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Photoionized plasma analysis. Jelle Kaastra. Introduction. What is a photoionised plasma?. Plasma where apart from interaction with particles also interaction with photons occurs Photon spectrum needs to affect the particles (e.g. heating)

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Photoionized plasma analysis

Photoionized plasma analysis

Jelle Kaastra



What is a photoionised plasma
What is a photoionised plasma?

  • Plasma where apart from interaction with particles also interaction with photons occurs

  • Photon spectrum needs to affect the particles (e.g. heating)

  • Thus, plasma with resonant scattering has photons involved but is not photoionised (although resonance scattering also occurs in photoionised plasmas)


It is all about the optical depth
It is all about the optical depth

  • Optical depth τ = 0: collisional

  • Optical depth τ≠ 0 but not τ >> 1: classical photoionised plasma

  • Optical depth τ >>1: more atmosphere-like or stellar interior-like, not discussed here

  • Note: optical depth depends on photon energy – the above is rather crude


Examples of photoionised plasmas
Examples of photoionised plasmas

  • Accreting sources:

    • Galactic X-ray binaries

    • Active galactic nuclei

  • Tenuous gas (like some components of the ISM/IGM)

  • Nova shells


Feeding the monster
Feeding the monster

  • Gas transported from 1020 to 1012 m scale

  • Disk forms due to viscosity / B-fields / loss angular momentum

  • Only few Msun/year reach black hole


Outflows from the monster
Outflowsfrom the monster

  • Notall gas reaches black hole

  • Outflowsthroughmagnetised jets, disk winds, outflowsfrom torus surrounding disk

  • Gives feedback tosurroundings, but howmuch?


B asics
Basics


Something to think about
Something to think about

  • Most important line features:

    • O-lines (1s-np of O I – O VIII)

    • Fe UTA & other n = 1-2 transitions

    • Fe-K

    • Si lines (see e.g. NGC 3783)

  • Multiple absorption components

  • Blending with foreground galactic features (example: Mrk 509 O IV with Galactic O I)

  • Contamination by emission lines


Photoionisation equilibrium
Photoionisation equilibrium


Key parameter ionisation parameter
Key parameter: ionisation parameter

  • Spectrum depends on ratio photons / particles

  • Common used (Xstar, SPEX): ξ = L / nr2 with:

    • L = ionising luminosity between 1 – 1000 Ryd (13.6 eV – 13.6 keV; note the upper boundary!)

    • n is hydrogen density (NB, different from ne!)

    • r is distance from ionising source

  • Alternative (Cloudy): UH = QH / 4πcnr2 with:

    • QH number of H-ionising photons (13.6 Ryd – ∞)


Photoionised plasmas
Photoionisedplasmas

  • Irradiated plasma

  • Twobalanceequations:

Photons:

Photo-ionisation

Heatingbyphoto-electrons

Electrons:

Radiative recombination (electron capture)

Cooling by collisional excitation (followed by line radiation)


Photoionisation modelling
Photoionisationmodelling

  • Radiation impacts a volume (layer) of gas

  • Different interactionsof photonswithatomscauseionisation, recombination, heating & cooling

  • In equilibrium,ionisation state of the plasma determinedby:

    • spectral energy distributionincomingradiation

    • chemicalabundances

    • ionisation parameterξ=L/nr2withLionisingluminosity, ndensityandrdistancefromionising source; ξessentially ratio photondensity / gas density


First balance equation ionisation stages 1
First balance equation: ionisation stages (1)

  • Same rates as for CIE plasmas:

  • Collisional ionisation

  • Excitation auto-ionisation

  • Radiative recombination

  • Dielectronic recombination

  • At low T, charge transfer ionisation & recombination


First balance equation ionisation stages 2
First balance equation: ionisation stages (2)

  • New for PIE plasmas:

  • Photoionisation

  • Compton ionisation (Compton scattering of photons on bound electrons; for sufficient large energy transfer this leads to ionisation)


Second balance equation energy
Second balance equation: energy

  • Balance: heating = cooling

  • Take care how heating etc is defined: we use here heating/cooling of the free electrons

  • For instance, for e-+ione-+ion++e-we assign the ionisation energy I to the cooling of the free electrons


Heating processes
Heating processes

  • Compton scattering (photon looses energy)

  • Free-free absorption

  • Photo-electrons

  • Compton ionisation

  • Auger electrons

  • Collisional de-excitation


Cooling processes
Cooling processes

  • Inverse Compton scattering (photon gains energy)

  • Electron ionisation

  • Recombination

  • Free-free emission (Bremsstrahlung)

  • Collisional excitation


Heating cooling ngc 5548 in 2013
Heating & cooling (NGC 5548 in 2013)

Inverse Compton

Recombination

Free-free emission

Collisional excitation

Electron ionisation

-------------------

Compton scatter

Photoelectrons

Auger electrons

Compton ionisation

(Coll. de-excitation)

(Free-free absorption)


Heating cooling ngc 5548 obscured
Heating & cooling (NGC 5548 obscured)

Inverse Compton

Recombination

Free-free emission

Collisional excitation

Electron ionisation

-------------------

Compton scatter

Photoelectrons

Auger electrons

Compton ionisation

(Coll. de-excitation)

(Free-free absorption)


Performance 151 grid points
Performance (151 grid points)

  • Same run on NGC 5548 obscured SED:

  • XSTAR: 40 hours (& crashed for kT > 10 keV)

  • Cloudy: 4 hours

  • SPEX: 5 minutes

  • Okay the above may depend on optimalisation flags etcetc, but ….


