What can emission lines tell us? lecture 2 Grażyna Stasińska

What can emission lines tell us?

lecture 2

Grażyna Stasińska

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Diagnostics based on emission lines

• plasma diagnostics: electron temperature, density

• ionic and elemental abundances - direct methods

• elemental abundances - statistical methods

• estimation of the effective temperature of the ionizing star

- or of the effective hardness of the ionizing radiation field

• determining the star formation rate

• how to distinguish normal galaxies from AGN hosts?









[OIII]4363/5007 [SII]6731/6716



1S0

1D2

23P1

0

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2P

2D

4S

The most popular Te diagnostic The most popular ne diagnostic

Some plasma diagnostics in X-rays

Porquet & Dubau (2000)

He-like ions emit three main lines (n = 2 shell), which are close in wavelengths:

• resonance lines (called w), • intercombination lines (x + y),• forbidden lines (z).

the combination of the ratio of these lines can be used to derive

• the ionizing process (pure photoionized plasma or hybrid plasma)

• the electron density : R(ne) = z / x + y• the temperature : G(Te) =[(x + y) + z] /w

Plasma diagnostic diagrams

Plasma diagnostic diagram for the planetary nebula NGC 7027









The method for abundances from opticanl or IV lines

• Te and ne are obtained from plasma diagnostics

• Ionic abundance ratios are determined from line intensity ratios

eg: O++/H+ = ([OIII]5007/H) / ([OIII]5007(Te)/H(Te))

• Elemental abundance ratios are obtained

• either by adding all the observed ions

eg: O/H = O+/H+ + O++/H+ + O+++/H+ + …

• or by using ionization correction factors (icfs)

The method for abundances from IR lines

• as above except that Te is not needed

a note on ionization correction factors

Ionization correction factors based on ionization potentials

• a first approximation promoted by Torres-Peimbert & Peimbert 1977• but risky: eg (O+++ + ..)/O ≠ He++/He (although O++ and He+ have the same ionization potential :54.4 eV)

• there is nothing which empedes O++ ions to be present in the He++ zone

Ionization correction factors based on model grids may be risky too

• observations often pertain only to a small fraction of the object while grids usually consider entire nebulae• there is no robust formula to correct for He°

Cases when no icf is needed

• when all the expected ionization stages are observed however in this case beware of errors in determining ionic abundances

• from different spectral ranges • from lines extremely sensitive to Te (lines with high excitation potential as UV lines or transauroral lines)

a rough evaluation of Te-based methods

• the methods are easy to implement

• they depend on a very limited amount of assumptions

• error bars are relatively easy to estimate

• the abundances of the most important elements are expected to be correct (within

error bars)

• they are very close to abundances obtained from successful tailored

photoionization modelling

• from optical spectra abundances can be derived for He, N, O, Ne, S, Cl, Ar, Fe

• C is however a difficult subject

a case of failure of Te-based abundances: metal rich HII r.

Stasinska 2005

• with very large telescopes [OIII]4363/5007, [NII]5755/6584, [SIII]6312/9532 can be measured even at high metallicities (eg Bresolin et al 2005)

the problem

• at Z > Z strong Te gradients are predicted

• Te sensitive ratios strongly overestimate Te in the emitting zones

O/H is strongly biased !

the bias depends on

• what line is measured to derive Te

• what relation is adopted between T(O+) and T(O++)

T(O+)=T[NII]5755/6584

T(O++)=[T(O+)-3000]/0.7T(O+) =T[NII]5755/6584

T(O++) =T[OIII]4363/5007

a further problem to derive Te at high metallicity

contamination of collisionally excited lines (CELs) by recombination

• at low Te, CELs with high excitation energy such as [OII]7330 or [NII]5755 may be dominated by recombination

• this effect, very strong in the case of T[OII]3727/7330 is usually not well corrected for in the literature (one should use the Te representative of the zone emitting the recombination line to correct for it)

• a similar effect is likely to occur for T[SII]4070/6720









In many cases, the weak [OIII] 4363 or [NII]5755 lines are not available because

• the temperature is too low • the spectra are of low signal-to-noise • the data consist of narrow band images in the strongest lines only

Strong line methods to derive abundances • are statistical • have to be calibrated

