What can emission lines tell us? lecture 4 Grażyna Stasińska

1. What can emission lines tell us?lecture 4Grażyna Stasińska

2. Some pending questions and some strategies to solve them • aperture correction • dereddening • underlying stellar absorption • escape of ionizing radiation • dust • temperature fluctuations • chemical inhomogeneities • the role of shocks

3. Aperture correction • When the studied objects are more extended than the observing beam • aperture correction is needed if the observing beams are not the same for all wavelenghts. For example • combining ground optical and UV spectra from IUE or HST • combining FIR measurements with optical measurements • Aperture correction are usually done • using line ratios that have a known intrinsic value (e.g. HeII 1640 / He II 4686) • using ratios of apertures • Such procedures bear uncertainties • they do not take into account the ionization stratification of the nebulae • The best way to do this (collaboration Stasinska, Morisset, Simon-Diaz ...2006) • build a photoionization model reproducing the observed H surface brightness distr. • compute the intensities through each observing slit • compare the observed intensities with the model intensity throughappropriate slit

4. Correction for dust extinction • The method: • The “logarithmic extinction at Hb”, C, is derived from the observed Ha/Hb ratio by comparing it to the theoretical one for case B recombination assuming an extinction law f(l) C = [ log (FHa / FHb)B - log (FHa/ FHb)obs ] / (fa- fb) • Emission line ratios are then dereddened using the formula log (Fl1 / Fl2)corrB = log (Fl1 / Fl2)obs + C (fl1- fl2) • Problems: • The “extinction law” is not universal • The intrinsic Ha/Hb ratio may be different from the theoretical case B (collisional excitation, case C) • If some dust is mixed with the ionized gas and strongly contributes to the extinction, no “extinction law” applies

5. the “extinction law” is not universal • the canonical extinction law corresponds to RV=AV / E(B-V) = 3.2 • in Orion RV = 5.5 • towards the Galactic bulge, RV ~ 2.5(eg Stasinska et al 1994) • larger values of RV are found for lines of sight crossing molecular clouds where dust grains are expected to be larger Extinction laws corresponding to various values of RV=AV / E(B-V) .1989 Histogram of RV for 95 galactic O stars Patriarchi et al 2001

6. checks on the reddening correction • Before dereddening check that conditions for case B are likely satisfied • If not, consider building a photoionization model with a code that treats the H atom correctly, and redden the resulting the emission lines to fit the observed Balmer decrement. • If case B is relevant for the object under study • check that, after reddening correction , Hg / Hb is close to the case B value • If not, [OIII]4363/5007 is likely to be in error by the same amount as Hg/Hb differs from the case B value • If many Balmer lines are measured with good accuracy • rather than using an “extinction law” • fit the observed H/Hb Hg/H to the theoretical case B values • this method is valid (in first approximation) also if dust is mixed with the HII gas • but, of course, it does not allow to derive C(H)

7. underlying stellar absorption • Some nebular spectra may contain a lot of light from stars • giant HII regions • small size PNe • nuclei of galaxies • entire galaxies • This light may contaminate the emission lines • This is an important problem • if the emission lines are faint • eg faint nuclei of galaxies • if one desires a high accuracy in line measurements • eg determination of the pregalactic helium abundance

8. underlying stellar absorption • Possible ways out • correct for reddening and stellar absorption at the same time • in order to obtain the correct Balmer decrement (eg Izotov et al) • observe with good spectral resolution • stellar lines are generally broader than emission lines • use « template » continuum spectra • this was common practise in the last decade to study faint nuclear emission regions • do a model for the stellar light and subtract it from the observed spectrum • this is now routinely done by many groups studying galaxy spectra and using stellar population synthesis techniques • Cid Fernandes et al, Tremonti et al etc...

