Download Sternentstehung - Star Formation

Sternentstehung - Star Formation
Sommersemester 2006
Henrik Beuther & Thomas Henning
24.4
1.5
8.5
15.5
22.5
29.5
5.6
12.6
19.6
26.6
3.7
10.7
Today: Introduction & Overview
Public Holiday: Tag der Arbeit
--Physical processes, heating & cooling, cloud thermal structure
Physical processes, heating & cooling, cloud thermal structure (H. Linz)
Basic gravitational collapse models I
Public Holiday: Pfingstmontag
Basic gravitational collapse models II
Accretion disks
Molecular outflow and jets
Protostellar evolution, the stellar birthline
Cluster formation and the Initial Mass Function
17.7 Massive star formation, summary and outlook
24.7 Extragalactic star formation (E. Schinnerer)
More Information and the current lecture files: http://www.mpia.de/homes/beuther/lecture_ss06.html
[email protected]
Summary Clusters and the IMF
- The most widespread mode of star formation is clustered.
- The IMF is almost universally valid, Salpeter-like for >1Msun, characteristic
mass plateau around 0.5Msun.
- Mass segregation in clusters.
- Hertzsprung-Russel diagrams allow to estimate cluster ages.
- Anomalous IMFs (e.g., Taurus or Arches).
- Cloud mass distributions versus pre-stellar core mass distributions.
- Gravitational fragmentation.
- Gravo-turbulent fragmentation.
- Competitive accretion.
- General features of the IMF. The plateau at 0.5Msun maybe explicable by
thermal physics.
- High-Mass star formation differences?
- Order of star formation.
Massive Star Formation
- Why important?
- Although few in numbers, LM3 they
inject significant amounts of energy into
ISM during their lifetime (outflows,
radiation, supernovae).
- They produce all the heavy elements.
- Low-mass star formation is strongly
influenced by massive stars.
- Massive stars exclusively form in clusters.
- Short Kelvin-Helmhotz contraction time:
tKH=3x107 yr (M*/1Ms)2 (R*/1Rs)-1(L*/1Ls)-1
For 60Ms, 12Rs and 105.9Ls we find:
tKH ~ 11000 yr
--> no observable pre-main sequence
evolution.
- Radiation pressure important constraint.
Eddington luminosity and limit
The Eddington luminosity/limit gives the upper limit on the luminosity a star
can have before it becomes unstable and blows away the gas again.
- Assumptions: spherical symmetry and fully ionized hydrogen.
--> Radiation exerts force mainly on free electrons via Thomson
scattering sT=(q2/mc2)2 (q: charge, m: mass particle)
- Outward radial force equals rate at which electron absorbs momentum:
sTS/c
(S: energy flux)
--> In effect radiation pushes out electron-proton pairs against total
gravitational force GM(mp+me)/r2 = GMmp/r2
- With the flux S = L/4πr2, the force equilibrium is:
GMmp/r2 = sTL/(4πr2c)
--> The Eddington luminosity: LEdd = 4πGMc (mp/sT)
= 4πGMc/k
(with k=sT/mp mass abs. coefficient)
= 1.3 x 1038 (M/Msun) erg/s
[Lsun] ~ 3 x 104 (M/Msun)
or
Ledd
- If L > Ledd then
- accretion stops if L provided by accretion
- Gas layers pushed out and star unstable if provided by nuclear fusion.
- Scaling relations for massive (proto)stars: L  Ma with 2<a<4
Radiation pressure
- In contrast now the radiation pressure of the central massive (proto)star
on the surrounding dust cocoon. Same relation:
L/M = 4πGc/k
(k: mass abs. Coefficient)
- While k is very low for the fully ionized H plasma (k~0.3cm2g-1), at the
dust destruction front (T~1500K) it is considerably larger with k~10cm2g-1.
-->
L/M ~ 103 [Lsun/Msun]
--> In spherical symmetric accretion models, accretion is expected to stop
as soon as the luminosity is approximately 1000 times larger than the
mass of the protostar. --> No problem for low-mass star formation.
--> The critical ratio is reached for stars of approximately 11Msun.
Since more massive stars are know, the assumption of spherical
accretion has to be wrong and other processes are needed.
Possible ways out
- Decrease mass absorption coefficient k.
- Deviate from spherical symmetry. Disks are known for low-mass
star formation. High accretion rates necessary.
- Coalescence of low- to intermediate-mass stars at the center of
extremely dense massive protoclusters.
- Competitive accretion caused by the cluster potential.
Change of mass absorption coefficient k
G: radiative to
gravitational
acceleration.
