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Dust processing in debris disks
Mark Wyatt
Institute of Astronomy, University of Cambridge
What is a debris disk? Flux density (Jy)
Infrared emission of nearby main sequence
stars brighter than photosphere: e.g.,
Fomalhaut has 70K excess
Imaging shows emission
from dust in a ~130AU ring
(Acke et al. 2012)
Wavelength (µm)
Component of planetary system Debris disks are
components of
planetary systems,
e.g., the Solar
System’s debris
disk is the asteroid
and Kuiper belts
They come from
the protoplanetary
disk and are
directly indicative
of planetary
architecture
Kuiper belt Descendant of protoplanetary disk Age Protoplanetary disk Debris disk <10Myr 10Myr – 10Gyr Op7cal depth OpEcally thick The PPD vs DD figure OpEcally thin Dust mass >10Mearth <1Mearth Gas mass ~100x dust mass None (usually) Structure Dust from 0.1(?)-­‐100AU Confined to ~30AU ring Dust origin Primordial? Secondary (short lifeEme) Debris disk primer
Simplest model for a debris
disk has planetesimals
orbiting the star confined to a
belt
Collisions grind planetesimals into
smaller and smaller fragments
resulting in collisional cascade with a
size distribution:
n(D) ∝ D-3.5
r
What we see from this belt is the
result of the interplay between
collisions and radiation forces
n(D)
Dmin
Cross-sectional
area
Dmax
Mass
Diameter, D
Key questions about dust processing
How big are the biggest planetesimals?
What stirred them to destructive collisions, and at
what level of stirring is required?
What processing occurs in the collisional cascade?
What physics is particularly relevant for the dust?
What do we know about dust composition?
Dust observables
Thermal emission: dust is heated by the star to a temperature (T) that
depends on dust size (D) and distance from the star (r), and
composition, then spectrum is integral over r and D of QabsBν(T)
Small dust is hotter than large dust, so
spectrum constrains size distribution
Spectral features (e.g., 10μm silicate
feature) constrain dust size (must be
smaller than wavelength) and composition
Scattered light and polarisation: composition, size, and geometry
Example of Fomalhaut
The 130AU location of dust in Fomalhaut combined with the emission
spectrum shows the size distribution is that expected in a collisional
cascade: n(D) ∝ D-3.5
Wyatt & Dent (2002)
Sub-mm slope
The size
distribution slope is
also well
constrained from
sub-mm spectral
slope to 3.48±0.14
(Ricci et al. 2012)
Need more long
wavelength
observations of
debris disks, as
these probe largest
grains
See Woitke and Schüppler talks
Contribution of
different grain
sizes to fluxes
in different
wavebands for
a collisional
size
distribution
shows <5% of
850µm flux
comes from
>20cm dust
Fraction of flux per log(D)
What dust sizes are we seeing? 24µm
60µm
100µm
450µm
850µm
What’s feeding the cascade?
The observable portion of
the collisional cascade
extends up to 20cm
Without larger objects this
would disappear in ~1 Myr;
200Myr age of Fomalhaut
implies planetesimals ~4km
feed the cascade
Wyatt & Dent (2002)
We cannot tell if >4km planetesimals exist
Dust mass evoluEon Convert sub-mm fluxes to mass in mm-cm-sized dust using opacity 1.7cm2/g:
Mdust = F850 d2 / Bν(T) κν
Debris disks have <0.1
times less mass in mmsized dust than
protoplanetary disks,
and there is a sharp
transition at ~10Myr
But the mass needed in
km-sized objects to
sustain debris disks
means that total solid
masses may be similar
Including
planetesimals
Panić et al. (in press)
Collisional processing Size distribution means that
most collisions are with much
smaller objects
Such cratering collisions don’t
usually affect mass flow
Dispersal threshold sets
specific incident energy
required to remove half the
mass, giving the size of object
required
Collisions below this threshold
in gravity regime result in
rubble piles
Compare with talks on Tuesday afternoon
Model size distribution evolution
Solve dmk/dt = (dmk/dt)gain - (dmk/dt)loss
Rippple from small-size cut-off
(Thebault & Augereau 2001)
Change in slope from strength-gravity scaling
(O’Brien & Greenberg 2003)
Turn-over at
late times
when P-R
drag becomes
important
(Wyatt,
Clarke &
Booth 2011)
Cross-­‐sec7onal area, AU2 Ripple from that
change in slope
(Durda et al. 1998)
Slower evolution in
mass than in area
(Lohne et al. 2008)
Diameter, m
Simpler numerical model
Rather than solve:
dmk/dt = (dmk/dt)gain - (dmk/dt)loss
Assume steady state and solve: (dmk/dt)gain = (dmk/dt)loss
This reproduces
all features of
full simulations,
but at high
resolution and
in seconds
(Wyatt, Clarke &
Booth 2011)
Changing collision velocity affects the wiggles, but not overall shape
Are the wiggles observable?
The size distribution constrains the temperature distribution, so what
does disk temperature tell us about smallest dust sizes?
