Download AST4930 Star and Planet Formation

Stellar clusters and associations
• Stars do not generally form in isolation.
• The vast majority (90%) of stars form in aggregates with 100 or more members, for a total
stellar mass exceeding 50M
NGC1333
LH 95
IC 348
Trapezium cluster (Orion Nebula Cluster)
NGC 3603
R136 (30 Doradus)
Stellar clusters and associations
• Stars form in groupings (clusters and associations):
1.
Stellar aggregates contain young stars
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These stars are young because of their photospheric characteristics:
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Brighter than main sequence – larger radii - still contracting
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Show photospheric signatures of infalling gas – H recombination lines
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Show evidence of warm dust surrounding them – IR excess from heated dusty disks
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Show prominent X-ray emission – typical of young convective stars
These stars must have formed recently because they are surrounded by gas
2. Massive stars are found in stellar aggregates
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Massive stars are short lived (few Myr for M>20Msun) – they must have formed recently
Therefore the aggregates themselves must be young
• Not all these groupings survive as bound clusters
• Early dissolution of the stellar population into the field
``infant mortality’’, or slow expansion and stellar loss.
• This can be concluded from the comparison of the
current star formation rate (stars per unit of time per
unit of volume) with the number of stars in open
clusters per interval of age.
• 90% of the current star local star formation output
feeds the galactic field population
Trapezium cluster (Orion Nebula Cluster)
Pleiades (100Myr)
Embedded clusters
• At early stages 𝑡𝑡~ < 1 𝑀𝑀𝑀𝑀𝑀𝑀 young clusters are deeply embedded within their
parental molecular cloud.
• The high dust extinction produced by the high column density of dust limits the
detection of their stellar population at optical wavelengths.
Absolute magnitude:
+
Sun:
𝑀𝑀𝑉𝑉 ~5 𝑚𝑚𝑚𝑚𝑚𝑚
Distance modulus:
Dust extinction
𝜇𝜇 = 5 log 𝑑𝑑 − 5
+
10pc
𝜇𝜇 = 0
100pc
𝜇𝜇 = 5
1kpc
𝜇𝜇 = 10
10kpc
𝜇𝜇 = 10
Low:
𝐴𝐴𝑉𝑉 < 5
Intermediate 𝐴𝐴𝑉𝑉 ~5 − 15
High
𝐴𝐴𝑉𝑉 ~20 − 100
Trapezium cluster (Orion Nebula Cluster)
Approximate limit of optical telescopes: V=20-25 mag (ground based); 28-30 (HST)
Embedded clusters
• Dust extinction is significantly lower at longer
wavelengths.
ONC optical
Trapezium cluster (Orion Nebula Cluster)
• Note: at large distances within the disk of the MW, the buildup
of foreground extinction fro the disk itself becomes important
too.
ONC JH
Detection of embedded clusters, NIR surveys
Near infrared photometric surveys enable the discovery of embedded young clusters.
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an excess of density of sources in a given location may be cluster candidate
color-magnitude and color-color diagram, compared to adjacent fields, can identify the young
population and help estimating its properties (number of stars, average age, …)
AV
direction
ONC optical
Trapezium cluster (Orion Nebula Cluster)
IR
excess
sources
ONC JH
IR excess young stars
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Some young stars exhibit an excess of flux at long NIR wavelengths (𝜆𝜆 > 2𝜇𝜇𝜇𝜇), as their NIR colors cannot be
reproduced invoking only dust extinction.
Once the observed multi-wavelength fluxes are corrected for extinction (e.g., based on multi-band photometry
at shorter wavelengths, or e.g., knowing the spectral type of the source from spectroscopic measurements, the
spectral energy distribution (SED) of the excess is isolated.
