1 / 52

Supernova Remnants and their Emission

Supernova Remnants and their Emission. Aya Bamba (Aoyama Gakuin U.). 青学の宇宙物理 メンバー. Member of astro group in Aoyama Gakuin Univ. 井上 剛志 ：　星間物理学（理論）、プラズマ物理（理論） 　　　　　星形成、超新星残骸、ＧＲＢ、磁気流体シミュレーション 大平 豊 ：　高エネルギー 宇宙物理学（理論）、プラズマ物理（理論） 　　　　　　無衝突 プラズマ現象、宇宙線粒子加速、 SNR 、 CTA 、 CALET

misu
Download Presentation

Supernova Remnants and their Emission

An Image/Link below is provided (as is) to download presentation Download Policy: Content on the Website is provided to you AS IS for your information and personal use and may not be sold / licensed / shared on other websites without getting consent from its author. Content is provided to you AS IS for your information and personal use only. Download presentation by click this link. While downloading, if for some reason you are not able to download a presentation, the publisher may have deleted the file from their server. During download, if you can't get a presentation, the file might be deleted by the publisher.

E N D

Presentation Transcript

  1. Supernova Remnants and their Emission AyaBamba (Aoyama Gakuin U.)

  2. 青学の宇宙物理メンバー Member of astro group in Aoyama Gakuin Univ. 井上 剛志：　星間物理学（理論）、プラズマ物理（理論） 　　　　　星形成、超新星残骸、ＧＲＢ、磁気流体シミュレーション 大平 豊：　高エネルギー宇宙物理学（理論）、プラズマ物理（理論） 　　　　　　無衝突プラズマ現象、宇宙線粒子加速、SNR、CTA、CALET 坂本 貴紀：　宇宙物理学、雷雲ガンマ線（実験・観測） Swift、小型ロボット望遠鏡、ＧＲＢ、雷・スプライト/TGF 澤田 真理：　精密プラズマ分光（実験・観測） Suzaku/Astro-H、超新星残骸、銀河系中心 柴田 徹：　宇宙線物理学（実験） 　　　　　宇宙線の伝播、中間子多重発生、ＣＴＡ 田中 周太：　高エネルギー宇宙物理学（理論） 相対論的電子・陽電子プラズマ、パルサー、パルサー星雲、 馬場 彩：　高エネルギー宇宙物理学（実験・観測） Suzaku/ASTRO-H、ＣＴＡ，超新星残骸、パルサー星雲 山崎 了：　高エネルギー宇宙物理学（理論） 　　　　　ＧＲＢ，超新星残骸、宇宙線粒子加速、ＣＴＡ、Ｆｅｒｍｉ 吉田 篤正：　宇宙物理学（実験） 　　　　　　ＭＡＸＩ、ＣＡＬＥＴ、ＣＴＡ，ＧＲＢ、中性子星、Suzaku/Astro-H

  3. Once a star explodes … Anyhow, stars explode. (An-pan-man knows it ..) We can study the remnants even we do not know how stars explode. Parameters: Energy input: ~1051 erg Duration: delta function Mass: ~Msun remnants of dead stars: Supernova remnants (SNRs) • Evolution of SNRs and emission • Acceleration of particles on shocks of SNRs • and emission Goal: SNRs makes chemical and high-E of our diverseness

  4. An-pan-man March (one of the most famous super-star) time flies quickly, shining stars will disappear, so you should go ahead with smile !

  5. Once a star explodes … Anyhow, stars explode. (An-pan-man knows it ..) We can study the remnants even we do not know how stars explode. Parameters: Energy input: ~1051 erg Duration: delta function Mass: ~Msun remnants of dead stars: Supernova remnants (SNRs) • Evolution of SNRs and emission • Acceleration of particles on shocks of SNRs • and emission Goal: SNRs makes chemical and high-E of our diverseness

  6. 0. Why X-ray observation is important for SNR study ? X-ray detector electrons X-rays ne∝ EX X-ray detectors basically change the X-ray to number of electrons. We can measure - position - time - energy simultaneously !

  7. Evolution of supernova remnants • and emission

  8. 1.1. When SNRs are young The surrounding interstellar matter can be negligible -> All explosion energy is used to the kinetic energy of the exploded star (ejecta) when Eexp=1051erg and M*=10Msun, vshock = 3.2x108 cm/s (10% of the light speed) uniform expansion free from deacceleration radius ∝ vshockt“free expansion phase” only kinetic E -> no energy dissipation no emission

  9. Expansion is “visible” Tycho’s SNR (SN1572) (Katsuda+10) comparison of X-ray images taken on different period -> detection of movement of shock

  10. How long can we use free expansion approximation ? Stars explode in interstellar matter. ISM The shock sweep up the ambient ISM. When mass of swept-up mass exceed the mass of the exploded star, we cannot ignore the swept-up ISM.