Wolfire & Cassinelli 1987
- k can be changed if:
- radiation field is shifted from optical
to far infrared
- dust composition changes significantly:
- reduce dust-to-gas ratio
- reduce grain sizes
Yorke 2004
Accretion rate dependence
In spherical gravitational infall
gravity governs accretion rate.
dM/dt ~ 4πr2rv
with 1/2mv2 = GMm/r
--> v = √2GM/r
Hence to overcome the radiation
pressure to form more massive
stars larger accretion rates are
necessary. However Eddington
limit sets upper boundary.
Beyond dust
destruction radius
Wolfire & Cassinelli 1987
Turbulent accretion: scaling up low-mass sf
- A 100Msun star forms in ~105yrs --> average accretion rate ~10-3Msun/yr
- Standard low-mass accretion rate: dM/dt ~ cs3/G ~ 2x10-6Msun/yr.
- To form massive stars McKee & Tan (2002, 2003) developed “turbulent
core model”, in which massive stars form within gravitationally bound
cores supported by turbulence and magnetic fields.
--> The turbulent support raises the sound speed cs and hence dM/dt.
dM/dt ~ 0.5x10-3 (Mfinal/30Msun)3/4 Scl3/4 (M/Mfinal)0.5
t ~ 1.3x105 (Mfinal/30Msun)1/4 Scl-3/4
[Msun/yr]
[yr]
Scl: cloud surface density ~ 1g cm-2
Forming massive accretion disks
- 2D frequency-dependent hydrodynamic simulations
reveal that short-wavelength radiation (most
effective for radiative acceleration) can escape in the
polar directions. Long-wavelength radiation is more
or less isotropic.
- In these models radiation driven outflows in the
polar direction and disks in the equatorial direction
are forming.
--> These models imply that massive stars may
form in a similar fashion as low-mass protostars
due to disk accretion.
Yorke & Sonnhalter 2002
Other lines:
Non-accreting protostars
Yorke & Sonnhalter
2002
Protostellar cavities reduce radiation pressure
Models for
different
outflow angles
and cavity
shapes.
Black: grav.
force
Col: rad.
force for
varying
degrees from
north (q)
and opening
Angles (q0)
Krumholz et al.
2005
- Protostellar outflows produce cavities which can work as valves releasing
large fractions of the radiative pressure.
- Away from the outflow cavity, gravity still overwelms radiation.
Radiative bubbles and instabilities
1.5x104yr
21.3Msun
1.65x104yr
22.4Msun
2.0x104yr
25.7Msun
Krumholz 2005
- 3D radiation hydrodynamic simulations, starting with 100 to 200Msun cores.
- First, until approximately 17Msun protostellar mass: smooth accretion flow,
low angular momentum gas accretes directly on protostar, high angular
momentum gas forms Keplerian accretion disk.
- From 17Msun upwards, radiation pressure stars driving out gas, bubbles form.
Further infalling gas moves along the bubble walls and falls onto disk.
- At ~22Msun bubbles become Rayleigh-Taylor instable and collapse again. This
happens when radiation force gets too low and the bubble walls encounter
net gravitational force. (Rayleigh-Taylor instabilities naturally occur when heavy fluids
(gas) are accelerated by light fluids (radiation).)
- Mean accretion rate 2.8x10-3Msun/yr
- No further fragmentation likely due to the heating from accretion/nuclear L.
Trapped hypercompact HII regions
- Stars >~ 10Msun produce enough UV radiation to ionize the surrounding gas.
- Since the temperature of the ionized region is about 100 times larger than
that of the molecular gas (10000 vs. 100K), the pressure of the small HII
region should be 100 greater as well (P=raT2 with aT√T for ideal gas):
--> Expansion of hypercompact HII region could stop accretion.
- Possible solution: if size of HCHII region is below sonic radius where the
necessary escape speed is larger than thermal sound speed of ionized gas.
This can happen if the accretion flow is dense enough that the radius of
recombination were smaller than the sonic radius.
- Keto (2002, 2003) included gravity in classic expansion models of HCHII
regions and found that a steep density gradient and accretion flow can be
maintained in the ionized gas confining the HCHII to the required small sizes.
Accertion solution
Sonic point
Free-fall with nr-3/2
And hydrostatic equlibrium
With nexp(2rs/r)
Keto 2002
ionized
neutral
Sonic point
Coalescence and merging
Bally & Zinnecker 2005
Bonnell et al. 1998
-
Required (proto)stellar densities of the order 106 to 108 stars per pc3.
Very explosive events expected.
Collimated outflows and jets can barely survive.
The IMF should have a dip at intermediate masses.
