Atomic Spectra in Astrophysics

Transcription

1 Atomic Spectra in Astrophysics Potsdam University : Wi : Dr. Lidia Oskinova lida@astro.physik.uni-potsdam.de

2 Nebular spectra Wind blown bubble NGC 7635, Hα image

3 Introduction 02 Extended sources Emission line spectra Forbidden lines: [OIII] λλ4959, 5007 [NII] λλ6548, 6583 [OII] λλ3726, 3729 Permitted H lines: Hα λ6563 Hβ λ4861 Hγ λ4340 Weaker HeI λ 5876, HeII λ4686 CII, CIII, CIV, OII... Continuum bf in Pashen continuum of HI λ>3646 Balmer λ<3646 Thermal dust emission Radio: electron bremsstrahlung UV spectra MgII λλ2796, 2803 CIV λλ1548, 1551 IR spectra [NeII] λ12.8µm [OIII] λ88.4µm

4 Basic Physical Ideas 03 Photoionization of H is the main energy-input mechanism Thermal electrons recombine. The degree of ionization at each point in the nebula is fixed by number of photoionizations and recombinations. Degree of ionization in the nebula depends on the temperature of the star, not temperature of the nebula. [Ne V] and [Fe VII] can be observed around hot stars. Electrons recombine on upper levels. The excited atoms decay to lower levels via radiative transitions. This is the origin of the HI Balmer- and Paschen-line spectra. Recombination of H + gives rise to excited atoms of H 0, emission of HI spectrum. Also, recombination spectra of HeII (strongest line λ4686). Weaker spectra of CII, CIII, CIV et cet. In addition ot line spectrum - IR contium spectrum of warm dust.

5 Basic Physical Ideas: Forbidden Lines 04 An ion is excited by collision with an electron H (IP = 13.6 ev), HeI (IP=24.5 ev), and HeII (IP=54.1 ev) What are resulting photoelectron energies? =10.9 ev Low energy electrons: not enough to collisionaly ionize (!) But enough to exite the electron to a higher level within the ground state. Elements with excitation potential < 10eV from gournd level Even when these elemens are rare in comparison with H and He, nearly all photoelectrons will be used to excite these elements Such lines with low excitation potential are forbidden lines. Zanstra: intesity of forbidden line is proportianal to UV radition field intensity

6 Basic Physical Ideas: Forbidden Lines 05 Electron is excited to a higher level within the ground state. Singly ionized sulfur, S+, has 3 valence electrons. 1s 2 2s 2 2p 6 3s 2 3p 3. Radiative transition are allowed when l=± 1, m=0,± 1 If selection rules for dipole radiation are not fulfilled - dipole radiation is not possible Quadropole or magneto-dipole radiation may occure: But the probability of such transaction is 10 5 times smaller Metastable states - excited states which have a relatively long lifetime due to slow radiative and non-radiative decay Forbidden transitions: upper-state lifetimes of ms or even hr. Allowed transitions: upper-state lifetimes are a few Osterbrock & Ferland, 2005 nanos. The lifetime: the 5/2 and 3/2 levels are 3846s and 1136s.

7 Basic Physical Ideas: Forbidden Lines 06 Removal from metastable level by collisions. In air under standard conditions, an atom experinces collisions per second In a typical nebula, an atom experiences 1 collision per minute. Or less. Therefore, there is enough time for radiative transition from metastable level. Removal from the metastable level by absorption of radiation. The probability is proportional to the density of radiation? How density of radiation changes with distance? Eddington: lack of forbidden lines in the stellar spectra - too high density of exciting photons. E.g. denstities in chromospheres are low enough for forbidden lines of [FeII] to be seen, but radiative filed is too high Thus nebulae fulfill both conditions: low densities of radiation AND low gas denstiy

8 07 Filters: F502N [O III], FR505N [O III] and F658N (Hα+[N II])

9 A typical planetary nebula spectrum.

10 Further effects: The Bowen Fluorescence Mechanism 09 There is accidental coincidence between the wavelength of HeII Lyα λ and OIII 2p 2 3 P 2 λ303.8 The HeII Lyα photons are scattered many times before they escape. Hence there is a high density of these photons in a nebula. O ++ is also present in the same zone as He ++. Some of HeII Lyα photons are absorbed by O ++ and excite 3d 3 P level of OIII. This level quickly decays by radiative transitions: Probability 0.74: resonance scattering 2p 2 3 P 2-3d 3 P 2 Probability 0.24: 2p 2 3 P 1-3d 3 P 2 λ Probability 0.02: the 3d 3 P 2 level decays by emitting one of the six longer wavelength photons (see Fig. 4.6 in Osrebrock). These lines are in optical and UV: Bowen resonance-fluorescent mechanism. Bowen fluorescence: conversion HeII Lyα photons in optical and UV lines of OIII. These lines are commonly observed in planetary nebula. HI Lyβ λ and OI 2p 4 3 P 2-2p 3 3d 3 3 D 3 λ Some atomic oxygen exists in the He + zone, due to rapid charge transfer between oxygen and hydrogen. The far UV lines of OI are also observed.