  11. the radius of SNR: vshockt the swept-up volume: 4/3p(vshockt)3 the swept-up mass: 4/3p(vshockt)3rISM the expansion starts to stop when; (assumption: uniform density)

  12. Observed expansion rates (Moffett+93) R∝tm Shocks of famous SNRs already starts decelerate

  13. 1.2. When the shock starts decelerate Self-similar solution by Sedov R∝t0.4 v∝t-0.6 T∝t-1.2 loss of kinetic E -> thermal E of downstream plasma E loss by emission is still negligible “Sedov phase” or “adiabatic phase”

  14. X-ray gallery of young supernova remnants Kepler’s nova Tycho’s nova Cassiopeia A SN1006 beautiful fireworks in the universe ~1 event / 30 yrs

  15. Thermal emission from the heated plasma (1) bremsstrahlung Since the downstream is so hot (~106-8 K or 0.1-10 keV), the gas is almost ionized. + + bending the direction by coulomb interaction = acceleration radiation ! bremsstrahlung + - + +

  16. Thermal emission from the heated plasma (2) line emission Electrons in atoms orbit around nuclei. When they change their orbit, they emit/absorb photon with transition energy -> emission lines In the plasma, atoms are highly ionized into around He-like or H-like ions. He-like H-like

  17. X-ray emission lines -> we can know how much heavy elements are distributed into interstellar medium Si Ne Mg O S Ar Ca Fe Identification of major heavy elements ! (cal src) Tycho’s spectra by Suzaku

  18. Searching for minor elements There are many kinds of elements ! All should be made in stars and distributed by supernovae Important to measure the amount of elements near iron (produced in imcomplete Si burning) in order to understand how heavy metals are produced chromium manganese

  19. Suzaku detection of Cr and Mn emission lines from Tycho Suzaku100ks observation -> detection of Cr and Mn lines ! MMn/MCr = 0.5 (0.2-0.7) (Tamagawa+08) First detection of emission lines from minor elements Near-future observatories w. excellent E resolution will detect minor elements from many SNRs.

  20. How elements dissipate into the interstellar medium? It should have information of its explosion mixing ? onion-like? The line distribution of Cas A by Chandra (Hwang+04) Si Fe heavier elements are located inside of lighter elements ?? We need 3D information

  21. Lines have doppler broadening by expansion Velocity in km/s (Hayoto+11) blue-shift red-shift Si S Ar Fe Tycho spectrum (Suzaku) line shift -> expansion velocity Radial peak in arcmin (ASCA: Hwang & Gotthelf 97) Line broadening due to expansion Heavier elements stay inside of the remnant.

  22. The plasma age Ionization is mainly by collision of ions and electrons in SNR plasma. Plasma in SNRs are so tenuous and ionization takes long time. In order to reach the equilibrium between temperature and ionization, nt ~ 1012 s cm-3 if n ~ 1 cm-3, t ~ 3x104 yr Plasma before equilibrium: non equilibrium state or ionizing (check Yamaguchi-san’s talk)

  23. How emission changes with different nt

  24. 1.3. When the plasma cooled down below 2MK … Radiative cooling coefficient (Gehrels+93) Plasma emit more and more -> cool down easily -> more efficient emission -> … Cooling of plasma is accelerated with strong emission radiative cooling phase

  25. Emission of plasma -> taking E out from the shock -> shock speed slows down more cool and dense shell hot plasma ISM cool down -> pressure can be ignored shell collects ISM further like snowplow “snowplow phase” R∝t2/7 -> R∝t1/4

  26. mixed morphology SNRs shells are already cold to emit X-rays radio X-ray ejecta is still hot IC443 (Keohane+)

  27. 1.4. Disappearance of SNR The shock speed slows down more and more -> comparable to the proper motion of surrounding ISM (10-20 km/s) SNRs lose the boundary between the outside -> disappear of SNR time scale ~ 106 yrs

  28. 2. Acceleration of particles on shocks of SNRs

  29. knee=1015.5eV ankle=1018.5eV (Cronin 1999) 2.1. cosmic rays very high E particles in the universe uCR ~ 1eV/cc c.f. CMB 0.3 eV/cc star light < 0.3 eV/cc magnetic field 0.3 eV/cc turbulence 0.3 eV/cc thermal E 0.01 eV/cc 1 CR per your fingertip per 1 second. one of the main components of our Galaxy

  30. 2.2. Shocks of SNRs are cosmic ray accelerator ! chemical abundance of cosmic rays made through separation of heavier elements

  31. Be: basically made through separation including radio isotope 14Be typical age of CR: ~6x106 years (Garcia-Munoz+77) ~escape timescale from Galaxy Galaxy volume: 5x1066 cm3 CR energy density: 1.6x10-12 erg cm-3 -> we need energy input to CR of 1x1040 erg s-1 E input by SNRs: 1051 erg per 30 years = 1042 erg s-1 If 1% of SN energy is injected CRs, we can explain all of the E of CRs by SNRs.