The excess has a circumstellar origin, and represents the emission of the dust in the inner part of the
circumstellar disks, heated by the central star (𝑇𝑇𝑑𝑑𝑑𝑑𝑑𝑑𝑑𝑑 < 𝑇𝑇𝑠𝑠𝑠𝑠𝑠𝑠𝑠𝑠 )
Since circumstellar disks have short lifetimes (<10Myr) as their material is either photoevaporated, accreted, or
condenses in planets, circumstellar disk presence from IR excess is an indicator of youth.
𝝆𝝆 Ophiuchi dark cloud complex
ONC optical
Trapezium cluster (Orion Nebula Cluster)
Upper Sco OB association
IR spectral index 𝛼𝛼𝐼𝐼𝐼𝐼 =
ONC JH
𝑑𝑑 log 𝜆𝜆𝐹𝐹𝜆𝜆
𝑑𝑑 log 𝜆𝜆
IR excess young stars
Class I
𝛼𝛼2−10𝜇𝜇𝜇𝜇 > 0
Protostar: emission originates from
accreting disk + envelope
Class II
−1.5 < 𝛼𝛼2−10𝜇𝜇𝜇𝜇 < 0
Pre-main sequence star: IR emission
from dusty circumstellar disk
ONC optical
Trapezium cluster (Orion Nebula Cluster)
Class III
𝛼𝛼2−10𝜇𝜇𝜇𝜇 < −1.5
Older pre-main sequence star. Disk
partially cleared inside-out
ONC JH
T and OB-associations
• Over a time of ~1Myr embedded clusters start to be optically revealed. We can distinguish
different categories.
• T-associations:
• Systems that were born in low-density dark clouds that did not
contain massive stars.
• Loose distributions of young pre-main sequence stars
• Typical example: the Taurus star forming region.
• OB-associations:
• The unbound, expanding grouping of newly formed stars,
containing massive, bright OB stars.
• The high effective temperature of OB stars (𝑇𝑇𝑒𝑒𝑒𝑒𝑒𝑒 > 10,000𝐾𝐾)
ionizes the remaining gas creating HII regions
• More massive than T associations.
• Actual clusters:
• When the stellar population remains bound (𝐸𝐸 = 𝐾𝐾𝐾𝐾 + 𝑃𝑃𝑃𝑃 < 0)
after gas removal.
• Note that a fraction (large or small) of stars may escape the stellar
system in its early phases of evolution even when a bound clusters
survives afterwards.
Taurus association
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Distance: 140 pc
Low density of stars: 1-10 stars/pc3 spans over 15 degrees on the plane of the sky
Absence of massive stars.
Many stars are distributed in relatively compact clumps which follow the density of gas
~1 Myr old (with age spread largely unknown)
Luhman (2009)
Memberships of young associations
Especially (but not only) for the case of loose T associations, the isolation of the young members from field
(background and foreground) galactic field stars might be challenging. There are several memberships indicators that
can be used; they are in many ways complementary, as each of them alone does not guarantee to accurately identify
(if at all) the young nature of PMS stars. Limiting to low-mass PMS stars, these are:
1. Flux excesses and other spectral characteristics
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Disk excess: as it traces the presence of a circumstellar disk. Note that for more evolved class III sources, the signal might be
weak, requiring observations at long wavelengths (MIR, FIR, sub mm)
Accretion signatures: as matter falls from the inner disk on to the star, the loss of gravitational potential energy is converted
into luminosity, with a typical spectrum including UV excess, optical continuum, and Hydrogen recombination lines. The Hα
line at λ=6563 Angstrom excess, in particular, is most commonly used since practical to measure, and bright. A standard
classification is as follows:
Classical T-Tauri stars: E. W. H𝛼𝛼 > 20 𝐴𝐴
Weak-line T-Tauri stars: 0 < E. W. H𝛼𝛼 < 20 𝐴𝐴
Other accretion indicators are H and K-lines of Ca II at 3968 and 3934 Angstrom
X-rays: young low mass stars are generally X-ray emitters, originating in plasma explosively heated and confined in magnetic
loops following magnetic reconnection events. This is an analogous of solar flares, and is facilitated by the turbulent and highly
convective interiors of low mass PMS stars
2. Kinematic common motions
The velocity dispersion of the members of a young star forming
region is generally of a few km/s, whereas the dispersion of field stars
may be more than one order of magnitude higher. Proper motions
and radial velocity surveys help to isolate the candidate members
that share the same bulk motion.
proper motions of
the members in
Taurus
OB associations
An aggregate of young stars containing early (OB) type stars.