11 Partial energy-level diagram of [OIII] and HeII showing coincedence of HeII Lα and [OIII]2p 2 3 P 2-3d 3 P 0 2. Solid lines: fluorescence lines in optical and UV.

12 Exercise: Determiming the Gas Density 11 The Sulfur Lines Singly ionized sulfur S +, 1s 2 2s 2 2p 6 3s 2 3p 3 3 valence e All upper levels are metastable They can be populated only by collisions The lines of interest have very close wavelength Nearly all collisions which can exite level 3/2, can excite 5/2 But! g 5/2 =6, g 3/2 =4 Why? What is the meaning of statistical weight? Life-time 5/2 is 3846 sec, 3/2 is 1136 sec Collisions can excite and deexcite Which line is more likely to be deexcited by collisions? lowest ground-state levels of p electrons For low densities (<100 cm -3 ): deexcitation by ph emision Ratio of λ6716 to λ6731 is equal to the ratio of What? For high densities (>10000 cm -3 ): deexcitation by collisions Ratio of λ6716 to λ6731 is equal to the ratio of What? Short lived level can emit more photons What is air density?

13 Exercise: Determiming the Gas Density 12 Follow Variation of λ6716/λ6731 ratio with density

14 X-ray spectroscopy 13

15 X-ray and optical comparison 14

16 15 Most information about the Universe: EM radiation Different physics: different type of radiation Measurable quantities:

17 Attenuation of photons in the atmosphere I 16

18 Attenuation of photons in the atmosphere II 17 Optical depthτ E = κ E ρds κ E mass absorption coefficient [κ E ]=cm 2 g -1 The Universe in X-rays is visible only from space

19 Major Modern Telescopes I XMM-Newton 18 X-ray Multi-Mirror ESA (with NASA) kev Orbit: 7000 km peregee km apogee 58 hours = 170 ksec θ=6 arcsec X-ray all-sky survey catalog, currently objects best sensitivity achieved so far biggest science satellite ever built in Europe 200 m 2 polished gold mirrors

20 Major Modern Telescopes I Chandra 19 NASA s Great Observatory NASA kev Orbit: km peregee km apogee 64 hours = 240 ksec θ=0.5 arcsec (Unprecedented!) best imaging for many decades to come best spectral resolution

21 The astrophysical significance of X-ray observations 20 Direct insight into accretion onto compact objects the most efficient process known in E=mc 2 sense Physical properties of space-matter in the near environment of black holes Physics of coronae and shocks : stars and supernovae Metal enrichment of interstellar medium Eliptical galaxies and clusters: the profile of dark matter halo, the enrichment hystory Cooling flows provide estimate of the mean density in the Universe

22 Inner Shell Processes X-ray fluorescence An electron can be removed from inner K- shell (how many electornes are there?) 21 The vacancy is filled by a L-shell electron Kα-line. If the vacancy is filled by M-shell electron Kβ-line. Iron is abundnat element with relatively large cross-section for K-shell ionization: Kα line at 6.4 kev is commonly observed from astrophysical objects See Grotrian diagrmans in Kallman+ 04, ApJSS 155, 675

23 22 Schematic X-ray spectrum of AGN From W.N.Brandt "X-raying Active Galaxies" AAS see Gimenez-Garcia+ 15, A& A, Fig. B2

24 23 Famous sketch (unification model) 23 AGN with 10 8 M BH R G 3x10 13 cm Accretion disk cm BLR cm Torus cm?? NLR cm Jets cm Urry & Padovani 1995 PASP 107, 803

25 24 The speed of matter within the jets is large fraction of c. SR s effects must be taken into accout: relativistic beaming, relativistic Doppler effect, superluminal motion The inner parts of the discs are close to the BH. GR effects must be taken into account: gravitational red-shift. The emitted radiation interacts with the discs and the surrounding matter Wikipedia

