Brown dwarf & star formation and the bottom of the IMF: a critical look - Gilles Chabrier
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Brown dwarf & star formation
and the bottom of the IMF:
a critical look
Gilles Chabrier
Field: Resolved objects IMF down to the HB limit
Salpeter
Kroupa ’01
Chabrier ’03
Chabrier ’05
Field: Unresolved systems IMF down to the HB limit
local
HST
Chabrier ’03 ApJL
T Y Reid et al. ’04
NBD/N* = 1/4 - 1/3 WISE (now !)~1/5
Chabrier ’05
Young clusters and SFR: system IMF across the HB limit
Stars Brown dwarfs
field system IMF
Young clusters and SFR: system IMF across the HB limit
field system IMF
0.01 Msol 5 Msol 0.01 Msol 5 Msol
Andersen et al. 2008
Combined analysis of 7 SFR’s
star/BD ratio in the 7 clusters consistent w/ the same underlying IMF
Excess of very-low-mass BDs ?
Sumi et al. ’11
5 events !
even if real, possibility of planets kicked out by other companions (Vera et al. ’09, Janson et al.’12)
Multiplicity frequency in the stellar and substellar regime
Thies & Kroupa 2008 :
discontinuity between the stellar and the BD IMF
Kroupa ’01
dn
/m ↵ ↵ = 0.3 1
dmConclusion 1 • IMF similar between the field and young clusters/SFR’s (within some scatter 0.01 Msol ruled out !! (Message for the IAU !!)
Brown dwarf
formation
{
-
-
Disk fragmentation
Vorobyov & Basu; Stamatellos & Whitworth; Bate
Accretion - ejection
Reipurth & Clarke; Bate, Bonnell
- Gravo-turbulent fragmentation
Padoan & Nordlund; Hennebelle & Chabrier; HopkinsDisk Fragmentation
- Disk massive and cool enough to become gravitationally unstable,
according to the Toomre and Gammie criteria
cS Prot
Q=µ=1000 (hydro) µ=50
µ=50 µ=20
µ=20
300 AU
Low magnetic fields allow disk formation
but
the disk is stabilized and does not fragment
Machida et al. ’04; Vorobyov & Basu ’06; Hennebelle & Teyssier ’07; Price & Bate ’07;
Hennebelle & Ciardi ’09; Commerçon et al. ’10µ=5 µ=2 µ=1.25
With stronger fields, no centrifugally supported disk
Hennebelle & driven
forms, making rotationally Teyssier 2007
fragmentation even
more problematic
Confirmed by calculations with complete resistive MHD (Masson et al. in prep.)Comparison of the PdBI maps with MHD simulations
Maury et al. 2010; see also Stamatellos, Maury et al. ’11
Hydrodynamical simulations
produce too much extended
(+ multiple) structures if
compared to the observations.
Obs
MHD simulations ?
Taurus Perseus
Hennebelle & Teyssier (2008)
MHD simulations :
produce PdB-A
synthetic images with No B
typical FWHM ~ 0.2’’ - 0.6’’
Similar to Class 0 PdB-A
sources observed !
w/ B
need B to produce compact,
single PdB-A sources.
Massive disks prone to fragmentation not observed !Rodriguez et al. ’05
Isolated, massive, compact (Rout~20 AU) disk around class 0
consistent w/ MHD disks !Constraints on BD formation/ejection by disk instability: magnetic field
No B (pure hydro): no outflow -> too massive disks : inconsistent w/ observations (Maury et al. ’10)
B : outflow + magnetic breaking (Price & Bate ’09, Commerçon et al. ’10; Machida et al. ’11, Masson )
Constraints on disk instability from class-0 disks
Class-0 lasts ~ 105 yr (Evans et al., Enoch et al.)
=> most Class-0 objects should have massive disks: not observed !
