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T he metal-line emission of the intergalactic medium in OWLS

Swinburne University of Technology, Melbourne November 26 th 2009. T he metal-line emission of the intergalactic medium in OWLS. Serena Bertone (UC Santa Cruz) Joop Schaye (Leiden Observatory) & the OWLS Team Bertone et al 2009 arXiv:0910.5723 Bertone et al 2010a, 2010b, in prep.

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T he metal-line emission of the intergalactic medium in OWLS

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  1. Swinburne University of Technology, Melbourne November 26th 2009 The metal-line emission of the intergalactic medium in OWLS Serena Bertone (UC Santa Cruz) JoopSchaye (Leiden Observatory) & the OWLS Team Bertone et al 2009 arXiv:0910.5723 Bertone et al 2010a, 2010b, in prep

  2. Outline • Introduction: • the intergalactic medium: warm, warm-hot, hot… what is it? • how do we detect it? • New simulations: OWLS • Results: • X-ray & UV emission at z<1 • rest-frame UV emission at z>2 • what gas does emission trace? • dependence on physics

  3. Intergalactic medium: Why do we care? Dark matter: 25% Metals: 0.03% > 90% of baryonic mass is in diffuse gas (Persic & Salucci 1992, Fukugita et al 1998, Cen & Ostriker 1999) unbiased info on matter power spectrum on largest scales reservoir of fuel for star and galaxy formation IGM metallicityconstrains the cosmic SF history interplay IGM-feedback puts constraints on galaxy formation models Neutrinos: 0.3% Stars: 0.5% Free H & He: 4% Dark energy: 70%

  4. Springel 2003 z=6 z=2 z=0 IGM evolution 100 Mpc cosmictime 100 Mpc/h box

  5. Gas temperature distribution T<105 K 105 K <T<107 K T>107 K z=0 Cen & Ostriker 1999 etc Warm-hot gas: WHIM mildly overdense collisionallyionised shock heated by gravitational shocks Warm gas: diffused photo-ionised traced by Lyα forest Hot gas: dense metal enriched haloes of galaxies and clusters

  6. IGM mass and metals The bulk of the metal mass does not trace the bulk of the IGM mass most metal mass most mass Bertone et al. 2009 z=0.25

  7. Mass and metal fractions evolution IGM mass - Dave’ et al. 2001 Average IGM temperature increases with time Most metals locked in stars at z=0 Average metal temperature increases with time Metal mass - Wiersma et al 2009 halo gas diffuse IGM WHIM ICM

  8. How can we detect the IGM? Absorption ✔ z<1: UV Tripp et al 2007 Lehner et al 2007 Danforth & Shull 2005 z>1.5: optical Kim et al 2001 Simcoe et al 2006 ? z<1 Nicastro et al 2005 Rasmussen et al 2007 Buote et al 2009 ✔ T<106 K Rest-frame UV lines Lyα, OVI, CIV… T>106 K X-rays metal lines OVII, OVIII ✗ z<1: UV Furlanetto et al 2004 Bertone et al 2010a z>1.5: optical Weidinger et al 2004 Bertone et al 2010b ? z<1 Fang et al 2005 Werner et al 2008 Bertone et al 2009 ✔✗ Emission

  9. Emission vs. absorption

  10. OWLS • OverWhelmingly Large Simulations • The OWLS Team: • JoopSchaye(PI, Leiden) • Claudio DallaVecchia(MPE) • Rob Wiersma, Craig Booth • Marcel Haas, • Freeke Van De Voort (Leiden) • Tom Theuns (Durham) • Serena Bertone (UCSC) • Ian Mc Carthy (Cambridge) • Alan Duffy (Perth) • Many thanks to the LOFAR and SARA supercomputing facilities

  11. OWLS • many runs (>50) with varying physical prescriptions/numerics(Schaye et al 2009) • cosmological hydrodynamical simulations: Gadget 3 • run on LOFAR IBM Bluegene/L • WMAP 3 cosmology • largest runs: 2x5123 particles • two main sets: L=25 Mpc/hand L=100 Mpc/hboxes • evolution from z>100 to z=2 or z=0

  12. New physics in OWLS • New star formation (Schaye & DallaVecchia 2008): • Kennicutt-Schmidt SF law implemented without free parameters • New wind model (DallaVecchia & Schaye 2008): • winds local to the SF event • hydrodynamically coupled • Added chemodynamics(Wiersma et al. 2009): • 11 elements followed explicitly (H, He, C, N, O, Ne, Si, Mg, S, Ca, Fe) • Chabrier IMF • SN Ia & AGB feedback • New cooling module (Wiersma, Schaye & Smith 2009): • cooling rates calculated element-by-element • photo-ionisation by evolving UV background included

  13. Physics variations in OWLS • Cosmology: WMAP1 vs WMAP3 vs WMAP5 • Reionisation & Helium reionisation • Gas cooling: primordial abundances vs metal dependent • Star formation: • top heavy IMF in bursts • isothermal & adiabatic EoS • Schmidt law normalisation • Metallicity-dependent SF thresholds • Feedback: • no feedback • feedback intensity: mass loading, initial velocity… • feedback implementation • AGN feedback • Chemodynamics: • ChabriervsSalpeter IMF • SN Ia enrichment • AGB mass transfer Schaye et al 2009

