1. TEMPERATURE STRUCTURE OF GASEOUS NEBULAE AND CHEMICAL ABUNDANCES M. Peimbert  C.R. O’Dell  A. Peimbert  V. Luridiana  C. Esteban  J. García-Rojas  L. Carigi  F. Bresolin  M.T. Ruiz  A.R. López-Sánchez Lake Geneva, Wisconsin, April 2007Microstructures in the ISM: Bob O’Dell 70th birthday

2. OUTLINE Why is this problem important? Definitions T [O III ], T (Balmer),T (O II ), T (C II ) Which is the cause of temperature variations The Orion nebula and microstructures The Orion nebula and the solar abundances Calibration of the R 23 method The primordial helium abundance Conclusions

3. Why is the problem of temperature variations important? Physical conditions of gaseous nebulae Abundances in H II regions and PNe Solar abundances Galactic chemical evolution Primordial helium abundance, Y P Metal content and chemical evolution of the universe

4. Temperature Structure T e ( 4363/5007 ) = T 0 [ 1 + ( 90800/T 0 -3 ) t 2 /2 ] T e ( Bac/H  ) = T 0 ( 1 – 1.70 t 2 ) T e ( 4649/5007 ) = f 1 ( T 0, t 2 ) T 0 = t 2 =  T e N e N i dV  N e N i dV  ( T e - T 0 ) 2 N e N i dV T 0 2  N e N i dV T e ( He lines ) = T 0 ( 1 – k t 2 ) k~1.8 T e (4267/1909) = f 2 (T 0, t 2 )

6. Ups and downs of t 2 March 2007

7. How Important Are Temperature Variations? Photoionization homogeneous models predict values of t 2 in the 0.003 to 0.03 range, with typical values around 0.01 Observational values of t 2 are in the 0.00 to 0.09 range with typical values around 0.03 Typical ratios between the abundances derived from permitted lines and forbidden lines are in the 2 to 3 range (O, C, N, Ne), the so called abundance difference factor, ADF By adopting t 2 values different from 0.00 it is possible to reconcile the abundances derived from forbidden lines with those derived from permitted lines

8. Presence of Temperature Variations There are temperature variations that can not be explained by chemically homogeneous photoionization models The sources of these variations can be many and a specific model has to be made for each nebula The abundances derived from recombination lines are almost unaffected by temperature variations The abundances derived from collisionally excited lines, under the assumption of constant temperature, typically underestimate the abundances relative to hydrogen by a factor of 2 to 3

9. Liu & Danziger 1993 Balmer vs. [O III ] Temperatures

10. N(C ++ ) from Recombination Lines vs. N(C ++ ) from Forbidden Lines Peimbert, Luridiana, & Torres-Peimbert 1995

11. Recombination to Forbidden O ++ ratios (log ADF) vs. [O III ] – Balmer Temperatures Liu et al. 2001

12. What causes Temperature Variations? Deposition of mechanical energy Chemical inhomogeneities Presence of WR Stars Dust heating Time dependent ionization Density variations Deposition of magnetic energy Shadowed regions

13. Microstructures and t 2 in the Orion Nebula O´Dell et al. 2003 [O III] 5007 image

14. Based on HST data O´Dell et al. 2003 We derived 1,500,000 T C [ 4363 / 5007 ] columnar values

15. Noise vs. True Temperature Variations O´Dell et al. 2003 The face of the nebula is mottled with small scale variations in T C with angular dimensions of about 10” (~0.02 pc) and amplitudes of 400 K

16. Histogram of T C [ 4363 / 5007 ] O´Dell et al. 2003 We obtained a t 2 A (O ++ )=0.008 across the face of the nebula values

17. Small Scale Ionization Structure O´Dell et al. 2003 [ N II ] / H I [ O III ] / H I

18. t 2 in the Orion Nebula From HST narrow filter images: –t 2 A (O ++ )=0.008 From a very small region of Orion Esteban et al. (2004) estimated: –t 2 sr (O ++ )=0.020±0.002 from O II and [O III ] –t 2 sr (H + )=0.022±0.002 from T(He I ) vs. T([O II ]+[O III ]) O´Dell et al. estimated: t 2 Whole Object ( H + ) =0.028±0.006

19. The Low Te Regions behind Clumps within the Ionized Gas Proplyds  Shadows, as long as 0.2 pc, covering 0.5% of the field of view contribute with 0.0093 to t 2 (O + ) Neutral High Density Clumps  Shadows, as long as 0.025 pc, covering about 1/250 of the volume contribute with 0.0016 < t 2 (O + ) < 0.0075

