how do multiple planetary systems shape the dust disk? . - PowerPoint PPT Presentation

The HL Tau system . Giovanni Picogna Tübingen Universität, CPT & Kepler Center 6th May 2015 Exoplanets in Lund 2015 Lund Observatory how do multiple planetary systems shape the dust disk?

. introduction

. ∙ We are now obtaining pristine images of the protoplanetary disk evolution that we can use to constraint planet formation models. ∙ an outstandig example is the HL Tau system, imaged by ALMA in the mm continuum ∙ where axysimmetric ring structures and gaps are visible 2 the hl tau system

. ∙ We are now obtaining pristine images of the protoplanetary disk evolution that we can use to constraint planet formation models. ∙ an outstandig example is the HL Tau system, imaged by ALMA in the mm continuum ∙ where axysimmetric ring structures and gaps are visible HL Tau system. Source: http://www.eso.org/public/news/eso1436/ 2 the hl tau system Figure 1:

. ∙ We are now obtaining pristine images of the protoplanetary disk evolution that we can use to constraint planet formation models. ∙ an outstandig example is the HL Tau system, imaged by ALMA in the mm continuum ∙ where axysimmetric ring structures and gaps are visible HL Tau system - continuum 233 GHz image. Source: http://www.eso.org/public/news/eso1436/ 2 the hl tau system Figure 1:

. There are different physical processes capable of creating rings in a disk: ∙ multi-planetary system ∙ zonal flows 3 ring structure formation Figure 2: Meru et al., 2014

. There are different physical processes capable of creating rings in a disk: ∙ multi-planetary system ∙ zonal flows 3 ring structure formation Figure 2: Flock et al., 2014

. I focus on the planetary origin of those structure ∙ straightforward explanation ∙ no detailed analysis yet of dust filtration and dynamical evolution in a multi-planetary system 4 multi-planetary system scenario

. I focus on the planetary origin of those structure ∙ straightforward explanation ∙ no detailed analysis yet of dust filtration and dynamical evolution in a multi-planetary system 4 multi-planetary system scenario

. numerical method

6 256 Non–reflecting Outer boundary Open Inner boundary 512 Cells in azimuthal direction Cells in radial direction . Value Quantity cylindrical choordinates as in Zhu et al. (2014) ∙ integrated with a semi–implicit and fully implicit integrator in treated as Lagrangian particles (Müller) ∙ modified to study the dynamic of a population of small bodies ∙ I have used the FARGO 2D code (Masset et al., 2000) numerical method

6 256 Non–reflecting Outer boundary Open Inner boundary 512 Cells in azimuthal direction Cells in radial direction . Value Quantity cylindrical choordinates as in Zhu et al. (2014) ∙ integrated with a semi–implicit and fully implicit integrator in treated as Lagrangian particles (Müller) ∙ modified to study the dynamic of a population of small bodies ∙ I have used the FARGO 2D code (Masset et al., 2000) numerical method

6 256 Non–reflecting Outer boundary Open Inner boundary 512 Cells in azimuthal direction Cells in radial direction . Value Quantity cylindrical choordinates as in Zhu et al. (2014) ∙ integrated with a semi–implicit and fully implicit integrator in treated as Lagrangian particles (Müller) ∙ modified to study the dynamic of a population of small bodies ∙ I have used the FARGO 2D code (Masset et al., 2000) numerical method

7 . 25,50 Planet semi–major axes ( au ) 1,5,10 Planet masses ( M th ) 0.55 0.1,1,10,100 Dust size (cm) 2.6 Dust density (g/cm 3 ) Dust particles Isothermal EOS -1 Temperature profile -1 Surface density profile 0.004 0.05 Aspect ratio [2.5,100] Disk extent ( au ) 0.135 Value Physical quantity parameter space Disk mass ( M ⊙ ) Viscosity ( α SS ) 1 × 10 6 Star mass ( M ⊙ )

. ∙ no self-gravity ∙ no back reaction of the particles on the gas ∙ isothermal EOS ∙ 2D simulations 8 limitations

. ∙ no self-gravity ∙ no back reaction of the particles on the gas ∙ isothermal EOS ∙ 2D simulations 8 limitations

. ∙ no self-gravity ∙ no back reaction of the particles on the gas ∙ isothermal EOS ∙ 2D simulations 8 limitations

. ∙ no self-gravity ∙ no back reaction of the particles on the gas ∙ isothermal EOS ∙ 2D simulations 8 limitations

. major physical parameters

10 M p R R 2 ∙ Viscous criterion . R R ∙ Thermal criterion gap opening criteria ( H ) 3 M th = c s 3 = M ⋆ G Ω p ( H ) 3 = 1 . 25 × 10 − 4 = q > M ⋆ ( H ) 2 q ≥ 40 ν = 4 × 10 − 4 = 40 α SS p Ω p

