What can emission lines tell us? lecture 1 Grażyna Stasińska

What can emission lines tell us?

lecture 1

Grażyna Stasińska

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lecture 1

Grażyna Stasińska

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some basic literature dealing with the ionized ISM

books•Physical Processes in the Interstellar Medium, Spitzer, 1978 •Astrophysics of Gaseous Nebulae and Active Galactic Nuclei, Second Edition,

Osterbrock & Ferland, 2005

•Astrophysics of the Diffuse Universe, Dopita & Sutherland, 2003

lectures

•Stasinska 2002 astro-ph/0207500

« Abundance determinations in HII regions and planetary nebulae »

•Stasinska 2007 astro-ph « What can emission lines tell us? »

The mere presence of emission lines indicates

• the existence of gas

• eg emission line galaxies contain gas in large amount while galaxies emitting only a continuum with absorption features (such as elliptical galaxies) do not

• the existence of an ionizing agent (most emission lines come from ionized species)

• hot star(s)

• active nucleus

• (shocks) …

a gallery of nebular spectra

a g the galaxy

Te diagnostic

IFU data for the most metal-poor HII galaxy I Zw 18

Kehrig et al 2006

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[OIII] image H image

deep UV-FIR spectrum of the high excitation planetary nebula NGC 7027

Zhang et al 2005

observed

dereddened

high resolution spectrum of the PN NGC 6153 showing many recombination lines

Liu et al 2000

SPITZER IRS spectrum of the PN SMP83 in the LMCBernard-Salas et al 2004

XMM spectrum of the corona of α Cen Liefke & Schmitt 2006

line displacements tell about radial velocitiesand allow one to measure

• redshifts of galaxies

• dark matter mass in galaxies using PNe as test particles (eg Romanowsky et al 2003)

• internal motions in zones of line emission (eg line broadening in AGNs)

• expansion velocities

high resolution multislit spectroscopy of the PN NGC 7009 in the [NeIII] line showing expansion of the envelope

Wilson 1958

Basic mechanisms in ionized nebulae and emission line production

• ionisation and recombination processes• heating and cooling processes• line production mechanisms• about radiation transfer in nebulae• equilibrium versus out of equilibrium• the nebular physicist’s compendium

ionisation and recombination processes

ionization• Photoionization

• Collisions

• Charge exchange

Recombination• Radiative recombination

• Dielectronic recombination

• Charge exchange

heating and cooling processes

Heating

• Photoionization

• Collisional ionization

cooling• Free-free radiation

• Free-bound radiation

• Bound-bound radiation

what determines the ionic fractions and the temperature?

Ionization

Ionic fractions

Recombination

Heating

Electron temperature

Cooling

• Photoionization• Collisions• Charge exchange

• Radiative recombination• Dielectronic recombination• Charge exchange

• Photoionization (mainly H and He)• Collisional ionization

• Free-free radiation• Free-bound radiation• Bound-bound radiation (mainly O)

line production mechanisms

Recombination followed by cascade

(these lines are named with the recombined ion)

• H lines: Balmer …, Paschen etc …• He I lines (He I 5876…)• He II lines (He II 4686…)

Collisional excitation followed by radiative deexcitation

• Forbidden lines: [OIII]5007, [NII]6584• Semi-forbidden lines : CIII]1907 …• Resonance lines: CIV 1550, NV 1240, OVI 1035, SiIV 1400

Photoexcitation and fluorescence• Bowen lines: OIII 3133, 3444 (Bowen 1934)• Fe K line (probe of astrophysical black holes Fabian et al 2000)

• Each line can be produced by several processes, but usually only one dominates• H Ly is produced both by recombination and by collisional excitation

about radiation transfer in ionized nebulae

• Lyman continuum photons from the ionizing source

• ionizing photons produced by the nebula

• non-ionizing photons

Lyman continuum photons from the ionizing source

• They suffer geometrical dilution away from the source

• They suffer line-of-sight absorption (main absorbers H and He)

• The first ones to be absorbed are the ones with energies close to the ionization threshold (-3)

• => hardening of the ionizing radiation field in external zones

ionizing photons produced by the nebula

• those photons are emitted in all directions

The “on the spot approximation”

• assumes that all ground-state recombination photons are reabsorbed OTS• is justified for analytical order of magnitude estimations• but computed Te is incorrect by ~1000K-2000K (Gruenwald et al’s 3D code)

The “outward only approximation”• Radial outward only (Ferland’s Cloudy)• Full outward only (Stasinska’s PHOTO)

“complete” treatment• Traditional iterative way (Harrington, Rubin)• Using Monte-Carlo transfer (Ercolano’s Mocassin)

• resonance iine radiation is locally scattered many times• Can be treated with a “quasi on the spot” approximation (Ferland’s Cloudy)• Can be treated with an “local escape probability” approximation • Can be treated “exactly” (Dumont’s Titan)

non-ionizing photons (including lines emitted by the nebula)

In general • they escape freely (the optical thickness of the nebulae is small enough)• They can be attenuated by dust absorption

Exceptions• resonance lines

• which suffer scattering (eg H Ly ) and may be selectively destroyed by dust• FIR lines

• can suffer self-absorption as abs is larger than for optical lines (Rubin 1968)• but turbulent/expansion velocities favour their escape (Abel et al 2003)

equilibrium versus out of equilibrium

typical timescales for nebulae with n=103cm-3

Recombination time• trec= 1/(ne) 100 yr

Cooling time• tcool= (nkTe)/50 yr

Dynamical time• tdyn=R/vexp 2x104 yr for a single star HII region

Stellar evolution time • t* 5x106 yr for HII regions (=tMS)

1x104 yr for PNe (=tPAGB)

Most ionized nebulae are in ionization and thermal equilibrium

low density plasmas can be out of equilibrium

The nebular physicist’s compendium

• The Stromgren sphere

• The H luminosity and surface brightness

• What drives the electron temperature?