Low-mass binaries with BD-mass companions
e.g. BD LHS6343c (Johnson et al. 2011) (other examples: Joergens 2008, Irwin et al., 2010, Bachelet et al., 2012)
MA=0.37 Msol
MC=0.063 Msol MDISK ≃ 0.04 Msol
not enough mass in the disk to form LHS6343c formation by disk fragment’n
wide BD binaries : (eg 2MASS J12583501/J12583798 6700 AU (Radigan et al. ’09)
Constraints on disk instability from direct imaging
Lafrenière et al. (2010), Johnson et al. (2010), Janson et al. (2011, 2012)
archival data from the Gemini Deep Planet Survey:
statistics of companions around 85 ABFGKM stars in disks from 5 to 700AU
companion frequencyEjection, competitive
accretion
- Collapse of a cloud -> starts forming small N-body clusters of
small (~10-3 Msol) gravitationnally bound entities
- then accrete mass, Ṁ / M 2(Bondi-Hoyle) or Ṁ / M 1.5 (tides)
- dynamical interactions responsible for merging (-> high-mass stars)
or ejection (LMS/BD)
• dynamical interactions are essential to build the IMF
• all stars form in clusters
• the IMF is determined by the gas-to-star conversion
(no correspondence with the CMF)Bate ’1É: fragmentation of a cloud w/ radiatove feedback
L=0.4 pc ; n ~ 105 pc-3 ; Mach# =14
Falgarone et al. ’00
~ 50xdenser and 5x more turbulent than (observed) Larson’s relations ! => overfavor dynamical int’ns !!
idem NGC1333 : N*+BD (simus) = 191 Msol N*+BD (observed) ~ 50 Msol (Scholz et al. ’12)
+ no turbulent compressive mode (Federath et al; Hennebelle & Chabrier ’09)
uniform initial density (Girichidis et al.)
no B
stellarfeedback underestimated
IMF set up at the very begining of the simulation
brings support to the gravo-turbulent scenario !Bayo et al. ’11: radial dist’n of stars and BDs in Lambda-Ori
stars
* BD’s Members confirmed
spectroscopically !
stars
BD’s
No difference in the distributions of stellar and substellar membersMarsh et al. ’10: radial dist’n of stars and BDs in rho-Oph
Constraints on ejection scenario: 3) timescale argument
André et al. ’07, Walsh et al. ’07, Muench et al. ’07, Gutermuth et al. ’09, Bressert et al. ’10
~ 0.4 km/s
⌧collapse ⌧ ⌧collisions
suggest little dynamical evolution during prestellar core formationOther issues w/ competitive accretion / ejection scenario
- how do you preserve the correspondence between CMF and IMF ?
answer : deny it !
- how do you produce both the correct MF of unbound (CO) clumps and bound cores ?
- how do you form BDs in low-density environments (eg Taurus) ?
- how do you preserve wide BD binaries ?
(eg 2MASS J12583501/J12583798 6700 AU (Radigan et al. ’09)
- how do you preserve disks, outflows in ejected embryos (proto-BDs) ?
- how do you end up having a universal IMF ?
(Bate 2009 : feedback from LMS leads to a narrow range of Jeans mass in the irradiated gas )
- how do you form BD’s in isolation (eg Fu Tau A,B Luhman et al. ’09) ?
- all stars must form in clusters : contradicts obs (Bressert ‘10, ‘12)
+ formation of isolated massive stars (Bontemps et al. ’10, Bestenlehner et al. ’11, Bressert ’12)
+ formation of isolated proto-BDs (Palau et al. 12) and pre-BD cores (André et al. ’12)Gravoturbulent
fragmentation
- Large-scale turbulence (density PDF lognormal) sets up a lognormal
distribution of overdense regions at all scales in the cloud
-> leads to the CO clump dist’n (can be filamentary)
- Virial: regions M(R) with |Egrav (R)| > (Eth+Erms(R)+Emag) collapse
(introduces a scale dependence)
-> leads to the core mass function (CMF)
- magneto-centrifugally outflows expel ~30-50% of the mass
-> leads to the star mass function (IMF)
• turbulence sets up the density fluctuations at the very early
stages of star formation : «turbulent Jeans mass»
• the IMF derives from the CMF and is imprinted in the very cloud
conditions (, T, Mach)Turbulent Jeans mass
Hennebelle & Chabrier ’08; Chabrier & Hennebelle ’11
2
turb Vrms V02 R 2⌘ confirmed by simulations:
MJ (R) = R= ( ) R
3G 3G 1 pc Schmidt et al. ’10
Not a static support ! But a statistical “filter” due to rms velocity dispersion (see CH11)
Timescale problem
3 p
Thermal Jeans mass : M ⇠ MJ / CS / ⇢ ) ⌧f f / M
turb p
Turbulent Jeans mass : MJ / Vrms / ⇢ ) ⌧f f / M 1/4
3
Much weaker dependence of the core free-fall timescale upon the mass when the impact of
turbulence is accounted for. Confirmed by the complete time-dependent theory (HC12, HC13).