B-type stars: 𝑇𝑇𝑒𝑒𝑒𝑒𝑒𝑒 > 10,000 𝐾𝐾; 𝑀𝑀 ≳ 4𝑀𝑀⨀
O-type stars: 𝑇𝑇𝑒𝑒𝑒𝑒𝑒𝑒 > 30,000 𝐾𝐾; 𝑀𝑀 ≳ 20𝑀𝑀⨀
The large mass of these stars requires the stellar population to be
more numerous than T-associations; thus OB associations originate
from more massive clouds.
The UV radiation of early type stars is ionizes the surrounding gas,
which then emits Hydrogen lines as it recombines forming an HII
region.
LH 95 (Large Magellanic Cloud)
Note that massive star early evolution is fast
enough that these objects are already on the
main sequence as soon as the stellar system is
non embedded.
Cygnus OB2 – NGC6914
Spectroscopic analysis can identify if an early type
star is in main-sequence or already leaving the
main sequence. In the first case, it follows a given
temperature-luminosity relation, which can be
exploited to determine their distance (but
reddening must be corrected for).
Nearby young associations and Gould’s belt
T associations
Many close star forming regions, as well as many of the
nearby bright stars, lie on a ring-like structure called Gould’s
belt. It has a radius of ~300 pc, and is tilted by 20° with
respect to the Galactic plane.
Kinematic measurements show that the belt is expanding
and slightly rotating. Age ~ 25Myr
OB associations
Expansion and subgrouping of complexes of star forming regions
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Young star forming regions are often part of larger complexes including subgroups. An example is
shown below: the Scorpius-Centaurus association.
These subgroups generally formed as part of a joint structure of clouds, or molecular cloud complex.
However, different subgroups may have formed at different times Δ𝑡𝑡 ≲ 10 Myr.
As young association tend to expand over time, as a result of excess kinetic vs. potential energy, the
concentration of an association is a (very rough) indicator of its age.
Proper motions in this case can be used to track back the moment in the past when the stellar
population was at its maximum density.
Upper Scorpius
The Orion complex
Trapezium cluster
Orion Nebula
Blue: optical
Red: HII
Green: H2
Extinction map
(Lombardi+2007)
Blue: optical
Red: HII
Green: H2
Orion OB association groups
Orion OB Ia: ~12 Myr
Orion OB Ib: ~8 Myr
Orion OB Ic: ~3 − 6 Myr
Orion OB Id: ~2 Myr - Orion Nebula
Blue: optical
Red: HII
Green: H2
Megeath+ (2012)
Young low mass stars in the Orion
complex, part of the same population of
the youngest Orion OB Id population, are
identified through their NIR and MIR
excess emission from circumstellar
material. They are confined in filamentary
structures which follow the Orion A and B
clouds.
Subgrouping on smaller scales in the Orion complex
Alves+ (2012)
Analysis of the properties of low-mass pre-main sequence stars: lower AV and higher fraction of
Class III sources vs Class II to the south of the Orion Nebula may indicate an older population.
Ages of young stellar associations
1.
Embeddedness (high AV = young).
2.
IR excess (temporal evolution Class I → Class II → Class III).
3.
4.
Proper motions to track back the expansion of a stellar system
Presence of unevolved massive stars sets upper limits to cluster age
5.
Turn off age: the
luminosity (and mass)
of the turn off point
from the main
sequence decreases
with time
6.