Competetive Accretion
Bonnell et al. 2004
- Gas clump first fragments into a large number
of clumps with approximately a Jeans mass.
Hence fragmentation on smaller scales.
- Then each clump subsequently accretes gas
from the surrounding gas potential. Even gas
that was originally far away may finally fall
onto the protostar.
Distance of
gas that is
ultimately
accreted.
Single-dish maps of massive outflows
- Early observations claimed different
collimation degrees for massive
outflows --> different formation?
- Then single-dish data taking into
distance and spatial resolution were
consistent with low-mass outflows.
- Outflow-mass scales with core mass.
- Many outflow properties scale over
many orders of magnitude.
- Outflow force implies non-radiative
outflow driving.
- High outflow rates imply high accretion
rates.
Force vs. L
Beuther et al. 2002
Wu et al. 2004, 2005
Outflow rate vs. L
Collimated massive jets and outflows
IRAM 30m data
Grey: 1.2mm
Contours:CO(2-1)
Beuther et al. 2002
An evolutionary scenario
- Outflows are ubiquitous phenomena
- Jet-like outflows exists at least up to early-B and late-O-type stars
- The observations suggests tentatively an evolutionary scenario
Disks in Massive Star Formation
Theory: Flattened rotating disk-structure with a Keplerian velocity profile
and a mass that is usually dominated by central star.
Observation: Velocity gradient perpendicular to the outflow.
Velocity [km/s]
Best current example IRAS20126+4102, Cesaroni et al. 2005
Offset [‘’]
The data imply a 7Msun central source. It may still be in its main
accretion phase because L~104Lsun can be well explained by accretion
Luminosity: Lacc = G(dM/dt)M*/R*.
The disk-outflow system in Orion-source I
Potentially disk-tracing molecules
IRAS 23151+5912: CH3OH vt=1
Beuther et al., in prep.
IRAS 18089-1732: HCOOCH3
Beuther et al. 2004, 2005
IRAS 20126+4104: CH3CN & C34S
Cesaroni et al. 1997, 1999, 2005
G29.96:
HN13C(4-3)
Moment 1 for HN13C(4-3)
Beuther et al.,
in prep.
G24.78 & G31.41: CH3CN
Beltran et al. 2004, 2005
Fragmentation of a massive protocluster
• 12 clumps within
each core
• Integrated masses
98Msun (south)
42Msun (north)
--> 80 to 90% of
the gas in halo
• Core masses
1.7Msun to 25Msun
• Column densities
1024cm-2 ->Av~1000
Spatial filtering affects only large scale halo
on scales >20’’
Assumptions: - All emission peaks of protostellar nature
- Same temperature for all clumps (46K, IRAS)
Caveats: - Temperature structure
- Peaks due to different emission processes, e.g., outflows?
Beuther & Schilke 2004
Hot core chemistry
Beuther et al. 2004, 2005
Summary
- Massive stars are very important for energy budget and nucleosynthesis.
- They form exclusively in a clustered mode.
- They have very short Kelvin-Helmholtz contraction times and hence no
optically observable pre-main sequence evolution.
- Large radiation pressure has to be overcome.
- Two main proposals: (1) scale up low-mass star formation scenario
(turbulent core model) with accretion disks and enhanced accretion rates.
(2) Turn more dynamical, competitive accretion, coalescence and merging.
- Current observations support the accretion model, collimated outflows and
disk-like structures have been found. Fragmentation is also indicative of
large, massive bound cores from early on.
- No current evidence for coalescence and merging. Seems not necessary but
may exist is selected sources.
- Current debate largely between “turbulent core” and “competitive accretion”
people. Observational discrimination necessary, e.g., are largest fragments
stable? Do we find global collapse?
- Chemistry interesting additional feature of Hot Molecular Cores.
Sternentstehung - Star Formation
Sommersemester 2006
Henrik Beuther & Thomas Henning
24.4
1.5
8.5
15.5
22.5
29.5
5.6
12.6
19.6
26.6
3.7
10.7
17.7
Today: Introduction & Overview
Public Holiday: Tag der Arbeit
--Physical processes, heating & cooling, cloud thermal structure
Physical processes, heating & cooling, cloud thermal structure (H. Linz)
Basic gravitational collapse models I
Public Holiday: Pfingstmontag
Basic gravitational collapse models II
Accretion disks
Molecular outflow and jets
Protostellar evolution, the stellar birthline
Cluster formation and the Initial Mass Function
Massive star formation, summary and outlook
24.7 Extragalactic star formation (E. Schinnerer)
More Information and the current lecture files: http://www.mpia.de/homes/beuther/lecture_ss06.html
[email protected]