Transcript The Formation of Low Mass Stars: Overview and Recent

Star and Planet
Formation
Neal Evans
The University of Texas at Austin
Star Formation in Galaxies
CO(1-0): OVRO
Antennae: galaxy
merger.
Visible (HST) shows
copious star
formation, but misses
the main show. Most
intense star formation
in obscured region
traced by CO.
Whitmore et al. (1999)
Wilson et al. (2000)
New Views are Coming
The Archive is open:
http://ssc.spitzer.caltech.edu/
Star Formation/Galaxy
Formation
 Key


part of galaxy formation
Properties of molecular clouds in other
galaxies
Connect studies of distant galaxies to MW
• Collective, clustered, massive SF
• Molecular line probes of high density
• Dust continuum emission

Insights into high-z starbursts
Star Formation traced by HCN
Relation between LIR and
LCO becomes non-linear
for very high LIR.
12
H C N 1-0 on galaxies
G ao & S olom on, 2003
log[L IR ] (L SUN )
10
Stays linear for LHCN.
S lo p e is clo se to 1
C O 1-0 on galaxies
S olom on et al. 1997
8
Is this correlation still correct
for sm aller and larger scale ??
6
4
4
6
8 -2
log[L' M o lecu le ] (K km /s pc )
10
J. Wu et al. Data from Gao
and Solomon 2003.
Relation Same for Cores in MW
15
APM08279
Gao galaxy sample HCN1-0
Solomon galaxy sample CO 1-0
Paglione Galaxy sample HCN 3-2
(Stars) High Z CO 1-0 galaxies
Galactic massive cores HCN 3-2
Galactic massive cores CS 5-4
Cloverleaf
log[L IR ] (L SUN)
10
y=1.31X -1.19
y=0.98X +3.81
5
G alactic C O clouds
y=0.96X +1.19
(M ooney & S olom on,1988)
?
y=0.97X +3.10
LHCN or LCS for cores in
MW also linear with LIR.
(For LIR > 104 Lsun)
These points are for
HCN 3-2 and CS 5-4.
Same line as HCN 3-2
in galaxies. Parallel to
HCN 1-0 relation for
starbursts.
?
0
-4
-3
-2
-1
0
1
2
3
4
5
6
7
8
9
10 11 12
Checking HCN 1-0 in
MW.
-2
log[L' M olecule ] (K km /s pc )
J Wu et al. In prep.
What probes can we use?
 Dust

Extinction of background stars
• Probe Nd(b), Bperp

Emission in infrared to millimeter
• Probe Td( r), Nd(b), Bperp

Problem: need to know dust properties
 Molecules

Emission or absorption (infrared to radio)
• Probe TK( r), n( r), v( r), Bpar

Problem: chemistry
Studies of High Mass Regions

Many Detailed Studies


Ho, Zhang, …
Surveys



van der Tak et al. (2000) (14 sources)
Beuther et al. (2002) (69 sources)
Survey of water masers for CS
• CS survey Plume et al. (1991, 1997)

Dense: <log n> = 5.9
• Maps of 51 in 350 micron dust emission

Mueller et al. 2002
• Maps of 63 in CS J = 5–4 emission

Shirley et al. 2003
Luminosity versus Mass
Log Luminosity vs. Log M
red line: masses of dense
cores from dust
Log L = 1.9 + log M
blue line: masses of GMCs
from CO
Log L = 0.6 + log M
L/M much higher for dense
cores than for whole
GMCs.
Mueller et al. (2002)
Linewidth versus Size
Correlation is weak.
Linewidths are 4-5
times larger than in
samples of lower
mass cores.
Massive clusters form
in regions of high
turbulence, pressure.
Shirley et al. 2003
Cumulative Mass Function
Incomplete below 103 Msun.
Fit to higher mass bins gives slope
of about –0.93.
Steeper than that of CO clouds or
clumps (–0.5 on this plot).
Similar to that of clusters,
associations (Massey et al. 1995)
in our Galaxy and in Antennae
(Fall et al. 2004).
Shirley et al. 2003
Massive Cores: Gross
Properties

Massive, Dense, Turbulent



Mass distribution closer to clusters, stars than GMC
Much more turbulent than low mass cores
A model for starbursts?


