From Massive Cores to Massive Stars

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					From Massive Cores to
    Massive Stars
          Mark Krumholz
       Princeton University
 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
   Question: are these the
    progenitors of massive stars?

Cores in IRDC
MSX 8 m
1.2 mm IRAM
(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
   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
   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.

   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
     Simulation of a Massive Core

                           QuickTime™ and a
                      YUV420 codec decompressor
                     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
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

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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

   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
   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
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    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

                                 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
    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
   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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