Atomic Spectra in AstrophysicsAtomic Spectra in Astrophysics

Potsdam University : Wi 2016-17 : Dr. Lidia Oskinovalida@astro.physik.uni-potsdam.de

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