  32. 2.3. 3 min. recipe of particle accleration (1) terminology shock front (us) outer region (u1 = 0) inner region (u2) SN shock front (us = 0) upstream region (uu = us) downstream region (ud = us) g-1 g+1 lab. system: shock system: ideal gas: g = 5/3 → ud = 1/4us

  33. 2.3. 3 min. recipe of particle accleration (2) acceleration uy shock front lab. system x ux ud uu paticles change their direction with scattering by magnetic field turbulence E conversation in upstream/downstream in shock system particles get energy always crossing shock spectrum of particle: power-law <- same to cosmic rays !

  34. 2.4. Maximum E of particles Particles gyrate with magnetic field (gyro motion) The radius should be smaller than the size of the system larger system stronger magnetic field can accelerate particles to higher energy

  35. 2.5. Nonthermal emission from accelerated particles (1) synchrotron emission electrons gyrate by magnetic field = acceleration ! -> emission - In the case of e is relativistic: synchrotron emission typical emission frequency: in 1uG magnetic field, ~GeV e -> 1012 Hz (radio band) ~TeV e -> 1018 Hz (X-ray band)

  36. (Yamaguchi+08) Sync. X-rays (Bamba+08) thermal X-rays SNRs are really hot bubbles. Shock fronts accelerate (at least) electrons.

  37. Spectrum of synchrotron emission ∝ B2 (ve ~ c) When electron spectrum is power-law, The spectrum of the synchrotron emission is, We can know the index of electrons from the spectrum of synchrotron emission. CR has power-law spectrum, we can expect the index of synchrotron emission.

  38. 2.6. Nonthermal emission from accelerated particles (2) inverse Compton emission hn hn’ - (SNR case: photons are CMB) - collision between particles and photon -> photon get E from particles “inverse Compton emission” Uph: E density of scattered photon We can measure magnetic field

  39. 2.7. Nonthermal emission from accelerated particles (3) emission via pi-on decay When protons with E>1GeV collide with other protons, sometimes produce pi-0 meson. pi-0 meson decays into two photons. We need a lot of protons ! molecular cloud etc. Only footprint by CR protons !! + +

  40. Summary of emission lines bremss. pi-0 IC sync. If we can detect each component, we can know thermal and nonthermal matter in SNRs.

  41. SNRs are also detected in TeV gamma-ray band (Acero+10) shells of SN1006 is also TeV emitter lines bremss. pi-0 IC sync. IC (e origin) ? pi-0 (p origin) ? No confirmation yet…

  42. GeV emission ! (Abdo+10) lines bremss. pi-0 W44 IC sync. pi-0 emission …? association with molecular cloud Only several SNRs are detected in GeV gamma-rays. It is still unknown what makes difference.

  43. 2.8. Topics: thin filaments on shocks Sync. X-rays forms thin filaments on shock fronts -> gyro-radius of electrons is so small SN1006 NE shell 0.3 – 2.0 keV 2.0 – 10.0 keV ~0.01 – 0.1 pc ! (Bamba+03)

  44. e- shock Why diffusion is so small ? • frequent scattering of electrons • magnetic field turbulence scatter electrons • turbulent B -> frequent scatter -> small diffusion • small gyro radius (= large magnetic field) • a lot of accelerated electrons • -> ~ large current ~ induced magnetic field • -> more efficient acceleration -> … Very efficient acceleration ! (Bamba+05)

  45. 2(g-1) (g+1)2 kTd = mvs2 kinetic E thermal E of downstream plasma thermal E of downstream plasma E of accelerated particle kinetic E kinetic E Energy budget on shock From Rankine-Hugoniot relation, ideal gas~ 0.19vs2 with E loss < 0.19vs2 If the acceleration is very efficient, the downstream plasma becomes cooler than the case without acceleration

  46. 3. Summary • Stars and their explosions make • the chemical and high energy diverseness • of our universe. • X-ray observations are one of the best tools • to investigate such kind of diverseness.

  47. SNRs are HanasakaJiisan in our universe (flower-blossom-old man) Japanese old story “Hanasaka” When he sprays ash onto a deadwood, the wood becomes a cherry tree in full bloom

  48. 2(g-1) (g+1)2 kTd = mvs2 kinetic E thermal E of downstream plasma kinetic E Energy budget on shock From Rankine-Hugoniot relation, ideal gas~ 0.19vs2

  49. Power of bremsstrahlung ∝ne2 gff: gaunt factor (Bressaard+62) When the plasma is thermallized probability of ve = ve ~ ve+dve spectrum has the cut-off at hn=kT -> we can know the temperature of plasma

More Related