Abundances in the BLR Nathan Stock February 19, 2007

Motivation Metallicity affects properties of the AGN (Ferland, 1996) –Opacity –Kinematics –Structure Outflows enrich IGM (Friaca, 1998) –Dust increases obscuration of high-z objects

Motivation Provides information on chemical history of the gas (Ferland 1996) –Representative of much larger region (~100’s pc) than BLR itself (>1 pc) –Constrain early nucleosynthesis / SFR in the host galaxy

Typical QSO Spectra Hamann 1999

Collisionally excited lines Line strengths of collisionally excited lines High density limit Low density limit Hamann 1999

Recombination lines Especially relevant for H, He Line strength (oversimplification) Hamann 1999

Calculating abundances Take the ratio of two lines Adjust from abundance of element in ionized state to total abundance Express in standard form Hamann 1999

Typical QSO Spectra Hamann 1999

But wait… CIV/Lyα ratio does not depend on metallicity Hamann 1999 CIV acts as a cooling line: as C/H ↑, temperature ↓

And there’s more… We want the lines to be emitted from the same region of the BLR –i.e. elements have similar ionization energies Differing n e, β in different regions Different continua as well Hamann 1999

And more… We want similar values of n crit –Removes possible dependence on n e –Both emitters either in high density or low density limit Note: similar n crit does not help if the emitters are in different regions of the BLR! Hamann 1999

Problem 1: Appropriate element To find metallicity, we cannot use C/H when Z>.02 Z sun Moreover, C and other elements don’t have ions in the same range as HII Ratios of elements to each other would also seem counterproductive, but… Hamann 1999

Secondary production of N C, N, O produced in later stages of stars N is also produced in CNO cycle –Valid for Z>.2 Z sun Result: N/O α O/H α Z Hamann 1999

Caveat: secondary N is delayed Delayed production will be important if Z enrichment is faster than stellar lifetimes –This is true in dense environments So, –q is a delay factor q=0 in slow evolution limit q~.5 in fastest evolutions Hamann 2001

Problem 2: Appropriate ionization We want lines that are emitted from the same region of the BLR (co-spatial) Need to model ionization regions –Define U, n H, abundance, incident spectrum –Use CLOUDY to identify ionized regions Hamann 2001

Results for a ‘typical’ BLR n H = 10^10 U =.1 Solar abundances

Problem 3: Appropriate n crit We want critical densities to be similar –Ensures ions are in same density limit in each part of the BLR Hamann 2001

Putting it together We can model how emission line intensities (equivalent widths) vary… – …with the flux of incident ionizing photons (Φ H ) –… with hydrogen density (n H ) Integrating the intensities over (Φ H, n H ) space (with appropriate weighting) gives us the total line strength Hamann 2001

Example: CIII] λ977 Φ H too low, C is neutral Φ H too high, C is ionized –In both limits, the CIII] equivalent width is weak n H too high, exceed n crit –Emission line is collisionally suppressed n H too low, forbidden lines become efficient coolants –Gas temperature drops, weakening emission lines Hamann 2001

Equivalent Widths of lines Hamann 2001

Flux Ratios of lines

Total line strength Integrate over the space to find the line strength –7  14 for log n H –17  24 for log Φ H Assume equal weighting –Previously shown to reproduce AGN broad emission lines fairly accurately (Baldwin 1997) Similarly, we can find the total line ratio Hamann 2001

Finding line ratio – Z relations Previously considered solar abundance Now, vary the abundances –Recalculate emission line ratios Hamann 2001

Dependence affected by shape of incident spectrum Solid curve: MF87 spectrum Dotted curve: α=-1 power law Dashed curve: segmented power law

Things to note: Hamann 2001 NIII]/OIII] and NV/HeII: N found in narrower region  line ratios will underestimate abundances

Possible NV contamination? NV line at λ1240, Lyα at λ1216 –If NV is moving away from Lyα at the right velocity, it can absorb and rescatter Lyα emission –v ~ 5900 km/s  achievable in BAL winds Hamann 1999 Lyα NV

BAL contributions likely small Only ~30% of Lyα photons interact with NV in BAL –Most Lyα passes through Given 12% BAL covering factor & typical BAL velocity profile NV BAL < 25% NV BLR Moreover, BAL peak would be much wider – >10,000 km/s BAL vs km/s BLR half-widths –Do not see this in spectra Hamann 1996

Integ Hamann 2001 Line ratios taken from the literature on quasars imply they have BLR regions which have greater than solar metallicities.

Abundances in Quasars All the most robust line ratios show Z~2-3Z solar is typical all quasars Abundances constitute a lower limit on actual Z because we assumed q=0 –In quasars, q is almost certainly not 0! Hamann 2001

More quasar data 70 quasars, each Z found by averaging several N ratios Average metallicty of the sample: Z~4 Z solar Dietrich, 2003

What does this high Z tell us? High metallicity  significant chemical evolution occurred before our observations –Even at the highest redshift quasars! Chemical evolution models imply most of the original gas has been processed by stars Vigorous star formation and evolution likely precedes the epoch of quasar activity –Have ~1-2 Gyr of stellar formation time at z=5 Hamann 2001

Metallicity-Luminosity Relation? Metallicity-tracking line ratios in the BLR do not appear to correlate with redshift However, it DOES appear to correlate with quasar luminosity Nagao 2006

Metallicity-Luminosity Relation? Is this relation fundamental or apparent? –Evidence that it’s M BH that really matters Warner 2007

Metallicity-M BH Relation What brings about this relation? –M BH α M host (at least in local universe) –We expect a more massive host galaxy to have a higher metallicity More stars producing elements Evidence that M BH α M host relation applies at high redshift as well A possible way to find host masses even at high redshift? Hamann 2007

So what’s next? Direct measurements of host galaxies –Mass, age, metallicity More samples, over a variety of properties –Higher redshifts, lower luminosities The usual pushing the limits of what we can observe Other abundances –Not just C, N, O –Fe, Si? Hamann 2007

Bibliography Dietrich, M., Hamann, F., et al. 2003, 589, 722 Ferland, G.J., Baldwin, J.A., Korista, K.T., et al. 1996, ApJ, 461, 683 Hamann, F., Dietrich, F. & Ferland, G.J. astro-ph Hamann, F., & Ferland, G.J. 1999, ARA&A, 37, 487 Hamann, F., & Korista, K.T. 1996, ApJ, 464, 158 Hamann, F., Korista, K.T., Ferland, G.J., et al. Astro- ph Nagao, T., Marconi, A., & Maiolino, R. 2006, A&A, 447, 863 Sadat, R., Guiderdoni, B. Silk, J. 2001,°a, 369, 26 Warner, C., Hamann, F., & Dietrich, M. 2007, ApJ, submitted