• What determines the ionization structure of a nebula?

• Why is the gas temperature roughly uniform in photoionized nebulae?

• Comments on energy losses

• Other comments on the gas temperature

• Comments on line intensities

QH [ph/sec] M*[M] Mion [M] Mion [M]

(n=102) (n=104)

PN 3x1047 .6 15 .15

O7 star 3x1048 30 150 1.5

starburst 3x1050 104 15000 150

the Strömgren radius

• In a homogeneous medium of density n and filling factor , the radius RS up to which gas is fully ionized is obtained from

QH = 4/3 RS3 n2 B(H) => RS [cm] = 9720 (QHn-2-1)1/3

• nb: The transition region between ionized and neutral gas is usually small:

maximum nebular mass that can be ionized• Mion= 4/3 RS

3 n mH => Mion [M] = 5 x 10-45 QH n-1

density bounded case:

LH = Mneb n (H)/ mH

LH is independent of QH

ionization bounded case:

LH = QH (H)/ B(H)

LH is a measure of QH

Hluminosity: LH = 4/3 R3 n2 (H)

H surface brightness: SH = LH / (4 R2)

ionization bounded case:

SH = A (QH n4 2)1/3

narrow slit spectra are of better quality for denser nebulae

density bounded case:

SH = B (Mneb n5 2)1/3

narrow slit spectra are of better quality for denser nebulae

• energy gains :

G = i, jni j i

where the i j are the gains per ion (photoionization and collisional ionization)

• energy losses :

L = i, jni j i

where the i j the losses per ion

(recombination and collisional excitation followed by photon emission)

• net energy gain :dE / dt = G - L

• If thermal equilibrium is achieved, the temperature is determined by: G = L

ionization equilibrium equation

between ni and ni+1, at a fraction f of the Stromgren radius RS

ni 4J i d= ni+1 nei

where the mean intensity of the radiation field in photons s-1

4J = QH g (T*) / r2

=> expression of the ionization state as a function of QH, n, f,T*:

ni+1 / ni = (QHn2)1/3 f-2 g(T*)

ionization parameter • definition: U = QH / (4R2 n c)

• quivalent expression:e U = A (QHn2)1/3

the ionization state is fully defined by the product QHn2 once T* is specified

the average ionization is higher for larger QH and larger n

in a given nebula, regions of higher n are more recombined

• Energy gained by photoionization of H at a distance r from the source

G = n(H°) 4J (h-h°) derg cm-3 s-1]

• Ionization equilibrium equation of H at distance r

n(H°) J d= n(H+) ne B(H)

• Substituting:

G = n(H+) ne B(H) < E >

with < E > = 4J (h-h°) d4J d

<E> is the mean energy gain per absorbed photon

<E> ≈ A T*

Energy gains due to photoionization of H are • independent of distance to the star• proportional to the star temperature

• The most important cooling process is collisionally excited line radiation

• For a given ion in a two-level approximation, the cooling rate is given by

Lcoll = n2 A21 h21 erg cm-3 s-1]

where n2 results from the equilibrium equation of levels 1 and 2 :

n1 ne q12 = n2 (A21 + ne q21)

• In the limit of small ne one has

Lcoll = n1 ne q12 h21

where q12 is the collisional excitation rate

q12 = 8.629 10-6 (1,2) /1 Te-0.5 exp (-E12/kTe)

note for « normal abundances » • the most important cooling ion is O++

• H and He have too high excitation potentials to be excited at “normal temperatures”

Cooling by collisional line excitation is more important for• abundant ions • lines corresponding to large• levels that can be easily attained at the temperature of the medium

Spatial variations of Te

• are mostly determined by • the mean energy of the absorbed photons• the populations of the main cooling ions

• are generally small • except at high metallicities

• in the O++ zone cooling is very efficient through emission of [OIII]52, 88 lines which have very low excitation potentials

• in the O+ zone the cooling efficiency is smaller (O+ has no low-lying levels)

photoionization models showing

the effect of metallicity

Stasinska 1978- - - Z < Z __ Z > Z

General dependence of Te with the defining properties of the nebulae

• for a given T*, Te as Z

• for a given Z, Te as T*

• for given T* , ionization state and Z, Te if n above ncrit

Temperature dependence of emission lines

Collisionally excited lines (CEL)

I CEL= n1(X) ne 8.629 10-6 (1,2) /1 Te-0.5 exp (-E12/kTe) h21

Recombination lines (RL)

I RL= n(X) ne Te- ( with ≈ 1)

Temperature dependence of line ratios

• Ratios RL / RL are almost independent of Te

• Ratios CEL(IR) / RL almost independent of Te

• Ratios CEL(opt or UV) / CEL(opt or UV) usually depend on Te

• Ratios CEL(opt) / RL strongly depend on Te

temperature dependence of emission lines

The example of lines emitted by O++

vthe Orion Nebulao’Dell http://vis.sdsc.edu/research/orion.html

volume visualization of the Orion Nebulao’Dell http://vis.sdsc.edu/research/orion.html

QuickTime™ et undécompresseur codec YUV420

a new view of the Orion Nebulaimages resolved in velocity and ionization

Garcia-Diaz & Henney 2006

What can emission lines tell us? lecture 1 Grażyna Stasińska

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