From Massive Cores to Massive Stars

1. From Massive Cores to Massive Stars Mark Krumholz Princeton University Collaborators: Richard Klein, Christopher McKee (UC Berkeley) Kaitlin Kratter, Christopher Matzner (U. Toronto) Jonathan Tan (University of Florida) Todd Thompson (Princeton University) Hubble Fellows Symposium April 4, 2007

2. Talk Outline • What are massive cores? • From core to star • Fragmentation • Disk Formation and Evolution • Binary formation • Radiation pressure and feedback • Final summary

3. Sites of Massive Star Formation(Plume et al. 1997; Shirley et al. 2003; Rathbone et al. 2005; Yonekura et al. 2005) • Massive stars form in gas clumps seen in mm continuum or lines, or in IR absorption (IRDCs) • Typical properties: • M ~ 103 - 104 M • R ~ 1 pc •  ~ 1 g cm–2 •  ~ few km s-1 • Properties very similar to young rich clusters • Appear virialized pc pc Spitzer/IRAC (left) and Spitzer/MIPS (right), Rathbone et al. (2005)

4. Massive Cores in Clumps • Largest cores in clumps: M ~ 100 M, R ~ 0.1 pc, centrally condensed • Question: are these the progenitors of massive stars? Cores in IRDC 18454-0158, MSX 8 m (grayscale), 1.2 mm IRAM 30m (contours), Sridharan et al. (2005) Core in IRDC 18223-3, Spitzer/IRAC (color) and PdBI 93 GHz continuum (contours), Beuther et al. (2005, 2007)

5. The Core Mass Function(Motte, Andre, & Neri 1998, Johnstone et al. 2001, Reid & Wilson 2005, 2006, Lombardi et al. 2006) • The core MF is similar to the stellar IMF, but shifted to higher masses a factor of 2 – 4 • Correspondence suggests a 1 to 1 mapping from core  star with roughly constant efficiency • Caveat: blending in distant, massive regions Core mass function in Pipe Nebula (red) vs. stellar IMF (gray) (Alves, Lombardi, & Lada 2007)

6. Massive Star Clustering > 5 M Fraction of stars vs. radius for stars of different masses in the ONC (Hillenbrand & Hartmann 1998) < 5 M • Most (all?) massive stars are born in clusters • Massive stars in young clusters are strongly mass-segregated, but lower mass stars are not(Hillenbrand & Hartmann 1998, Huff & Stahler 2006)

7. The Clustering of Cores(Elmegreen et al. 2001, Stanke et al. 2006) • Cores in clumps are also mass segregated • Mass function below ~ 5 M is the same everywhere; more massive cores only found near center • Similar to stellar mass segregation, also suggests core  star mapping Core mass function for inner (red) and outer (blue) parts of  Oph, Stanke et al. (2006)

8. From Core to Star

9. Stage 1: Initial Fragmentation(Krumholz, 2006, ApJL, 641, 45) • Massive cores are much larger than MJ (~ M), so one might expect them to fragment while collapsing (e.g. Dobbs et al. 2005) • However, accretion can produce > 100 L even when protostars are < 1 M m*=0.05 M m*=0.8 M Temperature vs. radius in a massive core before star formation (red), and once protostar begins accreting (blue)

10. Radiation-Hydro Simulations • To study this effect, do simulations • Use the Orion code to solve equations of hydrodynamics, gravity, radiation (in flux-limited diffusion approximation) on an adaptive mesh (Krumholz, Klein, & McKee 2007a, ApJ, 656, 959, and KKM, 2007b, ApJS, submitted, astro-ph/0611003) Mass conservation Momentum conservation Gas energy conservation Rad. energy conservation Self-gravity

11. Simulation of a Massive Core • Simulation of 100 M, 0.1 pc turbulent core • LHS shows  in whole core, RHS shows 2000 AU region around most massive star

12. Massive Cores Fragment Weakly • With RT: 6 fragments, most mass accretes onto single largest star through a massive disk • Without RT: 23 fragments, stars gain mass by collisions, disk less massive, disrupted by collisions and N-body interactions • Conclusion: radiation inhibits fragmentation Column density with (upper) and without (lower) RT, for identical times and initial conditions

13. Stage 2: Massive Disks(Kratter & Matzner 2006, Kratter, Matzner & Krumholz, 2007, in preparation) • Accretion rate onto star + disk is ~ 3 / G ~ 10–3 M / yr in a massive core • Maximum accretion rate through a stable disk via MRI or local GI is ~ cs3 / G ~ 5 x 10–5 M / yr for a disk with T = 100 K • Conclusion: cores accrete faster than stable disks can process, so disks become massive and unstable. Depending on thermodynamics, they may fragment.

