synchrotron relative to
IC cooling.
If Bmin is correct, IC
dominates for
starbursts.
This contradicts the
linearity of the FIRradio correlation.
Magnetic Fields & FIR-Radio Correlation
• In the limit of very strong cooling (the “calorimeter” limit):
• The observed Schmidt Law says that
• Therefore, in the limit of strong cooling:
Magnetic Fields
Conclusion:
If a fraction ~1% of
1051 ergs per SN
goes to CR
electrons, and they
cool rapidly, the
observed trend is
reproduced.
Implies that B is in
fact larger than Bmin.
Thompson et al. (2006)
Magnetic Fields in Starbursts
• Observations thus imply rapid electron cooling.
– Strong evidence for the calorimeter theory for the FIRradio correlation: tcool< < tescape.
• So, how big is B?
– Well, B is big enough that the synchrotron cooling
timescale is << tesc. But, what is tesc?
Very uncertain:
Diffusion in MW tesc ~107.5 yrs.
Maybe advection (winds!) in starbursts tesc ~105.5 yrs (?).
Magnetic Fields in Starbursts
• Argument/Problem:
– The strongest objection to the calorimeter theory for FIRradio correlation: if synchrotron dominates cooling and
tcool< < tesc, the radio spectral indices of starbursts at GHz
should be steep “cooled” : F ~ - , with  ~ 1-1.2.
– This is not observed. Spectral indices at GHz are
~constant & not steep: F ~ - , with  ~ 0.7.
• Solution:
– If CRs interact with matter at mean density & B~Beq, then
Ionization losses dominate for low-energy CRs, not high.
This effect changes the expected slope of the radio spectrum
at a characteristic frequency ~GHz.
Magnetic Fields in Starbursts
Steeper
p=2.5
p=2.0
• Ionization losses flatten the
radio spectra
• Ionization is important only if
CRs interact with ISM of
~mean density.
• Prediction: spectral break
ubiquitous at GHz ’s for all
galaxies obeying FIR-radio.
Flatter
• Because this only works if
B~Beq, this is the best
argument for B >> Bmin in
starbursts.
Summary
• Observations indicate
– feedback is important, SF is inefficient, starbursts are dusty, disks have Q~1.
• Radiation pressure
–
–
–
–
–
can dominate feedback in the optically thick regions of starbursts.
yields qualitative change to Schmidt Law.
couples to the cold dusty component, most of the mass.
predict starburst structure: T, Teff, F, , , v, SFR/area, efficiency
are in good agreement with observations (local & high-z ULIRGs).
• Magnetic Fields in Starbursts
– are larger than Bmin and probably ~ Beq.
– are large enough that the “calorimeter” theory for FIR-radio is preferred.
– are consistent with starburst radio spectral indices only if CRs interact with
ISM of mean density so that ionization/bremsstrahlung losses are important.
• -Ray Observations of Starbursts
– will constrain the ISM density seen by CR protons.
– will constrain the energetics of CR acceleration. - Lastly, (CRp/CRe) ~ 10.
Thompson et al. (2005), (2006ab)
The Present & The Future
• Radiation pressure feedback:
– Embedded sources, porosity, transport, multi-phase ISM.
– The gravitational instability in radiation pressure dominated backgrounds.
– Starburst winds, scaling relations: Faber-Jackson, M-.
• Other mechanisms for feedback:
– HII regions, stellar winds, supernovae, gravity.
• The starburst-AGN fueling connection.
• The FIR-radio correlation:
– Test prediction of spectral breaks at GHz.
– Electron calorimetry in normal star-forming galaxies (?).
• Starbursts: what is the role of the secondary electron/positrons?
• Backgrounds: neutrino (MeV to >TeV), -ray, FIR, & radio.
• What is the energy density of cosmic rays in starburst galaxies?
The End
Constraining the Average Density “Seen”
by Cosmic Rays
-Rays from Starbursts
• Assume SNe accelerate both CR protons & electrons.
• The GeV protons collide with ambient gas:
• Proton-proton collisions produce
• If pp<< esc, then the starburst is a “proton calorimeter,” and
all of the proton energy goes into ’s (1/3), e+,-’s (1/6), and
’s (1/2).
• What is esc? As for CR electrons, very uncertain.
Thompson et al. (2006)
-Rays from Starbursts
• Massive star formation  IR emission  Supernovae:
where  is the fraction of 1051 ergs per supernova to CRp’s.
This is a FIR--Ray correlation analogous to FIR-radio.
• How do we constrain ? Assume the e+,-’s from p-p cool
via only synchrotron in the starburst:
• Observed FIR-radio correlation:
Thompson et al. (2006)
-Rays from Starbursts
NGC 253
Arp 220
-Rays from Starbursts
• If GLAST sees a larger flux from NGC 253:
– Then  > 0.05  more energy per SN to CR protons.
– Because
from secondary
electrons/positrons, another process (not synchrotron) must
dominate CR electron cooling.
• If GLAST sees a smaller flux from NGC 253:
– Either the CRs interact ISM below mean density, rapid escape,
– or,  < 0.05  less energy per SN to CR protons.
– These options can in principle be distinguished by modeling the IC
and relativistic bremsstrahlung emission at -ray energies since the
latter also depends on density.
The Diffuse -Ray Background
• Massive star formation  IR emission  Supernovae.
+ star formation rate history of the universe.
+ the fraction of all star formation at high-z that occurs in “proton
calorimeters” (high density).
• For an individual galaxy:
• For the history of star formation:
Thompson et al. (2006)
The -Ray Background