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
Luminosity correlates well with core mass
Less scatter than for GMCs as a whole
L/M much higher than for GMCs as a whole
L/Mdust ~ 1.4 x 104 Lsun/Msun ~ high-z starbursts
L/L(HCN) similar to starbursts
Starburst: all gas like dense cores?
Hints of Dynamics
A significant fraction of the massive core
sample show self-reversed, blue-skewed
line profiles in lines of HCN 3-2.
Of 18 double-peaked profiles, 11 are blue, 3
are red.
Suggests inflow motions of overall core.
Vin ~ 1 to 4 km/s over radii of 0.3 to 1.5 pc.
J. Wu et al. (2003)
Open Questions for ALMA

Studies of gas, dust, high n tracers in galaxies
 Detailed structure of massive cores

Can we separate into fragments/clusters?
• Simulations predicting properties
• Understand IMF?
• See SMA early results as preview for ALMA

Can we study dynamics?
• Test inward motion hints in single-dish spectra
• Separate dynamics of fragments

Evolution of dust, ice, gas-phase chemistry
• Combine ALMA with Spitzer, SOFIA, Herschel, …
What do we need?
 High
resolution
 High dynamic range, image fidelity

Bright, but complex, sources
 Flexible

correlator
Very rich spectrum, need many diagnostics
 Full
complement of receivers
 For exgal clouds, excellent sensitivity
Low Mass vs. High Mass
 Low

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
“Isolated” (time to form < time to interact)
Low turbulence (less than thermal support)
Nearby (~ 100 pc)
 High

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
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Mass star formation
Mass star formation
“Clustered”
Time to form may exceed time to interact
Turbulence >> thermal
More distant (>400 pc)
High vs. Low Early Conditions
Property
Low
High
p
~1.8
~1.8
nf (median)
2 x 105
1.5 x 107
Linewidth
0.37
5.8
n( r) = nf (r/rf)–p ; rf = 1000 AU
Even “Isolated” SF Clusters
Taurus Molecular Cloud
Prototypical region of
“Isolated” star formation
Myers 1987
But Not Nearly as Much
Taurus Cloud at same scale
4 dense cores, 4 obscured stars
~15 T Tauri stars
1 pc
Orion Nebula Cluster
>1000 stars
2MASS image
The Basic Features
Envelope
Disk
Protostar
Jet/wind/outflow
T. Greene
Studies of the Envelope
 All

quantities vary along line of sight
Dust temperature, Td( r)
• Heating from outside, later inside

Gas temperature, TK( r)
• Gas-dust collisions, CRs, PE heating



Density, n(r), predicted to vary
Velocity, v(r), connected to density
Abundance, X(r), varies
• Photodissociation, freeze-out, desorption
Combined Modeling of Dust and Gas
nd(r), L

Radiative
Transfer
TD(r)
Gas to
dust
S
I(b)
Dust
PhysicalModel
Model
Physical
n(r),v(r)
v(r)
n(r),
Iterate
TD(r) to
TK(r)
n(r), T K(r)
v(r)
Simulate
Observations
Observations
Observations
Gas
Chemistry
X(r)
Monte
Carlo
nJ(r)
Simulate
Obs.
TA (v,b)
An Evolutionary Model

A physical model from theory


Sequence of Bonnor-Ebert spheres of increasing nc
e.g., Shu (1977) “Inside-out collapse”

Calculate luminosity of central star+disk
 Dust temperature through envelope
 Gas temperature
 Chemical abundances