20. Neutral High Density Clumps O´Dell et al. 2003

21. Different Components of t 2 The total value of t 2 (H + ) has to consider both the O + and the O ++ regions

22. Chemically inhomogeneous H II regions: Pros + In favor is the study of the N excess in NGC 5253 studied by Angel Sanchez-Lopez et al.(2007). who found from the O II and C II recombination lines t 2 values of 0.052 and 0.072, and that the excess N is due to pollution by massive WR stars + Also in favor is the study by Tsamis and Pequignot (2005) that produced a chemically inhomogeneous model of 30 Doradus that also reproduces the observed line intensities of the forbidden and permitted O, C, and N lines

23. Chemically inhomogeneous H II regions : Objections – One of the problems with the model of TP is that the excess abundance of O in the clumps is of a factor of 8, and that it requires an excess of 14 for C. Models of chemical evolution of irregular galaxies by Carigi, Colin, and Peimbert predict that 64% of the C is due to IMS and 36% to massive stars. Therefore for an excess of a factor of 8 in O the TP model should predict an excess of only a factor of 3 for C – An even larger discrepancy between the model by TP is present in the case of N for which ~80% is due to IMS – The small dispersion in abundances of H II regions in irregular galaxies and in the abundance gradient in our galaxy are against this idea

24. Chemically inhomogeneous H II regions: Implications The two phases chemically inhomogeneous model by Tsamis and Pequignot and the observations of 30 Doradus of A. Peimbert give: 12 + log O/H = 8.45, while the chemically homogeneous model gives 8.33 for t 2 = 0.000 and 8.54 for t 2 = 0.033 Therefore the TP model is closer to the abundances given by the O II lines than to those given by the [O III ] lines and the T[O III ] temperature

25. Orion and the Galactic gradient vs. the Solar abundances Galactic abundances from collisionally excited lines (assuming t 2 =0.00) are almost a factor of 2 lower than those we found from solar studies and Galactic chemical evolution models –Pilyugin et. al (2003, A&A, 401, 557) find O/H = 8.52 dex in the solar vicinity –Deharveng et. al (2000, MNRAS, 311, 329) find O/H = 8.53 dex in the solar vicinity

26. Galactic Abundance Gradients Esteban et al.ApJ, 2005

27. Determinations from Recombination Lines (Equivalent to t 2 ≠0.00 ) We have found the O/H abundance as a function of Galactocentric distance. From observations of H II Regions we found a solar vicinity abundance of 8.79 dex with a gradient of -0.044 dex kpc -1 (Esteban et. al, 2005, ApJ, 618, 95) –The slope of this gradient is similar to those derived from [O III ] and t 2 =0.00 This value is consistent with the O/H = 8.66 dex Solar value derived by Asplund et al. (2005), and with Galactic chemical evolution models that estimate that, in the 4.6 Gy since the Sun was formed, there has been an 0.13 dex increase in oxygen abundance of the ISM (Carigi et al. 2005, ApJ, 623, 213)

28. Additional Support for a Higher O/H Initial Solar Value There are two results that indicate that the initial solar abundance was higher than the one adopted by Carigi et al., and that correspondingly the ISM t 2 values are even higher than those derived by Esteban et al. 2005 1)Estimates of the gravitational settling indicate that the original oxygen solar abundance was higher by about 0.05 dex than the present photospheric one, e. g. Piersanti et al. (2007), Bahcall et al. (2006), Basu & Antia (2004)… 2)There is a strong discrepancy between the Asplund et al. 2005 photospheric abundances and the solar interior ones determined from helioseismic measurements that amounts to ~ 0.1 dex

29. Determination of O/H abundances in distant extragalactic H II regions: Calibration of the O 23 method 1)Calibration with observed T e [O III ] values 2)Calibration with models 3)Calibration with O II recombination lines

30. Peimbert et al. 2006

32. Which Calibration for O 23 ? The best way to calibrate the O 23 method is to use O II recombination lines to obtain the O/H values The O II recombination lines provide abundances that are about 0.2 to 0.3 dex higher than those given by the observed T( 4363/5007 ) values The use of the observed T( 4363/5007 ) values provides a lower limit to the O/H values Since nebular lines are less sensitive to temperature variations than auroral lines, model calibrations (that adjust the nebular lines) are closer to our calibration than those derived using the observed T( 4363/5007 ) values

33. Implications of the O 23 Calibration Our new calibration has implications on the metal production in the Universe and therefore on the star formation rate With this calibration and observations at different z values of strong nebular lines it will be possible to study the chemical evolution of the Universe as a whole