Disk imaging with
Herschel shows
there is a range of
temperatures from
that expected for
black body dust up
to several times that
(Booth et al. 2013)
Is the range from
wiggles or dust
composition, e.g.,
porosity (Kirchschlager &
Wolf 2013)
Radiation Pressure
Radiation pressure truncates the
collisional cascade at small particles:
β = Frad/Fgrav ≈ (0.4/D)(L*/M*)
β>0.5 blown out on
hyperbolic orbits
0.1<β<0.5 put on eccentric orbits
In addition to halo of unbound
grains, bound particles close to
blow-out limit extend far
beyond their birth ring (Wyatt
1999, Krivov et al. 2000, Thebault &
Augereau 2001, Strubbe & Chiang 2006)
Evidence for halos Quick comparison
shows that sub-mm
emission from
Fomalhaut
(orange) is well
aligned with optical
(blue) (Boley et al. 2012)
However, at 16AU the sub-mm ring is much
narrower than scattered light, since it sees
mm-cm-sized dust whereas optical traces
dust affected by radiation pressure
PolarisaEon of halos Increasing
polarisation with
distance seen in
AU Mic’s halo is
consistent with
smaller grains
further from star
(Graham et al. 2007)
High degree of
polarisation is
inconsistent with
compact grains
Vega’s disk also looks
more extended at
shorter wavelengths that
trace the dust halo
(Holland et al. 1998; Su et al.
2005; Sibthorpe et al. 2010)
Surface brightness
Halos also seen in thermal emission
70μm
850μm
24μm
Radius (AU)
It was claimed that the implied temperature requires dust below the
blow-out limit, that must have been produced very recently (Su et al. 2005)
But a steady state solution was also found that explains the halo as
bound grains (Muller et al. 2010)
Need more halo characterisation and modelling
Size distribution at small sizes
Cratering collisions
become important for
the size distribution at
small collision velocities
(Thebault & Wu 2008)
Two observables:
•  sharp outer disk
edge (e.g., HR4796)
•  no small hot dust,
so black body
temperatures
Cold disks
Cold debris disks
only detected at
long wavelengths
(Wyatt, Dent & Greaves
2003; Eiroa et al. 2011)
See Marshall talk
Could be disks of
unstirred
material at large
distances (e.g.,
Heng & Tremaine 2010;
Krivov et al. 2013)
Origin of stirring in debris disks
Planet formation models
predict Pluto-formation
stirs planetesimal disks
causing a bright ring that
propagates outward (Kenyon
& Bromley 2008)
Alternative explanation is
secular perturbations of
close-in giant planets, or
binary companion (Mustill &
Wyatt 2009)
Or passing stars, or
related to protoplanetary
disk (dispersal)?
See Ormel talk
Mid-IR spectra indicate dust composition
Spitzer IRS took 5-35μm
spectra of many debris
disks (e.g., Chen et al. 2006)
Many, but not all, include
spectral features that are
identifiable with features
and materials seen in Solar
System objects
To extract compositional
information, spectra need
to be modelled, as also
information about dust size
and temperature
Lisse et al. (2012)
See Menard talk
HD113766: composition is asteroidal
Composition is mixture of
pyroxenes and olivines, similar
to S-type asteroid (Lisse et al. 2008)
Dust is 0.1-20μm, total
observed mass ~320km radius
asteroid
Could be recent collision or
ongoing planet formation, but
massive asteroid belt is possible
Other systems are similar, such
as BD+20307 (Weinberger et al. 2011)
HD172555: Giant collision
Composition is dominated by
silica (Lisse et al. 2009)
Dust is 0.1-100μm in size
Total observed dust mass is
similar to 150-200km radius
asteroid, but parent body (and
total debris mass) likely larger
Possible SiO gas? (Also OI, see
Riviere-Marichalar et al. 2012)
Points to a recent collision at
>10km/s between two massive
(>Mars-mass) protoplanets;
e.g., Earth-moon forming
collision
Problem with radiation pressure
Recent collision? Maybe
but was detected in 1983
(IRAS), so continual
replenishment needed, and
if so bound grains should
dominate smaller dust due
to their much longer
lifetimes…
Size distribution of HD172555 dust
Radiation pressure blow-out limit
Inferred size distribution
includes 0.1-100μm dust,
but <4μm dust should be
removed on dynamical
timescales (a few years)
Forsterite
Mid-IR features only
detectable in systems with
hot dust
Herschel spectroscopy
could detect 69μm feature
from crystalline olivine
Shape gives temperature
(85K) and composition
(1% Fe) for β Pic (de Vries et
al. 2012)
See Bouwman talk
Low Fe content similar to
SS comets, and implies
origin in unequilibriated
<10km bodies
CO: rare, secondary ALMA images
show CO is
abundant in β
Pic (Dent et al. in
prep)
Height (AU)
CO photodissociates in 120yrs in interstellar radiation field, so is not
commonly detected in optically thin debris disks, but a handful have
CO in emission (Zuckerman et al. 1995; Moor et al. 2009)
40
20
0
-20
-40
Telesco et al. (2005)
Under embargo -150
-100 -50
0
50
100
Distance along midplane (AU)
This gas is thought to be secondary, and to be produced in the
destruction of planetesimals, either in high velocity collisions,
photodesorption, sputtering
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Conclusions