26 25 Reminder: β v γ 1 c 1 β 2 25 Classical Doppler effect λ=λ o ±λ o β cosϑ Relativistic Doppler effect λ=γ 1 (λ o ±λ o β cosϑ) Superluminal velocity and relativistic beaming A source moves nearly as fast as its radiation motion which is greater than the speed of light illusion of apparent transverse The apparent transverse velocity has a maximum All radiation is confinded to a narrow cone: relativistic beaming

27 26

28 27 Reminder: 27 Gravitational Redshift λ=λ o +λ o ( (1 r s /r)) 1, where r s = 2GM/c 2 is Schwarzschild radius

29 28 Relativisitc broadening Fe-line 28 Fabian et al PASP 112, 1145

30 29

31 30 Time average (ASCA) observations of AGN MGC / average 1997 / average Lightcurve minimum Flare Fabian et al PASP 112, 1145

32 Thermal spectrum: observations of Perseus cluster of galaxies 31 Hitomi Col Nature

33 Thermal plasma 32 Thermodynamic equilibrium occurs if N e > Te 0.5 Ei 3 j cm-3 For T=10 MK and H-like Iron, N e >10 27 cm -3 For T=0.1 MK and H-like Oxigen, N e >10 24 cm -3 These are very high densities occuring hardly anywhere outside stars Coronal/Nebular N e <10 16 cm -3 Astrophiscally important plasmas * * * * * * kt e I p Ionization and excitations are by collisions is balanced by radiative and dielectronic recombinaiton The state of ionization is determined by the temperature Excited ions return to the ground state t(recomb) < time(collision) Cooling is radiative Produced X-rays leave without interacting with the plasma,

34 Ionisation 33 Collisional ionization: e - +I I + +2e - Photoionization: γ+i I + +e - Inner shell ionization: e -,γ+i I *+ +2e - I + e -,γ Inner shell ionization:k-shell electron (ie 1s electron) is removed. Remining ion is very unstable. It will either emit a photon (radiatively stabilize) or an electron, called an Auger electron. Whether a photon or an electron is emitted depends upon chance and the ion involved. As Z increases, the probability of a photon being emitted increases; for iron, it is ~30%. For oxygen, it is ~ 1%. Innershell ionization of Fe I - Fe XVI tends to emit a 6.4 kev photon, commonly called the cold or neutral iron line.

35 Equilibrium in thermal plasma 34 Thermal plasma can be in equilibrium or out of it. Ionization equilibrium (CIE plasma) Ionization of ion z of element Z is balanced by recombination C Z,z 1 ionization rate,α Z,z recombination rate n Z,z 1 C Z,z 1 = n Z,z (α rad Z,z +αdi Z,z ) Plasma codes: e.g. Astrophysical Plasma Emission Code - APEC Large variety of astrophysical sources: stars Non-equilibrium ionization (NIE plasma) - ionization rate is higher than recombination - or recombination rate is higher than ionization dynamic time scale is shorter than required to establish IE NEI - codes occurs e.g. in supernova remnants

36 APEC simulated spectra for two different T(Chandra MEG+1) 35 FeXXVI FeXXV CaXIX SXV SiXIV SiXIII MgXII MgXI NeX NeX NeIX FeXVIII FeXVII OVIII FeXVII FeXVII OVIII OVII NVII 30 Counts/sec Line emission dominates at kt=0.6 kev (T=7MK) Strong continuum at kt=6 kev (T=70MK) NB! Instrumental responce No interstellar absorption Wavelength [A o ]

37 High-Resolution X-ray Spectra? Find He-like ions? 36 SiXIII MgXII MgXI NeX NeX NeIX FeXVIII FeXVII OVIII FeXVII FeXVII OVIII OVII NVII NVI CVI XMM RGS enhnaced N ζ Pup O4I XMM RGS ζ Ori O9.7I Wavelength (A o ) * Overall spectral fitting plasma model, abundunces * Line ratios T X (r), spatial distribution * Line profiles velocity field, wind opacity

38 Common X-ray diagnostics: lines of He-like ions Ratio of forbidden to intercombination line flux depends on? resonance UV flux dilutes with 37 r -2 f/i ratio estimator for distance where the hot gas is located Requires knowledge of stellar UV field intercombination UV forbidden OVII Gabriel & Jordan 1969

39 Comparing OVII in early and solar type stars 38 B0.5IV+BV at d=98 pc X-ray brightest massive star on sky Soft spectrum, narrow lines compare to solar type star OVII α Cru Capella Flux (Counts/sec/Angstrom) λ R λ I λ F Oskinova etal. in prep Chandra has 0.6 arcsec resolution

40 Inner Shell Processes Auger Ionization Additional source of ionization in plasma 39 See Grotrian diagrmans in Kallman+ 04, ApJSS 155, 675