9. escape of ionizing radiation • most nebular studies assume that the nebulae are ionization-bounded • modelling of planetary nebulae • estimate of T* via the Zanstra method • estimation of star formation rates in galaxies etc ... • there is growing evidence that many nebulae are density bounded • in at least some directions • planetary nebulae • giant HII regions • high redshift galaxies

10. the effects of dust • evidence for the presence of dust mixed with ionized gas • the effect of dust on the ionization structure • the effect of dust on the nebular thermal balance • the effect of dust on resonance lines • relevance of elemental abundances in case of depletion

11. dust coexists with ionized gas • Evidence from IR spectroscopy • strong IR continuum due to dust heated by stellar radiation • emission features attributed to silicate or carbon-based particles • the dust temperature indicates that grains are not limited to the neutral outskirts of the nebulae • Evidence from IR imaging • dust emission is seen in the ionized region (eg Graham et al 1993) • Evidence from optical spectroscopy • Refractory elements (Fe, Mg, Si,Ca) are largely depleted in HII regions and PN, indicating that dust is intimately mixed with ionized gas IR flux distribution in the PN NGC 3918Harrington et al 1988

12. the dust-to-gas mass ratio in ionized nebulae • Its estimate depends on the adopted grain size distribution • Very different values are quoted • in HII regions • md/mg = 10-4 - 10-3 Hoare et al 1991 • in planetary nebulae • md/mg = 10-4 - 10-2Natta & Panagia 1981, Stasinska & Szczerba 1999 • an extreme case: the dusty PN in the globular cluster M22 • md/mg = 0.4 Borkowski & Harrington 1991

13. dust and the ionization structure • the optical depth of dust in the EUV can be significant • D= sD (nD/nH) nH R ≈ 0.6 (U/10-3) if (nD/nH) has the local ISM value • absorption of ionizing photons by dust reduces • the intrinsic Hbluminosity and the ionization parameter effect of the wavelength dependence of sAanestad 1989 • speaks at 700A • dust will absorb H-ionizing photons more efficiently than He-ionizing photons • the ionization level increases with respect to dust-free case • in Orion the He+ zone merge with the H+ zone Baldwin et al 1991

14. effects of the presence of grains on the thermal balance • consequences of depletion • coolants such as Mg, Si, Fe (also C to a lesser extent) are partly tied up in grains • collisional line cooling is therefore reduced • especially in outer zones where Mg, Si and Fe cooling is most efficient • and Te is enhanced relative to the dust-free case • gas-grain collisions • are a cooling factor for the gas • photolectric effect on dust grains • electrons ejected from grains by photoelectric effect heat the gas Spitzer 1948

15. rdust- and H- heating • thermal gains of the gas due to photolectric effect on dust : GD= nD 4pJn/ (hn) an (D)(hn-E°) dn • thermal gains of the gas due to H ionization: GH = A n(H+)ne • ratio of dust- to H-heating : GD / GH= nD 4pJn/ (hn) an (D)(hn-E°) dnn(H+)ne) GD / GHa nD  nH U • heating by photoelectric effect on dust grains becomes relatively important • when dust-to-gas ratio is high • when ionization parameter is high

16. the effects of grains on Te • dust can be important for the thermal balance of the gas Baldwin, Ferland, Martin et al 1991 Fraction of total heating due to photoelectric effect and fraction of total cooling due to grain-gas collisions in the Orion nebula Baldwin et al 1991

17. the effects of grains on Te Heating by dust is more efficient when a population of small grains(10A) is presentDopita & Sutherland 2000 heating contributions from photoionization of small grains __ large grains ---- hydrogen __ _________________________________________ Te as a function of fractional radius

18. the effects of grains on Te small grains can give rise to important “temperature fluctuations”in filamentary or knotty nebulae Stasinska & Szczerba 2001 ___Te _____ ne ___Te _____ ne filamentary dust free model filamentary model with small dust grains if dielectronic recombination is enhanced at high Te, small grains could also perhaps help solving the recombination line conundrum

19. the effects of dust on resonance lines • Attenuation of resonance lines • resonance lines experience important scattering in the nebulae • can be selectively attenuated by dust absorption compared to other lines • should not be used for abundance determinations without caution Attenuation of resonance lines by dust in NGC 3918 Harrington et al 1988 • Departure from case B • destruction of H Lyman lines by dust absorption 100% conversion of high-n Lyman lines into H Lya and Balmer lines (the case B assumption is no more verified) Cota & Ferland 1988

