Y. Matsuo A), M. Hashimoto A), M. Ono A), S. Nagataki B), K. Kotake C), S. Yamada D), K. Yamashita E) Long Time Evolutionary Simulations in Supernova until SNR phase Included the Uncertainties of CSMs A :Kyushu university B : RIKEN C : Fukuoka university D : Waseda university E : Yamanashi university MMCOCOS December 2 6, 2013 Cassiopeia A (Cas A) Hwang et al This indicates the mixing between Si- and Fe-rich matters during the expansion. Badences 2010 From the observations, There is the emission line of Fe in Cas A. It is likely that Fe exist outside the Si. Nobody have simulated the evolution of the elements from the onset of SN explosion to SNR phase. So we try to simulate the SN shock expansion and trace the elements (in particular H,He,O,Si,Fe) from the onset of the explosion to SNR phase (330 yr). In this study, we investigate whether Fe collides with the reverse shock at which the fluid instability develops. Motivation Progenitors 6 M He core model (Hashimoto 1995) and 3.8 M CO core model Circumstellar mediums (CSMs) We assumed that CSM consists of RSG wind and/or WR wind. We adopted the several wind parameters ( for details, please see my poster) Initial models Density profiles of the progenitor models Density profiles of the CSM models Results In the case of 6M He core models, Fe does not collide with the reverse shock in most CSM models. In the case of 3.8M CO core models, Fe collides with the reverse shock in slow wind models. 6.0M He core and slow wind (T WR = 4000 yr) 3.8M CO core and slow wind (T WR = 4000 yr) Fe would be mixed because the fluid instability develops around the reverse shock region. Red: H, green:He, blue:O, pink:Si, aqua:Fe Initial model Progenitor M He core model (Hashimoto 1995) Circumstellar medium (CSM) --- RSG wind only SN explosion We input the thermal energy ( 410 51 erg) to simulate the SN ejecta. Perturbation: Resolution: 2000 (r) 100 () When the blast wave reaches near the outer boundary of the computational domain, we increase the domain by a factor of 1.1 to trace the shock until SNR phase. Numerical methods Circumstellar Medium: only RSG wind Results X Fe Elements distribution He O-rich Si-rich Fe-rich We can see the mixing between He and O. However, we cannot see the mixing between Si and Fe. 1D analysis forward shock reverse shock Time distribution of particle positions Red: H, green:He, blue:O, pink:Si, aqua:Fe Fe and Si would be mixed if I change the density distribution of CSM so that Fe-rich ejecta will collide with the reverse shock. In this model, He and O are mixed but Si and Fe ejecta are not mixed. O collides with reverse shock at which RTI develops, but Fe does not. We try other CSM models to investigate the possibility that Fe collides with the reverse shock. Circumstellar medium models We assume that CSM consists of RSG wind and/or WR wind. We adopt the following parameters: Wind velocity (V RSG and V WR ) and duration time of WR wind (T WR ) Fast wind models V RSG = 20 km/s, V WR = 3.4 10 3 km/s, T WR = 0, 1000, 2000, 4000 yr Middle wind models V RSG = 10 km/s, V WR = 1.7 10 3 km/s, T WR = 0, 2000, 4000, 8000 yr Slow wind models V RSG = 5 km/s, V WR = 8.5 10 2 km/s, T WR = 0, 4000, 8000, yr We investigate whether Fe collides with the reverse shock under these CSM models. The position of shocks and Fe RSG wind modelR FS (pc)R RS (pc) Slow RSG wind Middle RSG wind Fast RSG wind WR wind modelR FS (pc) R RS (pc) Slow wind model (T WR =4000yr) Slow wind model (T WR =8000yr) Slow wind model (T WR =16000yr) Middle wind model (T WR =2000yr) Middle wind model (T WR =4000yr) Middle wind model (T WR =8000yr) Fast wind model (T WR =1000yr) Fast wind model (T WR =2000yr) Fast wind model (T WR =4000yr) Fe 1.2 pc 2D results of the slow wind model Slow wind model ( T WR = 8000 yr) RTI develops around the O-rich region. But Fe rich matter are not mixed because 1) the mixing time of Fe are not enough to mix 2) Fe are not mixed in the star before the shock pass through the stellar surface. The expansion velocity of Fe is too slow. From the observations, Fe reaches at the reverse shock ( r = 1.6 pc ) 1.6 pc / 330 yr = 4740 km/s ~ 5000 km/s So, we need that the Fe-velocity is about 5000 km/s at least. Now, V Fe in our models is km/s at most. We need to increase the expansion velocity of Fe by a factor of ~ 1.5 at least. Fe need to be mixed up to the O-rich layers when the shock pass through the stellar surface. How do we need the expansion velocity of Fe ? Fluid element velocity (slow RSG wind) The ejecta expand at constant velocity after the ejecta reach the stellar surface. Expansion velocities of O ~ 4000 6000 km/s Si ~ 3500 4000 km/s Fe ~ 3000 3500 km/s Red: H, green:He, blue:O, pink:Si, aqua:Fe From the observation, Fe reach at the reverse shock ( r = 1.6 pc ) 1.6 pc / 330 yr = 4740 km/s ~ 5000 km/s So, we need that the Fe-velocity is about 5000 km/s Now, V Fe in our models is km/s. We need to increase the expansion velocity of Fe by a factor of ~ 1.5 at least. Fe need to be mixed up to the O-rich layers when the shock pass through the stellar surface. How do we need the expansion velocity of Fe ? We try to simulate the formation of SNR from onset of SN explosion and compare the results and the observation of Cas A. From the observation, Fe reaches at reverse shock and it is likely that Fe and SI mixed during the explosion. In all our models, Si and Fe are not mixed because the mixing between SI and Fe at SN explosion are not enough that Fe collides with the reverse shock. We find that expansion velocity of Fe need to be increased by a factor of ~ 1.5 at least and this problem is not solved if CSM distribution are changed. We should the estimate of growth rate of RTI to investigate the reason why Fe are not mixed in the star. So, we will try to mix Fe to O-shell artificially and investigate the possibilities that Fe would exist outside the Si. Summary Our simulations have two parameters Duration time of Wolf-Rayet wind (T WR ) Wind velocity (V wind ) ( Progenitor star have H envelop or not ) Our simulations Until blast wave pass through the stellar surface ( time scale ~ sec ) Rayleigh-Taylor instability (RTI) Richtmyer-Meshkov instability (RMI) Fluid instabilities during SN explosion Joggerst et al Hammer et al. 2010 Supernova remnant phase ( time scale ~ 1 1000 yr ) Rayleigh-Taylor instability (RTI) Instabilities develop in front of reverse shock. Fluid instabilities during SN explosion Veelen et al. 2009 After core bounce ( time scale ~ 1 sec ) Standing accretion shock instability (SASI) Neutrino driven convection However, these instabilities are not key process to mix between Fe and Si because the temperature in shocked matter is too high to synthesize the Fe. Fluid instabilities during SN explosion Henke et al Couch 2013