Follow gas as it falls, using evolving conditions
Line Profiles including all effects
Theory gives n(r,t), v(r,t)
QuickTime™ and a
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t<0: Series of Bonnor-Ebert spheres
t>0: Inside-out collapse model (Shu 1977)
C. Young
L(t) from Accretion,
Contraction
L(t) calculated.
First accretion.
First onto large (5 AU) surface
(first hydrostatic core).
Then onto PMS star with
R = 3 Rsun, after 20,000 to
50,000 yr. And onto disk.
Prescriptions from Adams and
Shu.
Contraction luminosity and
deuterium burning dominates
after t ~100,000 yr.
C. Young and Evans, in prep.
Evolution of Dust Tracers
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Assumes distance of 140 pc and typical telescope properties.
C. Young and Evans, in prep.
Calculate Gas Temperature
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Use gas energetics code (Doty) with gas-dust collisions, cosmic rays, photoelectric
heating, gas cooling. Calculate TK( r, t).
C. Young and
J. Lee et al.
Calculate Abundances
Chemical code by E. Bergin
198 time steps of varying
length, depending on need.
Medium sized network with
80 species, 800 reactions.
Follows 512 gas parcels.
QuickTime™ and a
GIF decompressor
are needed to see this picture.
Includes freeze-out onto
grains and desorption due
to thermal, CR, photo
effects. No reactions on
grains. Assume binding
energy on silicates for this
case.
J. Lee et al. In prep
Calculate Line Profiles
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GIF decompressor
are needed to see this picture.
Line profiles calculated from
Monte Carlo plus virtual
telescope codes. Includes
collisional excitation, trapping.
Variations in density,
temperature, abundance,
velocity are included.
Assumes distance of 140 pc
and typical telescope
properties.
J. Lee et al. In prep
J. Lee et al. In prep
A Closer Look
Lines of HCN (J = 1–0).
Shown for four times.
Top plot with 50” resolution.
Bottom plot with 5” resolution.
ALMA will probe the desorption
wave.
J. Lee et al. In prep
Evidence for Infall
Good evidence in a few.
(e.g., Zhou et al. 1993)
Surveys indicate infall is
common at early stages.
Gregersen et al. 1997, 2000
Mardones et al. 1997
Observing Infall with ALMA



A key observation is to
observe the infalling gas in
redshifted absorption against
the background protostar
Very high spectral resolution
(<0.1 km/s) is required
High sensitivity to observe in
absorption against disk.
Low Mass Cores: Gross
Properties

Molecular cloud necessary, not sufficient



Centrally peaked density distribution

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
High density (n>104 cm–3)
Low turbulence
Power law slope ~ high mass
Fiducial density ~ 100 times lower
Complex chemistry, dynamics even in 1D



Evidence for infall seen, but hard to study
Outflow starts early, strong effect on lines
Rotation on small scales
Open Questions

Initial conditions



Cloud/core interaction
Trace conditions in core closer to center
Inward motions before point source?

Timescales for stages
 Establish existence and nature of infall



Envelope-Disk transition


Inverse P-Cygni profiles against disks
Chemo-dynamical studies
Inner flow in envelope
Outflow dynamics

Nature of interaction with ambient medium
Sub-stellar Objects
Brown dwarfs, free-floating planets, …
 BDs clearly exist, clearly have disks,
accretion,…
 How do these form?




Ejection from multiples, clusters
Formation like stars
Properties of disks
 Do they form in low-mass, dense envelopes?
 Low end of core mass function
Planet Formation
 Best
studied around isolated stars
 Origin and evolution of disk
 Gaps, rings, …
 Debris disks as tracers of planet formation
 Chemistry in disks

Evolution of dust, ices, gas
Planet Formation
SMM image of Vega shows dust peaks off
center from star (*). Fits a model with a
Neptune like planet clearing a gap.
This is with 15-m at 850 microns and 15”
resolution.
ALMA can do at higher resolution.
SMM image of Vega
JACH, Holland et al.
Model by Wyatt (2003), ApJ, 598, 1321
With higher resolution
Vega also observed by Wilner et al. (2003). Model of resonance with planet.
Predicts motion of dust
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are needed to see this picture.
Model and Animation by Marc
Kuchner
ALMA Resolution
Simulation Contains:
* 140 AU disk
* inner hole (3 AU)
* gap 6-8 AU
* forming giant planets at:
9, 22, 46 AU with local
over-densities
* ALMA with 2x over-density
* ALMA with 20%
under-density
* Each letter 4 AU wide,
35 AU high
Observed with 10 km array
At 140 pc, 1.3 mm
Observed
Model
L. G. Mundy
Chemistry of Planet-forming
Disks
LkCa15 with OVRO.
Trace the composition
changes with evolution.
QuickTime™ and a
TIFF (LZW) decompressor
are needed to see this picture.
Qi et al. 2004
ALMA will have resolution
and sensitivity to do this
kind of study in many
disks.
The Icy Component
Rich spectrum of ices:
CO2, H2O, CH3OH, OCN–
and others.
Can study ice evolution
in regions forming sunlike stars. Little
processing at T>50 K,
some evidence for lower
temperature processing.
Spitzer IRS plus Keck/NIRSPEC or VLT/ISAAC
Boogert et al. ApJS, submitted
Open Questions