20. relevance of elemental abundances in case of depletion • Mg, Si, Fe, Ni, Ca • these elements can be almost entirely in the form of grains • their abundances in the gas phase cannot be easily used as indicators • of chemical evolution of galaxies • or nuclear processes in PN progenitors • C • can be heavily depleted by carbon-based grains (graphite, PAHs...) • O • can be slightly depleted (20% for Orion, estimated from depletion pattern of metals, Esteban et al 1998) • He, Ne, Ar • rare gases, do not combine into grains

21. Temperature fluctuations • Were postulated by Peimbert (1967) to explain discrepancies betwen Tefrom various diagnostics • Peimbert’s formalism • Do temperature fluctuations exist? • see reviews by Peimbert 1995, 2001, Mathis 1997, Stasinska 1998, Esteban 1998, 2001 • Numerous studies point towards t2 ~ 0.04 • But little direct evidence is seen

22. example of indirect evidence for t2 ≠ 0 In planetary nebulae Te from Balmer discontinuity is smaller than T[OIII] 4363/5007 (Liu & Danziger 1993) t2 ~ 0.04 is a representative value

23. If Te fluctuations exist they affect abundance determinations • e.g. abundance derived in M8 (Peimbert et al 1993) using Taylor series expansion of the line emissivites for various values of t2 Abundances derived from optical forbidden lines with respect to H are underestimated when ignoring t2 Abundance ratios like N/O or C/O are less affected Abundances derived from recombination or FIR lines are not affected

24. visualisation of the Peimbert formalism on a two-zone toy model f = N2n2V2 / N1n1V1 in this case the values of T0 and t2 are simply V2n2 N2 T2 V1n1 N1 T1 volume electron density ionic density temperature • Even in such a simple model , the temperature distribution requires 3 parameters to be defined (T0, t2, f), not 2 (T0, t2 ) • f >> 1 may represent a photoionized nebula with small shock-heated regions of veryhigh T1 • f << 1 may represent a nebula with high metallicity clumps of low T2 t2 __0 __0.02 __0.04 __0.06 variations of T1 and T2 with f for fixed t2 and fixed T0

25. effect of t2 on derived abundances t2t __0 __0.02 __0.04 __ 0.06 • if t2 is not accounted for • O++5007 is underestimated because T4363/5007 overestimates the temperature characteristic of the [OIII]5007 emission • the bias depends on T0 and f

26. does the Peimbert formalism give correct abundances ? t2 __0 __0.02 __0.04 __ 0.06 • t2obs≠ t2 • the result depends on f • Even if the computed t2 is not equal to the true one, does the Peimbert formalism lead to accurate abundances? • not quite • the bias depends on T0 and f

27. frequent misuse of the Peimbert’s formalism • from expansion of the emission coefficient in Taylor series, and integrating over the observed volumes, one obtains: • from which T0 and t2are obtained • but, except if O is entirely in the form of O++ in the nebula, • t2(H+)≠ t2(O++) • T0(H+) ≠ T0(O++)

28. Visualisation of energy requirements The simplest example: for t2 = 0.04 and T0 = 10000K, f =1 implies T1 = 12000K and T2 = 8000K blacklog of heating rate in arbitrary units red: log of cooling rate in the O++ zone By shifting the heating curve up and down one understands how Te varies with energy input t2 = 0.04 requires D log G = 0.3, ie a factor 2 difference in heating rates between regions 1 and 2 !

29. What fluctuates? • Te ? • Natural gradients in photoionized nebulae are small • except at high metallicities • (Stasinska 1980, Garnett 1992, Kingdon & Ferland 1995, Perez 1997) • Ne ? • In high density clumps collisional dexcitation increases Te with respect to the ambient medium Kholtygin 1998, Mathis et al. 1998 • (this is not sufficient to explain t2 ~ 0.04) • densities above 105cm-3 boost [OIII] 4363/5007 (Viegas & Clegg 1994) • but there is no evidence of such high densities in the O++ zones

30. What fluctuates? • Ni ? • Te is lower in C-rich zones(Torres-Peimbert et al 1990) • The O++ discrepancy between CEL and ORL requires the existence of O-rich zones (Stasinska 1998, Liu et al 2000, 2001, Péquignot 2001)