Best known strong line methods: the ones based on oxygen lines

• Pagel et al 1979 used ([OII]+[OIII])/Has an indicator of O/H

• this method, la, has been calibrated many times

• Mc Gaugh 1994 refined the method to account for the ionization parameter U• Pilyugin (2000, 2001 ..., 2005) proposed the most sophisticated approach

Rationale of Mc Gaugh’s method

there are 4 independent strong line ratios• HH, [OII]/H, [OIII]/H, [NII] /H

there are 5 parameters determining them• C(H , <T*>, U, O/H, N/O

underlying hypothesis of the method• <T*> is related to O/H

• (this is expected statistically for giant HII regions)

the procedure• both O/H and U are derived simultaneously from• ([OII]+[OIII])/H, and [OIII]/[OII]

a problem• ([OII]+[OIII])/Hvs. O/H is double valued

a way out• [NII]/[OII] indicates whether O/H is high or low

because N/O increases with O/H (“astrophysical” argument)

McGaugh diagrams for the O23+ method

versus /

what lies behind the [OIII]5007/H vs O/H relation

Intensity ratio: [OIII]5007/H = A n(O++) / n(H+) Te0.5 exp (-28800/Te)

Thermal balance equation: n(H+) ne T* ≈ B ni j ne Te

-0.5 exp (- Eexc/Te)

• if 12 + log O/H << 8.2 • cooling is due to H Ly ,

• Te is independent of O/H

• [OIII]5007/H ≈ C T* O/H

• if 12 + log O/H > 9 • cooling is due to [OIII]52,88[OIII]5007/H ≈ C T* f(Te)

where f(Te) = Te exp (- 28800/Te)

• which decreases • with increasing O/H

An evaluation of strong line methods

Perez-Montero & Diaz 2005

• uses a data base of 367 objects with measured Te

• including some giant HII regions in the inner parts of galaxies (expected to be metal rich)

• but ignores the strong bias due to low Te evidenced by Stasinska 05

the strong line method recalibrated

Pilyugin Thuan 2005

upper branch calibration

(ie high O/H)

lower branch calibration

(ie low O/H)

• uses a data base of over 700 objects with measured Te

• including some giant HII regions in the inner parts of galaxies (expected to be metal rich)

• uses only Te-derived abundances

• but ignores the strong bias due to low Te evidenced by Stasinska 05

the last word on abundances from strong line methods is not said

more on strong line methods for Giant HII Regions

Stasinska 2006

Requirements for an ideal metallicity indicator

• should be single valued• should have a behaviour dominated by a well understood physical reason • should be unaffected by the presence of diffuse ionized gas• should be independent of chemical evolution

Looking for an ideal metallicity indicator

• data base of 670 objects in spirals, SDSS DR3 and BCDs galaxies with Te measured

• using P calibration of Pilyugin 2001 when Te is not measured

results: two new well behaved metallicity indicators

[ArIII]/[OIII] [SIII]/[OIII]

=0.23 =0.25

but the lines are only moderately strong ...

nb: all strong line methods will need recalibration when we undertand better the physics of metal-rich HII regions, (Stasinska 2005)

comparison of O/H from various metallicity indicators

[ArIII]/[OIII] vs [NII]/H

• larger dispersion (effect of N/O and ionization variations)• slight bias

[ArIII]/[OIII] ~[SIII]/[OIII]

• very tight correlation (as expected)• dispersion mostly from measurement errors









Estimation of T* by counting photons

Zanstra 1931

• TZH is obtained assuming that all stellar Lyc photons are absorbed by the nebula, from the observed stellar visual magnitude and the total nebular H flux

• for very hot stars (PN nuclei), one can also define TZHe using the He II 4686 flux as a measure of the number of photons with energies above 54.4 eV





notes on Zanstra-type methods and on the ionization of He

results from model computations with PHOTO

a . He I 5876 / H measures T* only in a small range (T* < 40 kK)

due to competition between H° and He+ to absorb photons with energies > 54.4 eV

c . HeII 4868 / Hsaturates at T* > 150 kK

c . HeII 4868 / Hdepends on U at T* > 100 kK

dependence on He/H

c . HeII 4868 / H does not depend on He/H

e . HeII 4868 / He I 5876 depends on He/H

not considered in empirical methods

f . the H+ and He++ zones may have different Te

___ U=10-2 He/H=0.1___ U=10-3 He/H=0.1___ U=10-2 He/H=0.15

a

c

b

d

e f

T* from observed ionization structure

Kunze et al 1996The ionization structure depends on T*-> line ratios of two successive ions measure T*

but the ionization structure also depends on U !!!