14. Massive Disks in Simulations(KKM 2007a) • Disks reach Mdisk ~ M* / 2, r ~ 1000 AU • Global GI creates strong m = 1 spiral pattern • Spiral waves drive rapid accretion; eff ~ 1 • Radiation keeps disks radially isothermal • Disks reach Q ~ 1, unstable to fragment formation Surface density (upper) and Toomre Q (lower); striping is from projection

15. Observing Massive Disks Integrated TB in simulated 1000 s / pointing ALMA observation of disk at 0.5 kpc in CH3CN 220.7472 GHz (KKM 2007c, ApJ, submitted)

16. Stage 3: Binary Formation • Most massive stars in binaries (Preibisch et al. 2001) • Binaries often very close, a < 0.25 AU • Mass ratios near unity (“twins”) common (Pinsonneault & Stanek 2006, Bonanos 2007) • Most massive known binary is WR20a: M1 = 82.7 M, q = 0.99 ± 0.05 (Rauw et al. 2005) Mass ratio for 26 detached elcipsing binaries in the SMC (Pinsonneault & Stanek 2006)

17. Close Binaries(Krumholz & Thompson, 2007, ApJ, in press, astro-ph/0611822) • Stars migrate through disk to within 10 AU of primary • Some likely merge, some form tight binaries, a < 1 AU • Protostars with masses 5 – 15 M reach radii ~ 0.1 AU due to deuterium shell burning Radius vs. mass for protostars as a function of accretion rate

18. Mass Transfer and “Twins” WR20a • Large radii likely produce RLOF, mass transfer • Transfer is from more to less massive  transfer unstable • System becomes isentropic contact binary, stabilizes at q 1, contracts to MS • Result: massive twin Minimum semi-major axis for RLOF as a function of accretion rate

19. Massive protostars reach MS in a Kelvin time: This is shorter than the formation time  accretion is opposed by huge radiation pressure on dust grains Question: how can accretion continue to produce massive stars? Stage 4: Radiation Pressure

20. Radiation Pressure in 1D(Larson & Starrfield 1971; Kahn 1974; Yorke & Krügel 1977; Wolfire & Cassinelli 1987) • Dust absorbs UV & visible, re-radiates IR • Dust sublimes at T ~ 1200 K, r ~ 30 AU • Radiation > gravity for • For 50 M ZAMS star, In reality, accretion isn’t spherical. Investigate 3D behavior with Orion.

21. Simulations of Radiation Pressure (KKM, 2005, IAU 227)

22. Beaming by Disks and Bubbles • 2D and 3D simulations reveal flashlight effect: disks and bubbles collimate radiation • At higher masses, radiation RT instability possible Collimation allows accretion to high masses! Density and radiation flux vectors from simulation

23. Beaming by Outflows(Krumholz, McKee, & Klein, ApJL, 2005, 618, 33) • Massive stars have outflows launched from dust destruction zone • Outflow velocity ~103 km s–1 no time for grains to re-grow until gas is far from star • Result: outflow cavities optically thin, radiation can leak out of them • Simulate with MC radiative transfer code Gas temperature distributions with a 50 M star, 50 M envelope

24. Outflows Help Accretion • Simulations show that this effect can produce an order-of-magnitude reduction in radiation force • This can move the system from radiation being stronger than gravity to weaker than gravity Radiation and gravity forces vs. radius for a 50 M star and a typical outflow cavity geometry

25. Summary • Massive stars form from massive cores • Massive cores fragment only weakly • They produce disks that are massive and unstable. This is probably observable. • Close massive binaries likely experience mass transfer, which explains massive twins • Radiation pressure does not halt accretion • Mass and spatial distributions of massive stars are inherited from massive cores • However, every new bit of physics added has revealed something unexpected…

26. Plan B Give up and appeal to intelligent design…