31. is photoionization the only heating source in photoionized nebulae? • In a number of nebulae, classical photoionization models produce T[OIII] lower than observed • Giant HII regions:Campbell 1990, Garcia-Vargas et al 1997, Stasinska & Schaerer 1999, Luridiana et al 1999, Luridiana & Peimbert 2001 • PNe: Peña et al 1998 • Additional energy sources have been proposed: • Shocks(Peimbert et al 1991) • Conduction fronts (Maciejewski et al 1996)

32. the ORL /CEL discrepancy • Expected properties of optical recombination lines (ORLs) • their emissivity is roughly proportional to Te-1 • they should give correct abundances with respect to H • even in presence of temperature fluctuations • ORL abundances versus CEL (collisionally excited lines) abundances • ORL abundances are larger than CEL abundances by important factors • Wyse 1947, Peimbert et al 1993, Liu et al 1995 (O) , Kaler 1986 (C) • Esteban et al 1998, Liu et al 2000, 2001 (C,N,O) C++1909 / O++5007 versus C++4267 /O++5007 in planetary nebulae Compilation Rola & Stasinska 1994

33. ORL versus CEL abundances • ionic abundances in the planetary nebula NGC 6153Liu et al 2000

34. invoked causes of ORL-CEL discrepancy Faintness of the ORLs biased measurements flux calibration is difficult over a large dynamical range they may suffer from blends Heavy element recombination coefficients are not reliable Temperature fluctuations Density condensations no (from high S/N spectroscopy) no: [OIII]4931/[ [OIII]4959 agrees with theory: 4 10-4Mathis & Liu 1999 no (from echelle spectra ) have been recomputed with the R-matrix method Storey 1994 ORL abundances from numerous transitions are in agreement t2 explaining ORL-CEL discrepancy >> t2 explaining Te[OIII] -Te(BJ) IR-CEL abundances are consistent with optical -CEL abundances Liu 2000,2001 no (high order Balmer lines)Liu 2000...

35. invoked causes of ORL-CEL discrepancy Faintness of the ORLs biased measurements flux calibration is difficult over a large dynamical range they may suffer from blends Heavy element recombination coefficients are not reliable Temperature fluctuations Density condensations no (from high S/N spectroscopy) no: [OIII]4931/[ [OIII]4959 agrees with theory: 4 10-4Mathis & Liu 1999 no (from echelle spectra ) have been recomputed with the R-matrix method Storey 1994 ORL abundances from numerous transitions are in agreement t2 explaining ORL-CEL discrepancy >> t2 explaining Te[OIII] -Te(BJ) IR-CEL abundances are consistent with optical -CEL abundances Liu 2000,2001 no (high order Balmer lines)Liu 2000... Recombination coefficients computed so far do not include dielectronic recombination for n > 10 which is likely to be efficient at Te > 20 kK Chemical inhomogeneities:they require super metal rich inclusions with solar C/N/O/Ne Liu 2000, Tsamis 2003

36. possible origin for chemical inhomogeneities • in planetary nebulae • ejecta from the central star ? • photoevaporating planetesimals or planet debris ?

37. possible origin for chemical inhomogeneities galactic fountainwith a spray superbubble galactic disk no more supernova galactic disk t=40-100 Myr supershell supernova giant HII region warm oxygen-rich cloudlets neutral oxygen-rich droplets t=1-40 Myr cold oxygen-rich cloudlets ionized oxygen-rich droplets galactic disk ionized ISM cold oxygen-rich droplets fully mixed HII gas t=0 galactic disk t=100 Myr new ionizing stars • HII regions droplets containing matter from supernova ejecta Tenorio Tagle 1996 Stasinska Tenorio Tagle Rodriguez Henney 2007

38. The t2 problem and the ORL/CEL discrepancy are still a subject of debate

39. The role of shocks • The effects of shocks on emission line spectra as compared to stellar ionization • high densities due to gas compression • possible presence of highly ionized species (He++) • existence of important warm low-ionization zone (emitting [OI], [SII], [OII] lines • higher Te , [OIII]4363/5007 enhanced • very high Te close to the shock, producing X-ray radiation Effects of X-ray ionization as compared to stellar ionization 2, 3, 4 Effect of gas compression without shock • local compression of gas lowers U • effects 1, 3 are produced (without the need of shock heating) NB • what is often attributed to shocks may actually be only due to compression • photoionization is much more efficient than shocks to ionize gas

40. What can emission lines tell us? a lot !