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From Massive Cores to Massive Stars

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					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
                  Talk Outline
   What are massive cores?
   From core to star
       Fragmentation
       Disk Formation and Evolution
       Binary formation
       Radiation pressure and feedback
   Final summary
  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
              pc                pc                ~ few km s-1
                                             Properties very similar
Spitzer/IRAC (left) and Spitzer/MIPS
                                              to young rich clusters
(right), Rathbone et al. (2005)              Appear virialized
         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),                                   Core in IRDC 18223-3,
Sridharan et                        Spitzer/IRAC (color) and PdBI 93
al. (2005)                                GHz continuum (contours),
                                          Beuther et al. (2005, 2007)
         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                     Core mass function in Pipe Nebula
    distant, massive regions                (red) vs. stellar IMF (gray) (Alves,
                                            Lombardi, & Lada 2007)
          Massive Star Clustering
             > 5 M
                                                      Fraction of stars
                      < 5 M                          vs. radius for stars
                                                      of different masses
                                                      in the ONC
                                                      (Hillenbrand &
                                                      Hartmann 1998)



   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)
           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
Core mass function for inner (red) and
outer (blue) parts of  Oph, Stanke et
                                              suggests core  star
al. (2006)                                    mapping
From Core to Star
      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                     Temperature vs. radius in a
                                           massive core before star
    can produce > 100 L                   formation (red), and once
    even when protostars                   protostar begins accreting (blue)

    are < 1 M
     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
     Simulation of a Massive Core


                           QuickTime™ and a
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                     are neede d to see this picture.




   Simulation of 100 M, 0.1 pc turbulent core
   LHS shows  in whole core, RHS shows 2000 AU
    region around most massive star
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
Column density with (upper) and
                                       Conclusion: radiation
without (lower) RT, for identical       inhibits fragmentation
times and initial conditions
            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.
     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
       Observing Massive Disks



                                   QuickTime™ and a
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                             are neede d to see this picture.




Integrated TB in simulated 1000 s / pointing ALMA observation of disk at 0.5
kpc in CH3CN 220.7472 GHz (KKM 2007c, ApJ, submitted)
         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 =
Mass ratio for 26 detached elcipsing       82.7 M, q = 0.99 ±
binaries in the SMC (Pinsonneault &
Stanek 2006)
                                           0.05 (Rauw et al. 2005)
                   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
        Mass Transfer and “Twins”
                                            Large radii likely
WR20a                                        produce RLOF,
                                             mass transfer
                                            Transfer is from
                                             more to less massive
                                              transfer unstable
                                            System becomes
                                             isentropic contact
                                             binary, stabilizes at q
 Minimum semi-major axis for RLOF as a        1, contracts to MS
 function of accretion rate
                                            Result: massive twin
     Stage 4: Radiation Pressure
   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?
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.
Simulations of Radiation Pressure
            (KKM, 2005, IAU 227)




                    QuickTime™ and a
               YUV420 codec decompressor
              are neede d to see this picture.
    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
            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
Gas temperature distributions with
a 50 M star, 50 M envelope              radiative transfer code
          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
                    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…
                Plan B




Give up and appeal to intelligent design…

				
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