Atomic Spectra in Astrophysics - uni-potsdam.delida/TEACH.DIR/L10spec.pdf · The excited atoms...
Transcript of Atomic Spectra in Astrophysics - uni-potsdam.delida/TEACH.DIR/L10spec.pdf · The excited atoms...
Atomic Spectra in AstrophysicsAtomic Spectra in Astrophysics
Potsdam University : Wi 2016-17 : Dr. Lidia [email protected]
Nebular spectra
Wind blown bubble NGC 7635, Hα image
02 Introduction
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
03 Basic Physical Ideas
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 H0 , 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.
04 Basic Physical Ideas: Forbidden Lines
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? 24.5-13.6=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
05 Basic Physical Ideas: Forbidden Lines
Osterbrock & Ferland, 2005
Electron is excited to a higher level within the ground
state.
Singly ionized sulfur, S+, has 3 valence electrons.
1s2 2s2 2p6 3s2 3p3 .
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 105 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
nanos.
The lifetime: the 5/2 and 3/2 levels are 3846s and 1136s.
06 Basic Physical Ideas: Forbidden Lines
Removal from metastable level by collisions.
In air under standard conditions, an atom experinces 1010 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
07
Filters: F502N [O III],
FR505N [O III] and
F658N (Hα+[N II])
A typical planetary nebula spectrum.
09Further effects: The Bowen Fluorescence Mechanism
There is accidental coincidence between the wavelength of HeII Lyα λ303.78 and
OIII 2p2 3 P2 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 3d3 P level of OIII. This
level quickly decays by radiative transitions:
Probability 0.74: resonance scattering 2p2 3 P2 -3d3 P2
Probability 0.24: 2p2 3 P1 -3d3 P2 λ303.62
Probability 0.02: the 3d3 P2 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 λ1025.72 and OI 2p4 3 P2 -2p3 3d3 3 D3 λ 1025.76. 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.
Partial energy-level diagram of [OIII] and HeII showing coincedence of HeII Lα and[OIII]2p2 3 P2 - 3d3 P0
2 . Solid lines: fluorescence lines in optical and UV.
11Exercise: Determiming the Gas Density
lowest ground-state
levels of p electrons
The Sulfur LinesSingly ionized sulfur S+ , 1s2 2s2 2p6 3s2 3p3 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! g5/2 =6, g3/2 =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?
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?
12Exercise: Determiming the Gas Density
Follow www.williams.edu/Astronomy/research/PN/nebulae
Variation of λ6716/λ6731 ratio with density
13
X-ray spectroscopy
14X-ray and optical comparison
15
Most information about the Universe: EM radiation
Different physics: different type of radiation
Measurable quantities:
16Attenuation of photons in the atmosphere I
17Attenuation of photons in the atmosphere II
Optical depth τE =
∫κEρds
κE mass absorption coefficient
[κE]=cm2 g-1
The Universe in X-rays
is visible only from space
18Major Modern Telescopes I XMM-Newton
X-ray Multi-Mirror 1999 -
ESA (with NASA) 0.2 - 12.0 keV
Orbit: 7000 km peregee
114 000 km apogee
58 hours = 170 ksec
θ=6 arcsec
X-ray all-sky survey catalog, currently 250000 objects
best sensitivity achieved so far
biggest science satellite ever built in Europe
200 m2 polished gold mirrors
19Major Modern Telescopes I Chandra
NASA’s Great Observatory 1999 -
NASA 0.2 - 12.0 keV
Orbit: 16000 km peregee
150 000 km apogee
64 hours = 240 ksec
θ=0.5 arcsec (Unprecedented!)
best imaging for many decades to come
best spectral resolution
20The astrophysical significance of X-ray observations
Direct insight into accretion onto compact objects
the most efficient process known in E=mc2 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
21Inner Shell Processes
X-ray fluorescence An electron
can be removed from inner K-
shell (how many electornes are
there?)
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
2222 Schematic X-ray spectrum of AGN
http://www.astro.psu.edu/users/niel/papers/
From W.N.Brandt "X-raying Active Galaxies" AAS’04
see Gimenez-Garcia+ 15, A& A, Fig. B2
2323 Famous sketch (unification model)
Urry & Padovani 1995 PASP 107, 803
AGN with
108 M BH
RG 3x1013 cm
Accretion disk
1013..14 cm
BLR 1016..17 cm
Torus 1017 cm ??
