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What can emission lines tell us? lecture 2 Grażyna Stasińska

What can emission lines tell us? lecture 2 Grażyna Stasińska. Diagnostics based on emission lines. plasma diagnostics: electron temperature, density ionic and elemental abundances - direct methods elemental abundances - statistical methods

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What can emission lines tell us? lecture 2 Grażyna Stasińska

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  1. What can emission lines tell us?lecture 2Grażyna Stasińska

  2. 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?

  3. 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?

  4. 1S0 1D2 2 3P1 0 2P 2D 4S The most popular Te diagnostic The most popular ne diagnostic [OIII]4363/5007 [SII]6731/6716

  5. 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

  6. Plasma diagnostic diagrams Plasma diagnostic diagram for the planetary nebula NGC 7027

  7. 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?

  8. 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/Hb) / (e[OIII]5007(Te)/eHb(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

  9. 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)

  10. 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

  11. 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[OIII]4363/5007 T(O+)=T[NII]5755/6584 T(O++)=[T(O+)-3000]/0.7

  12. 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

  13. 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?

  14. 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 • arestatistical • have to be calibrated • Best known strong line methods: the ones based on oxygen lines • Pagel et al 1979used ([OII]+[OIII])/Has 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

  15. Rationale of Mc Gaugh’s method • there are 4 independent strong line ratios • HH, [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])/Hb, and [OIII]/[OII] • a problem • ([OII]+[OIII])/Hvs. 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)

  16. McGaugh diagrams for the O23+ method  versus /

  17. what lies behind the [OIII]5007/Hb vs O/H relation Intensity ratio:[OIII]5007/Hb = A n(O++) / n(H+) Te0.5 exp (-28800/Te) Thermal balance equation: n(H+) ne T* ≈ B ni jne Te-0.5 exp (- Eexc/Te) • if 12 + log O/H << 8.2 • cooling is due to H Ly a, • Te is independent of O/H • [OIII]5007/Hb ≈ C T* O/H • if 12 + log O/H > 9 • cooling is due to [OIII]52,88m [OIII]5007/Hb ≈ C T* f(Te) where f(Te) = Te exp (- 28800/Te) • which decreases • with increasing O/H

  18. 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) • butignores the strong bias due to low Teevidenced by Stasinska 05

  19. 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 • butignores the strong bias due to low Te evidenced byStasinska 05 the last word on abundances from strong line methods is not said

  20. more on strong line methods for Giant HII RegionsStasinska 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

  21. results: two new well behaved metallicity indicators • [ArIII]/[OIII] [SIII]/[OIII] • s=0.23s=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)

  22. comparison of O/H from various metallicity indicators • [ArIII]/[OIII] ~[SIII]/[OIII] • very tight correlation (as expected) • dispersion mostly from measurement errors • [ArIII]/[OIII] vs [NII]/Ha • larger dispersion (effect of N/O and ionization variations) • slight bias

  23. 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?

  24. 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 TZHeusing the He II 4686 flux as a measure of the number of photons with energies above 54.4 eV

  25. 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 / H saturates at T* > 150 kK c. HeII 4868 / Hdepends 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 b a ___ U=10-2 He/H=0.1 ___ U=10-3 He/H=0.1 ___ U=10-2 He/H=0.15 d c e f

  26. T* from observed ionization structure • Kunze et al 1996 • The ionization structure depends on T* • -> line ratios of two successive ions measure T* • but theionization structure also depends on U !!! T* (SIV/SIII) / (NeIII/NeII) vs NeIII/NeII Morisset 2004 determination of T* using a full grid of atmospheres with WM-basic and taking into account T*, U and metallicity

  27. 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

  28. 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?

  29. 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

  30. star formation rate using L(H) • Kennicutt 1998 SFR [M yr -1] = 7.910-42L(H)[erg s-1] A(H) / f • where A(H)is the extinction • f is the fraction of Lyc photons absorbed by H Scharer 1999 the SFR from L H strongly depends on assumed parameters for the stellar population temporal evolution of models with cst SFR IMF Mup Z/Z  _____ Salpeter 100 1 .......... Salpeter 100 .05 _ _ _ _ Salpeter 100 2 _ . _ . Salpeter 30 1 ------ Scalo 100 1

  31. 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 • data from subsample of SDSS DR3 • normal star forming galaxies • AGN host galaxies • hybrid

  32. 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 star forming galaxies from AGN hosts?

  33. Segregation of emission line objects in emission-line ratio diagrams • 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 [OIII]/Hb PNe AGNs GHRs [NII]/Ha

  34. The next step • Veilleux & Osterbrock 1987 • more diagrams, more points GHRs form a sequence in the [OIII]/Hb vs [NII]/Haand [OIII]/Hb vs [SII]/Ha • comparison with sequences of photoionization models [OIII]/Hb vs [NII]/Ha [OIII]/Hb vs [SII]/Ha [OIII]/Hb vs [OI]/Ha

  35. 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 diagram now remind the wings of a seagull

  36. modelling of the upper envelope of the left wingStasinska 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

  37. can one distinguish AGN hosts and NSF galaxies with their [NII]/Haonly ?

  38. distinguishing AGN hosts and NSF galaxies using only [NII]/Ha NSF hybrid • feasible • allows one to consider more galaxies of the initial sample (intensities of [OIII] and Hb not needed) • allows one to see relations with another parameter (here D4000) AGN all

  39. end of lecture 2

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