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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 150 Conclusions