11 . ∙ It quantifies the coupling between the solid and gas components ∙ We adopted the formula by Haghighipour & Boss (2003) that smoothly combines the Epstein and Stokes regimes. 8 fC D v rel stopping time ] − 1 [ τ s = τ f Ω K = ρ • a • ( 1 − f ) v th + 3 Ω K ρ g F D = − 1 ∆ v p τ f ∙ For our parameters the dm-size particles have a stopping time ∼ 1

. results

∙ mm cm dm m –sized particle evolution . For videos look at this webpage 13 10 thermal masses

. ∙ ring fragmentation in 5 high–mass stable points near 5:3 orbital resonance with the inner planet ∙ vortices formation in the inner gap ∙ coorbital regions destabilized by the outer planet ∙ mm-size dust migrate through the gap 14 10 thermal masses

. ∙ ring fragmentation in 5 high–mass stable points near 5:3 orbital resonance with the inner planet ∙ vortices formation in the inner gap ∙ coorbital regions destabilized by the outer planet ∙ mm-size dust migrate through the gap 14 10 thermal masses

. ∙ ring fragmentation in 5 high–mass stable points near 5:3 orbital resonance with the inner planet ∙ vortices formation in the inner gap ∙ coorbital regions destabilized by the outer planet ∙ mm-size dust migrate through the gap 14 10 thermal masses Figure 3: cm-sized particles

. ∙ ring fragmentation in 5 high–mass stable points near 5:3 orbital resonance with the inner planet ∙ vortices formation in the inner gap ∙ coorbital regions destabilized by the outer planet ∙ mm-size dust migrate through the gap 14 10 thermal masses Figure 3: cm-sized particles

. ∙ ring fragmentation in 5 high–mass stable points near 5:3 orbital resonance with the inner planet ∙ vortices formation in the inner gap ∙ coorbital regions destabilized by the outer planet ∙ mm-size dust migrate through the gap 14 10 thermal masses

∙ Final surface density distribution . ∙ Final eccentricity distribution 15 10 thermal masses

∙ Final surface density distribution . ∙ Final eccentricity distribution 15 10 thermal masses

∙ mm cm dm m –sized particle evolution . For videos look at this webpage 16 5 thermal masses

. ∙ ring fragmentation in 5 high–mass stable points ∙ no long–lived vortex ∙ coorbital regions destabilized by the outer planet ∙ particle exchange between coorbital regions 17 5 thermal masses

. ∙ ring fragmentation in 5 high–mass stable points ∙ no long–lived vortex ∙ coorbital regions destabilized by the outer planet ∙ particle exchange between coorbital regions 17 5 thermal masses

. ∙ ring fragmentation in 5 high–mass stable points ∙ no long–lived vortex ∙ coorbital regions destabilized by the outer planet ∙ particle exchange between coorbital regions 17 5 thermal masses

. ∙ ring fragmentation in 5 high–mass stable points ∙ no long–lived vortex ∙ coorbital regions destabilized by the outer planet ∙ particle exchange between coorbital regions 17 5 thermal masses

∙ Final surface density distribution . ∙ Final eccentricity distribution 18 5 thermal masses

∙ Final surface density distribution . ∙ Final eccentricity distribution 18 5 thermal masses

∙ mm cm dm m –sized particle evolution . For videos look at this webpage 19 1 thermal mass

. ∙ Ripples in m-sized particles distribution ∙ Particles in coorbital region very close to the planet location ∙ Ring is wider and do not fragment 20 1 thermal mass

. ∙ Ripples in m-sized particles distribution ∙ Particles in coorbital region very close to the planet location ∙ Ring is wider and do not fragment 20 1 thermal mass

. ∙ Ripples in m-sized particles distribution ∙ Particles in coorbital region very close to the planet location ∙ Ring is wider and do not fragment 20 1 thermal mass

∙ Final surface density distribution . ∙ Final eccentricity distribution 21 1 thermal mass

∙ Final surface density distribution . ∙ Final eccentricity distribution 21 1 thermal mass

. ∙ Generate ALMA–like images (e.g. with RADMC-3D + CASA package) ∙ Extend integration time ∙ Add particle back-reaction ∙ Include disk self–gravity ∙ Different mass planets ∙ Relax isothermal approximation 22 what’s next?

. ∙ Generate ALMA–like images (e.g. with RADMC-3D + CASA package) ∙ Extend integration time ∙ Add particle back-reaction ∙ Include disk self–gravity ∙ Different mass planets ∙ Relax isothermal approximation 22 what’s next?