Pre-main sequence
H-R diagram vs
evolutionary models
𝑡𝑡𝑀𝑀𝑀𝑀 ∝ 𝑀𝑀−2.5
Pre-main sequence ages
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The long contraction time needed for PMS stars to reach the MS can be used as a
clock to measure stellar and cluster ages.
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BUT: 1) theoretical models for PMS star are more uncertain than those for MS stars
2) determination of 𝑇𝑇𝑒𝑒𝑒𝑒𝑒𝑒 and log 𝐿𝐿 for PMS stars is challenging
Embedded massive stars
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Massive stars at very young ages can be still deeply embedded
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Their UV radiation can create Ultra Compact HII regions
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Smaller than HII regions (𝑟𝑟 < 0.3𝑝𝑝𝑝𝑝)
Radiation absorbed and re-emitted at longer wavelengths (MIR, FIR, Radio)
Runaway massive stars
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25% of O-type stars are not spatially associated with young clusters or clouds.
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Many are located farther from the Galactic plane, and show velocities
𝑣𝑣 = 50 − 100 km/s
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Their velocity originates from:
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Multi-body stellar encounters
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As ex-companions of supernovae
Open clusters
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The remain of OB associations that have not undergone complete disruption. Size: 2-10pc.
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Despite their survival (ages 30-1000Myr) not all of them are bound.
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No or little molecular gas left. No active star formation. The members are coeval.
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~1200 known systems (but sample is limited by galactic disk extinction), distributed.
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Precious laboratories of stellar evolution: simple stellar populations over a range of ages. Accurate photometric
measurements allow to test theoretical models of stellar evolution and spectra of stars.
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Draper (1930): discovery of dust extinction: smaller (= farther away) young clusters
appear dimmer.
Pleiades (130 Myr)
NGC4755 (16Myr)
Hyades (650Myr)
Hyades
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The proper motion are aligned towards a convergence point, where the open cluster appear to shrink towards.
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This indicates that the cluster is receding. The determination of the convergence point, together with measured
proper motions and radial velocities can be used to determine the distance of the cluster:
θ
𝑉𝑉𝜏𝜏
𝑉𝑉𝑟𝑟
= tan 𝜃𝜃 ; 𝑉𝑉𝜏𝜏 = 𝜇𝜇 � 𝑑𝑑
𝑉𝑉𝑟𝑟 tan 𝜃𝜃
𝑑𝑑 =
𝜇𝜇
The initial mass function
• The IMF is the distribution of the initial masses of stars.
𝑑𝑑𝑑𝑑
𝜉𝜉 𝑚𝑚 =
𝑑𝑑𝑑𝑑
1.
Provides useful clues on the processes that lead to the conversion of molecular
clouds into stars.
2.
Stars of different mass have different lifetimes; the IMF dictates the fate of stellar
systems up to entire galaxies.
3.
Massive stars provide energy feedback and chemical enrichment. How many does
a galaxy form?
4.
Very low mass stars and brown dwarfs are ``invisible’’ mass at large distances.
How many are there?
5.
Its form, universality or variations, is fundamental to interpret observations at
large distances (e.g., galaxies)
The initial mass function
• The IMF is the distribution of the initial masses of stars.
• Empirical distribution, based on counting stars of different mass.
• The stellar IMF is commonly fitted by a power-law:
𝑑𝑑𝑑𝑑
𝜉𝜉 𝑚𝑚 =
∝ 𝑚𝑚−𝛼𝛼
𝑑𝑑𝑑𝑑
Φ log 𝑚𝑚 =
𝑑𝑑𝑑𝑑
∝ 𝑚𝑚−Γ
𝑑𝑑 log 𝑚𝑚
𝛼𝛼 = Γ + 1
•
Salpeter (1995): the stellar IMF, above 1Msun is well fitted by a single power law
𝛼𝛼 = 2.35
• Extension of the IMF in the low mass stellar range and brown dwarf regime
(M<0.08Msun) shows flattening
•
Kroupa (2001,2001): the galactic field IMF
multi-part power law
•
Chabrier (2005) IMF:
lognormal for low-mass stars and brown dwarfs
From luminosity function to IMF
1.