How the disk initially forms
 Timescales for disk evolution
 How planets form in the disk


How unusual the solar system is


Core accretion or Gravitational Instability
Systems with giant planets out where ours are
Evolution of dust, ice, gas in disk

Building blocks for planets
Requirements
 Maximum

Image fidelity (gaps will be hard to see)
 Best

sensitivity
Especially for debris disks
 Flexible

Spatial resolution
correlator, receiver bands
Chemistry
In the ALMA era…
SOFIA 2005
Spitzer 2003
SMA, CARMA, eVLA,
LMT, GBT, APEX, ASTE,
JCMT, CSO, …
Herschel 2007
SAFIR ~2015
JWST 2011
AT-25 2012
Making the most of ALMA
 Complementary
Observatories
 User-friendly system


Low barriers to those from other wavelengths
Proposals, planning tools, reduction, analysis
 Scientific

support staff
Broad wavelength experience
 Financial
support tied to time
A Closer Look
A few abundance profiles at
t=100,000 yr.
Vertical offset for convenience
(except CO and HCN).
Big effect is CO desorption, which
affects most other species.
Secondary peaks related to
evaporation of other species.
J. Lee et al. In prep
Bolocam map of Ophiuchus
Bolocam map (1.2 mm)
of region in Spitzer
survey. Covers very
large area (> 10 sq.
deg.) compared to any
previous map.
Rms noise ~ 50 mJy,
with about half the
data.
K. Young et al. In prep
Early Results from Spitzer
 Based
on validation data (about 1%)
 Observed two small cores (IRAC/MIPS)


One (L1228) with a known infrared source
One (L1014) without
 Observed


a few IRS targets
B5 IRS
HH46/47 IRS (with ERO team)
A Typical Starless Core
L1014 distance ~ 200 pc, but
somewhat uncertain.
R-band image from DSS
A Surprise from Spitzer
Three Color Composite:
Blue = 3.6 microns
Green = 8.0 microns
Red = 24 microns
R-band image from DSS at
Lower left.
We see many stars through
the cloud not seen in R.
The central source is NOT
a background star.
L1014 is not “source-less”.
Larger size in red is PSF.
C. Young et al. ApJS, submitted
Source Peaks on mm Emission
Both long-wave
maps are 3-sigma
contours.
C. Young et al.
ApJS, submitted
Left: 8 micron on 1.2 mm MAMBO dust continuum emission (Kauffmann & Bertoldi)
Right: 24 micron on 850 micron SCUBA data (Visser et al. 2002)
Models
Model of SED for d = 200 pc.
Central object has very low
luminosity: 0.09 Lsun.
Requires BB plus disk
(red line) in an envelope.
M(envelope) about 2 Msun.
Cannot be a stellar-mass
object with significant
accretion. Probably substellar at this point.
Alternative model: more
distant (2.6 kpc) object lined
up by chance with peak of a
foreground core (dashed line)
C. Young et al. ApJS, submitted
Lessons from L1014
 “Starless”


Or may have substellar objects
1 out of 1 has a source (will soon have more)
 Very


cores may not be
low luminosity sources may exist
Must be low mass and low accretion
Caveat: possible background source
HH46/H47 Cloud
QuickTime™ and a
MPEG-4 Video decompressor
are needed to see this picture.
NASA/JPL-Caltech