Consistent w/ obervations: André, Basu, Inutsuka ’09 :
ka ’09Comparisons with numerical simulations
Jappsen et al. ’05 Schmidt et al. ’10
n̄crit ⇥ 4.3 106 cm 3
thermal
n̄crit ⇥ 4.3 107 cm 3
Hennebelle & Chabrier, ’09 turbulent
No free parameterChabrier
system IMF
HC analytical M 7
IMF
M⇠5
Hennebelle & Chabrier ’09
29log ⌃˙
40
Lcloud 10
(pc) 1
4
* Lada et al. ’10
Evans et al. ’09
+ Heiderman et al. ’10 Hennebelle & Chabrier 2011, 2013
Gutermuth et al. ’11gravoturbulent fragmentation
• supported by IMF (at least partly) determined by the prestellar Core MF
(André et al. 2010 HERSCHEL, Konyves et al. ‘12,10)
• can lead to binary properties consistent w/ observations (Jumper & Fisher ’12)
emerging observations of fragmentation at the Class-0 phase
• can explain both the unbound clumps and bound cores MF’s (HC08,09)
• emerging observations of isolated proto-BD and pre-BD cores
VeLLO L1148–IRS (Kauffmann et al. ’11); IRAS 16253-2429 (Wiseman et al.)
J042118 + J041757 1-5 MJup (Palau et al. ’12)
Oph B-11 30 MJupYoung BDs vs young star properties
same radial velocity dispersion see also. Luhman et al. ’07, ’12
same spatial distribution in young clusters
consistent with the same IMF
wide binary BDs
accretion + disk signature (large blue/UV excess, large asymetric emission lines, Hα
BD and star formation : common mechanism
ejection, disk frag’n, photoevaporation might play some rome but
are not essential in making it possible for BDs to formHow to produce binaries ?
Let us consider an m=2 perturbation with an amplitude of 50%
µ=20 µ=2 µ=1.25
If the perturbation has a large amplitude, the fragments develop
independently of rotation, the field is not amplified and does not prevent the
fragmentation except if it is initially strong (see also Price & Bate 2007).
But the fragments are initially strongly seeded….
Need to explore more realistic initial conditions.
Hennebelle & Teyssier 2008 (see also Price & Bate 2007)Observations and statistical properties of multiple Systems
About 50% of the stars are binaries (Duquenoy & Mayor 91)
There are evidences that the fragmentation occurs very early
when the star is still accreting.
Evidences for fragmentation Multiple embedded
in the Class0 objects IRAS4A Young stellar objects
(Observations realised with the in ρOph and Serpens
BIMA interferometer)
NGC1333 IRAS4 - BIMA - 2.7 mm
5’’ 3’’
A2
A1 0.5’’ 1’’
Looney, Mundy, Welch (2000) Duchêne et al. 2003Conclusion 2 • Can star/BD’s form dominantly by G.I. or competitive accretion/ejection ? - not supported - so far - by observations - face many issues (among the main ones : CO-clump MF ? CMF ? universality ?) - IMF should vary appreciably w/ environment - need more realistic intial/physical conditions - can produce some objects (eg Delorme et al. ’13) • Gravo-turbulent fragmentation - relies on a sound theory (same theory yields CO-clump MF and bound core CMF) - emerging population of isolated pre- and proto-BD cores - the CMF-to-IMFpart still to be fully understood (magneto-centrifugally driven outflows (Matzner & McKee) - binary formation at the 1st and 2nd core stages still to be fully understood • Low-mass part predicted to vary with Mach number, e.g. in extreme environments -> naturally leads to more low-mass stars - thus higher M/L - at large Machs expected in bursts, mergers (ETG’s)
Can the IMF vary in extreme environments (eg ETG’s) ?
Salpeter
Ṁ & 10 ⇥ ṀM W
Mach=50
(M/L) ⇠ 1.8 (M/L)M W HC Mach=5
Mach=40
Chabrier
Mach=20
ṀM W ⇡ 2M /yr
Chabrier & Hennebelle, in
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