Feedback in Starburst Galaxies

1. Feedback in Starburst Galaxies Todd Thompson Princeton University with Eliot Quataert, Norm Murray, & Eli Waxman

2. Outline • Goal: A model for the global structure of starbursts. • Why starbursts? The physical conditions. • Radiation pressure feedback. • Magnetic fields, cosmic rays, & -rays.

3. Systematics of Star Formation • Schmidt Law: • ``Star-forming” galaxies: • Extended, few-kpc scales. • ~ billion year timescales. • ``Starburst” galaxies: • Compact, 100’s pc scales. • 1-100 million year timescales. • Pressure: P ~  G g2 Starbursts Star-forming galaxies Kennicutt (1998)

4. Regulation & Feedback in Galaxies • Low star formation efficiency: Suggests feedback and/or regulation over a broad range of conditions. • Q~1 observed in disks. (Martin & Kennicutt 2001) • Stellar processes (?): Stellar winds, radiation, supernovae, HII regions, etc. • Non-stellar processes (?): MRI. (Sellwood & Balbus 99;Piontek & Ostriker 04) Starbursts Star-forming galaxies Kennicutt (1998)

5. Why Starbursts?

6. M82 IRAS 19297-0406 M51 NOAO NGC 253 Arp 220

7. Backgrounds & Starbursts Dole et al. (2006)

8. 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.

9. The Physical Conditions Arp 220 (d ~ 80 Mpc): • Two counter-rotating cores, ~100pc. • Circumbinary disk R~300pc. • 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 300 pc Solomon, Sakamoto Beswick 2006; Mundell et al; Lonsdale et al

10. Pressures • Accounting:

11. 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) Thompson, Quataert, & Murray (2005)

12. Radiation Pressure Supported Starbursts • Radiative flux: • Radiative diffusion: • Radiation pressure: • Obtain Eddington-limited starbursts:

13. Some Predictions • The “Schmidt”-law for optically-thick starbursts: Higher  implies more pressure support, which implies a lower star formation rate & efficiency. Kennicutt (1998)

14. 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)

15. Some Predictions • The “Schmidt”-law: • When = 0T2: no dependence on anything, but 0.

16. 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)

17. Evidence for a Characteristic Flux? Davies et al. (2006)

18. Why Radiation Always Wins • Schmidt law: • Flux: • Radiation pressure: • Hydrostatic pressure: • Critical surface density:

19. Magnetic Fields & Cosmic Rays

20. 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: tsynch < < tescape (Völk‘89; generally unaccepted) galaxy = CR beam dump Starbursts Star-Forming Galaxies Yun et al. (2001)

21. 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)

22. Magnetic Fields Conclusion: Magnetic fields in star-forming galaxies are both minimum energy & equipartition. and

23. Magnetic Fields Conclusion: Either the minimum energy estimate is wrong, or magnetic fields are dynamically weak in starburst galaxies. Thompson et al. (2006)

24. BminMust 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 FIR-radio correlation. UBmin /Uph

25. 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:

26. 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)

27. Magnetic Fields in Starbursts • Observations thus imply rapid electron cooling. • Strong evidence for the calorimeter theory for the FIR-radio 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 (?).

28. Magnetic Fields in Starbursts • Argument/Problem: • The strongest objection to the calorimeter theory for FIR-radio 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: • IfCRs 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.

29. Magnetic Fields in Starbursts • 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. • Because this only works if B~Beq, this is the best argument for B >> Bmin in starbursts. Steeper p=2.5 p=2.0 Flatter

30. 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)

31. 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?

32. The End

33. Constraining the Average Density “Seen” by Cosmic Rays

34. -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)

35. -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)

36. -Rays from Starbursts NGC 253 Arp 220

37. -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.

38. 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)

39. The -Ray Background