NLR 1018..20 cm
Jets 1017..24 cm
24
Wikipedia
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
2525 Reminder:
β ≡ vcγ ≡ 1√
1−β2
Classical Doppler effect
λ = λo ± λoβ cosϑ
Relativistic Doppler effect
λ = γ−1(λo ± λoβ cosϑ)
Superluminal velocity and relativistic beaming
A source moves nearly as fast as its radiation illusion of apparent transverse
motion which is greater than the speed of light
The apparent transverse velocity has a maximum
All radiation is confinded to a narrow cone: relativistic beaming
26
http://www.mrelativity.net
2727 Reminder:
Gravitational Redshift
λ = λo + λo(√
(1 − rs/r))−1,
where rs = 2GM/c2 is Schwarzschild radius
2828 Relativisitc broadening Fe-line
Fabian et al. 2002 PASP 112, 1145
29
3030 Time average (ASCA) observations of AGN MGC6-30-15
Fabian et al. 2002 PASP 112, 1145
1994 / average 1997 / average
Lightcurve minimum Flare
31Thermal spectrum: observations of Perseus cluster of galaxies
Hitomi Col. 2016 Nature
3232Thermal plasma
Thermodynamic equilibrium occurs if Ne > 1014T 0.5e ∆E3
i j cm-3
For T=10 MK and H-like Iron, Ne >1027 cm-3
For T=0.1 MK and H-like Oxigen, Ne >1024 cm-3
These are very high densities occuring hardly anywhere outside stars
Astrophiscally important plasmas
Coronal/Nebular Ne <1016 cm-3
kTe ≈ Ip
* 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,
33Ionisation
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.
3434Equilibrium in thermal plasma
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
CZ,z−1 ionization rate, αZ,z recombination rate
nZ,z−1CZ,z−1 = nZ,z(αradZ,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 IENEI - codes
occurs e.g. in supernova remnants
3535 APEC simulated spectra for two different T(Chandra MEG+1)
Fe
XX
VI
Fe
XX
V C
aX
IX
SX
V
SiX
IVS
iXII
I
Mg
XII
Mg
XI
Ne
X
Ne
X
Ne
IX F
eX
VII
I
Fe
XV
II O
VII
I F
eX
VII
Fe
XV
II
OV
III
OV
II
NV
II
0
5
10
15
20
25
30
2 4 6 8 10 12 14 16 18 20 22 24 26Wavelength [A
o]
Co
un
ts/s
ec
Line emission dominates at kT=0.6 keV (T=7MK)
Strong continuum at kT=6 keV (T=70MK)
NB! Instrumental responce
No interstellar absorption
36High-Resolution X-ray Spectra
ζ OriO9.7I
XMM RGS
0.0
0.1
0.2
0.3
6 7 8 9 10 11 12 13 14 15 16 17 18 19 20 21 22 23 24 25 26 27 28 29 30 31 32 33 34Wavelength (A
o)
ζ PupO4I
XMM RGS
SiX
III
Mg
XII
Mg
XI
Ne
X
Ne
X
Ne
IX F
eX
VII
I
Fe
XV
II
OV
III
Fe
XV
II F
eX
VII
OV
III
OV
II
NV
II
NV
I
CV
I
0.0
0.1
0.2
* Overall spectral fitting plasma model, abundunces
* Line ratios TX (r), spatial distribution
* Line profiles velocity field, wind opacity
enhnaced N
? Find He-like ions ?
37Common X-ray diagnostics: lines of He-like ionsRatio of forbidden to intercombination line flux depends on ?
UV flux dilutes with
r-2
f/i ratio estimator for
distance where the
hot gas is located
Requires knowledge
of stellar UV field
OVIIGabriel & Jordan 1969
intercombination
resonance
forbidden
UV
38Comparing OVII in early and solar type stars
Chandra has 0.6 arcsec resolution
α CruCapella
λR λ I λF
OVII
0.021.5 21.6 21.7 21.8 21.9 22.0 22.1 22.2
Flu
x (C
ount
s/se
c/A
ngst
rom
)
Oskinova etal. in prep
B0.5IV+BV at d=98 pcX-ray brightest massive star on skySoft spectrum, narrow linescompare to solar type star
39Inner Shell Processes
Auger Ionization Additional source
of ionization in plasma
See Grotrian diagrmans in Kallman+ 04, ApJSS 155, 675