34. Determination of the Primordial Helium Abundance, Y P, with t 2 = 0.000 and t 2 ≠ 0.000 ∆Y (Hc) Y (t 2 = 0.000) Y (t 2 ≠ 0.000) Y P (t 2 ≠ 0.000) NGC 3460.0015 ± 0.00050.25370.2507 ± 0.0027 ± 0.00150.2453 ± 0.0027 ± 0.0019 NGC 23630.0057 ± 0.00160.25510.2518 ± 0.0047 ± 0.00200.2476 ± 0.0047 ± 0.0022 Haro 290.0047 ± 0.00130.25770.2535 ± 0.0045 ± 0.00170.2500 ± 0.0045 ± 0.0019 SBS 0335-0520.0144 ± 0.00380.25940.2533 ± 0.0042 ± 0.00420.2520 ± 0.0042 ± 0.0042 I Zw 180.0114 ± 0.00310.25290.2505 ± 0.0081 ± 0.00330.2498 ± 0.0081 ± 0.0033 Y (sample)0.0056 ± 0.00150.25540.2517 ± 0.0018 ± 0.00210.2477 ± 0.0018 ± 0.0023 Peimbert et al. 2007

35. The Y P Determination Error Budget SourceError Collisional Excitation of the H I Lines±0.0015 Temperature Structure±0.0010 O (∆Y/∆O) Correction±0.0010 Recombination Coefficients of the He I Lines±0.0010 Density Structure±0.0007 Underlying Absorption in the He I Lines±0.0007 Reddening correction±0.0007 Recombination Coefficients of the H I Lines±0.0005 Underlying Absorption in the H I Lines±0.0005 Ionization Structure±0.0005 Collisional Excitation of the He I Lines±0.0005 Optical Depth of the He I Triplet Lines±0.0005 He I and H I Line Intensities±0.0005 Systematic effects Peimbert et al. 2007

36. The Y P Determination Y P, D P, and WMAP Comparison MethodYPYP D P × 10 5  10 bh2bh2 Y P 0.2477 ± 0.0029*2.78 +2.28 − 0.98 5.813 ± 1.8100.02122 ± 0.00663 D P 0.2476 ± 0.00062.82 ± 0.28*5.764 ± 0.3600.02104 ± 0.00132 WMAP 0.2482 ± 0.00042.57 ± 0.156.116 ± 0.2230.02233 ± 0.00082* *Observed values Cosmological predictions based on SBBN and observations Peimbert et al. 2007

37. 1/5 Oxygen Abundance of: 30 Doradus Observational “Direct Method” ( T ( 4363 / 5007 ); t 2 =0.00 ) A. Peimbert 2003 8.33-0.21 Chemically Inhomogeneous Photoionization Model Tsamis & Pequignot 2005 8.45-0.09 Observational t 2 Method ( Oxygen Recombination Lines ) A. Peimbert 2003 8.54---

38. 2/5 Oxygen Abundance of: Orion Nebula Observational “Direct Method” ( T ( 4363 / 5007 ); t 2 =0.00 ) Osterbrock et al. 1992 Pilyugin et al. 2003 Deharveng et al. 2000 8.49 8.50 8.51 -0.16 -0.15 -0.14 Chemically Homogeneous Photoionization Models Baldwin et al. 1991 Rubin et al. 1991 8.58 8.60 -0.07 -0.05 Observational t 2 Method ( Oxygen Recombination Lines ) Esteban et al. 2004 8.65---

39. 3/5 Oxygen Abundance of: Solar Vicinity Photospheric Solar Value Asplund et al. 2005 8.65 Present Day ISM based on the Solar Value and Galactic Chemical Evolution Models Carigi et al. 2005 8.79+0.02 Present Day ISM based on the Solar Value and Young G Dwarf Stars Bensby & Feltzing 2005 8.80+0.03 Present Day ISM based on H II Regions + Dust Content (Mainly Orion) Esteban et al. 2005 8.77---

40. 4/5 Oxygen Abundance of: H II Regions Observational “Direct Method” ( T ( 4363 / 5007 ); t 2 =0.00 ) e.g. Pilyugin & Thuan 2005 -0.30 — -0.20 Photoionization Models (Fitting Nebular Lines) e.g. Kobulnicky & Kewley 2004 McGaugh 1991 -0.05 — +0.10 -0.15 — +0.00 Observational t 2 Method ( Oxygen Recombination Lines ) e.g. Peimbert et al. 2006 --- * log (O/H) + 12 = 8.2 - 8.9

41. 5/5 Primordial Helium Abundance: H II Regions Observational “Direct Method” ( T ( 4363 / 5007 ); t 2 =0.0; Incomplete Error Estimate ) Izotov et al. 2007 0.2516 ± 0.0011 O bservational “Direct Method” ( T ( 5007 / 4363 ); t 2 =0.0; Full Error Estimate ) Peimbert et al. 2007 0.2523 ± 0.0027 Observational t 2 Method ( Balmer continuum and He I lines with MLM ) Peimbert et al. 2007 0.2477 ± 0.0029 Primordial Deuterium + SBBN O´Meara et al. 2006 0.2476 ± 0.0006 Wilkinson Microwave Anisotropy Probe + SBBN Spergel et al. 2006 0.2482 ± 0.0004

42. The End