𝑀𝑀𝑉𝑉∗
General luminosity function:
Distribution of present-day number of stars per
absolute magnitude bin (per unit of volume, 𝑝𝑝𝑝𝑝 −3 )
𝑑𝑑𝑑𝑑
now
𝑑𝑑𝑀𝑀𝑉𝑉
Needs special care when derived: incompleteness, selection
effects, and stars of different mass (=luminosity) extend to
different scale height above the galactic midplane.
Φ 𝑀𝑀𝑉𝑉 =
2.
Bright,
𝒎𝒎 ≳ 𝟎𝟎. 𝟖𝟖𝑴𝑴⨀
The first formed
already died
Initial luminosity function:
Φ 𝑀𝑀𝑉𝑉 =
𝑡𝑡𝑔𝑔𝑔𝑔𝑔𝑔
�
Ψ 𝑀𝑀𝑉𝑉 =
𝑡𝑡𝑔𝑔𝑎𝑎𝑎𝑎−𝑡𝑡𝑀𝑀𝑆𝑆 (𝑀𝑀𝑉𝑉 )
𝑡𝑡𝑔𝑔𝑔𝑔𝑔𝑔
�
0
𝑑𝑑𝑑𝑑
at formation
𝑑𝑑𝑀𝑀𝑉𝑉
𝑑𝑑𝑑𝑑 𝑚𝑚̇ 𝑡𝑡 Ψ(MV )/2H(M𝑉𝑉 )
𝑑𝑑𝑑𝑑 𝑚𝑚̇ 𝑡𝑡 Ψ(MV )/2H(M𝑉𝑉 )
star formation rate
scale height
𝑖𝑖𝑖𝑖 𝑀𝑀𝑉𝑉 < 𝑀𝑀𝑉𝑉∗
𝑖𝑖𝑖𝑖 𝑀𝑀𝑉𝑉 >
𝑀𝑀𝑉𝑉∗
Faint,
𝒎𝒎 ≲ 𝟎𝟎. 𝟖𝟖𝑴𝑴⨀
Never died
m
0
t
1010
From luminosity function to IMF
𝜉𝜉 𝑚𝑚 =
𝑑𝑑𝑑𝑑
𝑑𝑑𝑑𝑑
=
𝑑𝑑𝑑𝑑 𝑑𝑑𝑀𝑀𝑉𝑉
𝑑𝑑𝑀𝑀𝑉𝑉 𝑑𝑑𝑑𝑑
= Φ(𝑀𝑀𝑉𝑉 )
𝑑𝑑𝑀𝑀𝑉𝑉
𝑑𝑑𝑑𝑑
Mass-magnitude relation for MS stars.
•
Theoretical: from stellar evolutionary simulations
and stellar atmosphere modeling.
•
Empirical: from multiple stellar systems with
orbital parameters well determined.
Unresolved stellar multiplicity affects the slope of the measured IMF. This correction is sensitive to
the actual fraction of stars in multiple systems, its variation with stellar mass (e.g., massive stars
more likely to be in multiple systems than low-mass stars), and the companion mass distribution.
“stellar IMF” – corrected for multiplicity
“system IMF” – uncorrected
IMF from young clusters
1.
Advantages:
- no evolutionary effects.
- one single distance for all members of one cluster.
- VLMS and BDs are brighter when they are young.
2.
Disadvantages:
- estimated masses are more uncertain.
(both from observational difficulties and theoretical uncertainties in stellar
evolutionary models)
Orion Nebula cluster, Da Rio+ 2012
Universality of the IMF
• There is no evidence of strong systematic variations of the IMF with initial conditions
after the first generation of stars in the universe.
• The IMF appears universal over several order of magnitude of cluster density, mass,
metallicity.
• Local variations of the IFM are statistical or exceptional.
Bastian+2010