Performance
Performance

  • Often people make a grid of models as function of few parameters  table grid  feed into favorite fitting program

  • SPEX pion model allows fast instantaneous calculation & simultaneous fitting of the continuum of any shape; multiple stacked layers


Stability photo ionisation equilibrium examples from detmers et al 2011
Stabilityphoto-ionisationequilibrium(examplesfromDetmers et al. 2011)

Ξ = Fion/nkTc= ξ/4πckT

Stable equilibrium fordT/d Ξ> 0


Stability curves differ here case ngc 5548 mehdipour et al 2014
Stability curves differhere case NGC 5548 (Mehdipour et al. 2014)


Other useful representations same data as previous slide
Other useful representations (same data as previous slide)


Differences photo ionisation models mrk 509 sed
Differencesphoto-ionisationmodels(Mrk 509 SED)


Differences photoionisation models ngc 5548 obscured case
Differences photoionisation models(NGC 5548 obscured case)


Photoionisation modelling1
Photoionisationmodelling


Practical examples from spex 1
Practical examples from SPEX (1)

  • Most simple model: slab

  • Input:

    • Set of ionic column densities (arbitrary, no physics involved)

    • Outflow velocity

    • Line broadening

    • Covering fraction fc

  • Transmission: T(E) = (1 – fc) + fc e-τ(E) with τ(E) containing all physics of absorption

  • Emission needs to be modelled separately


Practical examples from spex 2
Practical examples from SPEX (2)

  • next simple model: xabs

  • Input:

    • Set of ionic column densities pre-calculated using real photoionisation code

    • Ionisation parameter ξ = L/nr^2

    • Outflow velocity

    • Line broadening

    • Covering fraction fc

  • Transmission: T(E) = (1 – fc) + fc e-τ(E) with τ(E) containing all physics of absorption

  • Emission needs to be modelled separately


Practical examples from spex 3
Practical examples from SPEX (3)

  • next simple model: warm

  • Input:

    • Set of ionic column densities pre-calculated using real photoionisation code

    • Absorption measure distribution dNH(ξ)/dξ, parametrized by powerlaw segments

    • Outflow velocity

    • Line broadening

    • Covering fraction fc

  • Transmission: T(E) = (1 – fc) + fc e-τ(E) with τ(E) containing all physics of absorption

  • Emission needs to be modelled separately


Practical examples from spex 4
Practical examples from SPEX (4)

  • latest model: pion

  • Input:

    • Arbitrary SED (using SPEX emission components, or file, or …)

    • Does self-consistent photoionisation calculations

    • Ionisation parameter ξ = L/nr^2

    • Outflow velocity

    • Line broadening

    • Covering fraction fc

  • Transmission: T(E) = (1 – fc) + fc e-τ(E) with τ(E) containing all physics of absorption

  • Emission (still) needs to be modelled separately


Future extensions of the pion model
Future extensions of the pion model

  • Include also emission (using SPEX plasma code core; several processes need updates)

  • Cooling at low T not yet accurate enough (Rolf Mewe’s CIE model stopped at K-like ions or higher)

  • Thicker layers (simple radiation transport using escape factors)

  • NB only the Titan code takes full radiative transfer into account



Absorption measure distribution
AbsorptionMeasure Distribution

Discrete components

Emission measure Column density

Continuous

distribution

Ionisation parameter ξ

Temperature


Decomposition into separate
Decomposition into separate ξ

  • Early example: NGC 5548 (Steenbrugge et al. 2003)

  • Use column densities Fe ions from RGS data

  • Measured Nion as sum of separate ξ components

  • Need at least 5 components


Separate components in pressure equilibrium or continuous
Separate components in pressure equilibrium, or continuous?

Discrete components in pressure equilibrium?

Continuous NH(ξ) distribution?

Krongold et al. 2003

Steenbrugge et al. 2005


Discrete ionisation components in mrk 509 detmers et al 2011 paper iii
Discrete ionisationcomponents in Mrk 509?Detmers et al. 2011 paper III

  • Fitting RGS spectrum with 5 discrete absorber components (A-E)

  • Gives excellent fit


Continuous amd model mrk 509 detmers et al 2011
Continuous AMD model?Mrk 509, Detmerset al. 2011

  • Fit columns withcontinuous (spline) model

  • C & D discrete components!