Morisset 2004determination of T* using a full grid of atmospheres with WM-basic and taking into account T*, U and metallicity

(SIV/SIII) / (NeIII/NeII) vs NeIII/NeII

T*

T* from energy-balance methods

Stoy 1931 Stasinska 1980

L( CEL) / L(H) = f(T*) Te is a function of O/H and T*

calibration by Preite-Martinez & Pottasch 83











star formation rate

techiques• UV continuum, FIR continuum, recombination lines, forbidden lines...• each technique requires a calibration usually done with evolutionary stellar synthesis models

basic parameters• metallicity (Z)• star formation history (SFH)• description of the IMF• stellar evolutionary tracks• stellar model atmospheres

see reviews by Kennicutt 1998 and Schaerer 1999

star formation rate using L(H)

Kennicutt 1998

SFR [M yr -1] = 7.910-42 L(H)[erg s-1] A(H) / f

• where A(H) is the extinction

• f is the fraction of Lyc photons absorbed by H


sont requis pour visionner cette image. IMF Mup Z/Z _____ Salpeter 100 1.......... Salpeter 100 .05_ _ _ _ Salpeter 100 2_ . _ . Salpeter 30 1------ Scalo 100 1

Scharer 1999

the SFR from L H strongly depends on assumed parameters for the stellar population

temporal evolution of models with cst SFR

star formation rate using [OII]

advantage of [OII]• is seen in a broad redshift range, rather used at large redshifts (~ 1)

caution about [OII]• calibrations by different authors differ strongly (see Kennicutt 1998)• [OII]/H is expected to vary with metallicity and U • [OII] can be produced by ionization by an active galactic nucleus AGN and not by

stars

exemple of observed dispersion in [OII]/H


sont requis pour visionner cette image.data from subsample of SDSS DR3

• normal star forming galaxies• AGN host galaxies• hybrid








• how to distinguish normal star forming galaxies from AGN hosts?

Segregation of emission line objects in emission-line ratio diagrams

PNe

AGNs

GHRs

The BPT diagram

Baldwin, Phillips, Terlevich 1981

e i of the diagram

Interpretation

photons from PNe and AGNs are harder than those from massive stars that power GHRs

they provide more heating

collisionally excited lines will be brighter than in the case of ionization by massive stars only

[OII

I]/H

[NII]/H

The next step

Veilleux & Osterbrock 1987

• more diagrams, more points GHRs form a sequence in the [OIII]/H vs [NII]/Hand [OIII]/H vs [SII]/H

• comparison with sequences of photoionization models

[OIII]/H vs [NII]/H[OIII]/H vs [SII]/H[OIII]/H vs [OI]/H

the Sloan Digital Sky Survey revolution

Kauffmann et al 2003

• spectra of 100 000 galaxies

• subtraction of stellar continua obtained by population synthesis

galaxies hosting AGNs also form a sequence!

galaxies in the BPT diagramnow remind the wings of a seagull

modelling of the upper envelope of the left wing

Stasinska Cid Fernandes Mateus Sodre Vale Asari 2006

motivation• previous dividing lines were “too generous” for NSF galaxies

the model (uses Starburst99 & PHOTO)• constant star formation• abundance ratios taken from Izotov et al 2006

result • U decreases az Z increases• [OI] and [SII] lines less well fitted (because of 1-zone model)

of the 4 diagrams, the [OIII]/H vs [NII]/H is the best to distinguish NFSg and AGN hosts

can one distinguish AGN hosts and NSF galaxies with their [NII]/Honly ?

distinguishing AGN hosts and NSF galaxies using only [NII]/H

• feasible

• allows one to consider more galaxies of the initial sample (intensities of [OIII] and H not needed)

• allows one to see relations with another parameter (here D4000)

AGN

NSF

all

hybrid

end of lecture 2

What can emission lines tell us? lecture 2 Grażyna Stasińska

Documents

te lines

effective temperature

ionizing star

o h o h o h

electron temperature

uv lines

resonance lines

forbidden lines