  14. Gas cooling rates collisional ionisation eq. photoionisation eq. density dependent Wiersma, Schaye & Smith 2009 Photo-ionisationby UV BK + collisionalionisation equilibrium Cooling rates calculated element by element for 11 species: takes into account changes in the relative abundances

  15. z=0.25 Gas emissivity UV lines X-ray lines • 11 elements – comparable to cooling rates • Collisionalionisation+ photo-ionisationby UV BKimportant at low density • UV lines stronger than X-ray ones Density  Bertone et al 2009-2010a

  16. Emission at low redshift:Soft X-rays& UV lines

  17. Emission at low redshift: X-rays 100 Mpc/h boxes 20 Mpc/h thick slices 15” angular resolution 12 X-ray + 6 UV emission lines Bertone et al 2009

  18. X-ray lines O VIII strongest line lines from lower ionisation states and whose emissivity peaks at lower temperatures trace moderately dense IGM: C V, C VI, N VII, O VII, O VIII and Ne IX lines from higher ionisation states trace denser, hotter gas: C VI, O VIII, Ne X, Mg XII, Si XIII, S XV and Fe XVII Fe XVII emission has different spatial distribution than other elements: later enrichment by SN Ia Bertone et al 2009

  19. UV lines C III and C IV strongest lines: trace gas in proximity of galaxies O VI and Ne VIII trace more diffuse gas than C IV - different spatial distribution no emission from the hottest gas in groups UV emission is a good tracer of galaxies and of mildly dense IGM, but not of IGM in very dense environments Bertone et al 2010a

  20. Surface brightness PDFs • C IV and O VI lines detectable by FIREBALL (Tuttle et al 2008) • Detection of X-ray lines requires new instruments: WFI on the Interntional X-ray Observatory (IXO) Bertone et al 2009, 2010a

  21. What gas produces the emission? emission-weighted particle distributions Emission traces moderately dense gas, not the bulk of the IGM mass and metals. the peak temperature of the emission increases with atomic number and ionisation state X-ray emission traces gas with T>106K O VI and Ne VIII trace diffuse gas C III, C IV and Si IV trace the CGM Bertone et al 2009, 2010a

  22. Impact of physics What happens when changing the physical model?

  23. Impact of physics: X-rays • no feedback: no metal transport localised emission • primordial cooling rates: longer cooling times stronger emission at high density (≈100 times) • momentum driven winds: metals more spread in IGM  weaker emission (≈100 times) • AGN feedback: weaker emission in dense regions Bertone et al 2009

  24. Impact of physics: UV • no feedback: emission localised in galaxies • primordial cooling rates: stronger emission at high density • AGN feedback: weaker emission • changes in wind parameters: small effect Bertone et al 2010a

  25. Emission at high redshift:rest-frame UV lines emission at 2<z<5 25 Mpc/h simulations 2” angular resolution Bertone et al 2010b

  26. IGM emission at z>1.5 • At z>1.5 rest-frame UV lines are redshifted in to the optical band • A number of upcoming optical instruments might detect IGM emission lines at 1.5<z<5: • Cosmic Web Imager on Palomar (CWI, Rahman et al 2006) this year! • Keck Cosmic Web Imager (KCWI) • Antarctic Cosmic Web Imager (ACWI, Moore et al 2008) • IFUs with large fields of view & high spatial resolution • Great chance to observe the 3-D structure of the IGM for the first time!

  27. Higher ionisation states

  28. Lower ionisation states

  29. Emission PDFs Lower ionisation states: single lines shorter wavelengths C III up to 10x stronger than C IV Higher ionisation states: doublets – easy to identify weaker than lower ion. states

  30. CWI: flux limit: 100 photon/s/cm2/sr angular resolution: 2” central regions of groups and galactic haloes Lines detectable by CWI Bertone et al 2010b

  31. Summary • Dense cool gas in the haloes of galaxies is traced by: • UV lines: C III, C IV, Si IV • Low density WHIM gas in filaments is traced by: • UV lines: O VI, Ne VIII • X-ray lines from hydrogen-like atoms: C V, O VII, Ne IX • X-ray lines from elements with low atomic numbers: C VI, O VIII • Dense hot gas in clusters and groups is traced by: • X-ray lines from fully ionised atoms: C VI, O VIII • X-ray lines from elements with high atomic numbers: Mg XII etc. • Detection of WHIM emission by future telescopes: • challenging in low density regions • very likely in groups and cluster outskirts • CWI very likely to detect metal line emission at high z for the first time

  32. Visually… Temperature cuts: the strongest emission comes from the temperature range where the line emissivity peaks Density cuts: the strongest X-ray (O VIII) and UV (O VI) emission comes from the densest gas

  33. Summary:Dependencies on gas properties Median density, temperature and metallicity of gas particles vs. particle emission • Correlation of emission with density and metallicity: highest emission from densest and most metal enriched particles • Median temperature of highest emission corresponds to peak temperature of emissivity curve – as seen in T-nHI diagrams Bertone et al 2009, 2010a

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