  • FWHM <35% & <80%

  • B (& A) toopoorstatisticsto prove ifcontinuous

  • E harder determined: correlationξ & NH

  • Discrete components

D

E

C

B


Pressure equilibrium no
Pressure equilibrium? No!


A comparison between sources
A comparison between sources

  • All Seyfert 1s show similar trend

  • NH increases with ξlike power law

  • High ξ cut-off?

  • Same behaviour in Seyfert 2s (NGC 1068, Brinkman et al. 2002)


Time dependent p hotoionisation
Time-dependent photoionisation


Why study time dependent photoionisation
Why study time-dependent photoionisation?

  • Because most photoionised sources are time-variable

  • Gives opportunity to determine distance of gas from ionising source  mass loss, kinetic luminosity etc


The recombination time scale
“The” recombination time scale

  • Pure recombination equilibrium:

    0 = dni/dt = niRi-1 + ni+1Ri

  • This leads, with Ri = neαi to characteristic time

    trec = 1 / [ne (ni+1/ni – αi-1/αi)]

  • Thus, we see that trec~1/ne

  • However, there is always a point where ni(ξ) and ni+1(ξ) are such that trec∞, and this point is usually close to where ni(ξ) peaks!


Density estimates line ratios
Density estimates: line ratios

  • ξ = L/nr2

  • C III has absorption lines near 1175 Å from metastable level

  • Combined with absorption line from ground (977 Å) this yields n

  •  n = 3x104 cm-3 in NGC 3783 (Gabel et al. 2004)  r~1 pc

  • Onlyappliesforsome sources, low ξ gas

  • X-rayssimilarlines, sensitivetohighern (e.g. O V, Kaastra et al. 2004); no convincing case yet (in AGN, but Fe linesfromexcited levels are seen in X-raybinaries


Density estimates reverberation
Density estimates: reverberation

  • If L increases for gas at fixed n and r, then ξ=L/nr² increases

  •  change in ionisation balance

  •  column density changes

  •  transmission changes

  • Gas has finite ionisation/recombination time tr (density dependent as ~1/n)

  •  measuring delayed response yields trnr


Lightcurve mrk 509 during 100 days kaastra et al 2011 paper i
LightcurveMrk 509 during100 days(Kaastra et al. 2011, paper I)

  • Factor ~2 increase in soft X-ray

  • Correlated with UV

  • No correlation with hard X-ray

UV

Soft X-ray

Hard X-ray


Spectral energy distribution mehdipour et al 2011 paper iv
Spectral energy distribution(Mehdipour et al. 2011, paper IV)

DBB

Soft excess

Power law


Time dependent seds mrk 509 kaastra et al 2012 paper viii
Time-dependentSEDsMrk 509(Kaastra et al. 2012, paper VIII)


Predicted signal
Predictedsignal

Simplified case: predicted change transmission forinstantaneous 0.1 dexincrease L, at spectralresolution EPIC/pn

Signal is weak (1% level) but detectable


Time dependent calculation
Time-dependentcalculation

Total

Soft X

Hard X

Time evolution ion concentrations ni:

dni/dt = Aij(t) nj

Aij(t) contains t-dependentionisation & recombinationrates


Limits distance
Limitsdistance

  • Recombination time scaledensity n

    • Using ξ=L/nr2r=√(L/ ξn)

    • No variabilityseen: lower limit r

    • Variabilityseen, but sparse data: upper limit r

  • Using measured column densityN=nΔr withΔrthicknesslayer & Δr <r r<L/Nξ

  • [O III] 5007 has been imaged (Phillips et al. 1986) (r=3 kpc)


Summary distance limit methods for 5 components in mrk 509
Summary distance limit methodsfor 5 components in Mrk 509


Results where is the outflow kaastra et al 2012 paper viii
Results: where is the outflow?(Kaastra et al. 2012, paper VIII)


Abundances outflow steenbrugge et al 2011 paper vii
Abundancesoutflow(Steenbrugge et al. 2011, paper VII)

  • Relative metal abundances close to Solar

  • Absolute abundancesawait new COS data withhydrogenLymanseries

  • Onlydoableaftercarefullphotoionisationmodelling



Surprise v ery low soft x ray flux
Surprise: very low soft X-ray flux





Obscuring stream
Obscuring stream

  • Two components:

  • Main: log ξ = -1.2, NH=1026 m-2, fcov=0.86 (X-ray) and ~0.3 in UV; produces UV BAL

  • Second: almost neutral, NH=1027m-2, fcov=0.3 (X-ray) and <0.1 in UV

  • Partial covering inner BLR, v up to 5000 km/s, inside WA  distance few light days (~1014 m, 0.003 pc)

  • Obscuration already 3 years ongoing




Importance for feedback murray et al 1995
Importance for feedback(Murray et al. 1995)


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