METALS IN THE ICMMETALS IN THE ICM

LESSON 3LESSON 3

OUTLINE OF THE LESSONOUTLINE OF THE LESSON

• INSTRUMENTATION FOR X-RAY SPECTROSCOPY• MEASUREMENT TRICKS• ENRICHMENT PROCESSES AND RELATIVE

CONTRIBUTION OF SUPERNOVAE (SNIa AND SNII)• WHAT CAN WE INFER AND WHAT CAN WE DO WITH

METALS X-RAY LINES

AN X-RAY SPECTROMETERAN X-RAY SPECTROMETER•ENERGY RESOLUTION ΔE (keV)

•ENERGY BANDPASS δE = Emax – Emin (keV)

•QUANTUM EFFICIENCY Q counts/photon

•EFFECTIVE AREA A (cm2)

•RESOLVING POWER R = λ/Δλ = E/ΔE E (keV) x λ(Ǻ) = 12.4

tEQAFN

If you observe a source flux F (photons/cm2 s keV) for a time t the number of detected photons is

And the number of counts per energy bin is

tEQAFE

ENn

The precision of a narrow line measurement is determined by uncertainty in n

2121 //

tEQAFnn

n

AN X-RAY SPECTROMETERAN X-RAY SPECTROMETERSo the key quantity is the product A x Q x ΔE The overwhelmingly important development goal of improving energy resolution ΔE can be matched in practice by improvement of equal magnitude in effective area or throughput ( A x Q)

Better energy resolution requires also bigger X-ray mirrors.

The energy resolution is taken as the FWHM of the response to a monocromatic photon which is approximately Gaussian.

For example XMM 150 eV at 6 keV (EPIC) and 4650 cm2 at 1 keV

CHANDRA 210 eV at 6 keV (ACIS-S) and 800 cm2 at 1 keV

(We hope !) CON-X 2 eV at 6 keV and 15000 cm2

XEUS 2 eV at 6 keV and 23000 cm2

Remember that if AQ is increased in proportion to the decrease in ΔE, the detected count rate N/t increases for a given flux so to obtain spectra of higher energy resolution requires detectors of ever increasing count rate capacity (already a problem ! Pile up in CCD)

X-RAY SPECTROMETERSX-RAY SPECTROMETERS

•ENERGY DISPERSIVE: photon energies are derived from the signal pulse heights at some detector output

•WAWELENGHT DISPERSIVE: X-ray wavelength is derived from the photon arrival position in the spectrum dispersed by a grating or Bragg crystal

In this class of spectrometers we absorb single photons, we create a population of signal carriers and count the signal carriers to obtain information on position, arrival time and E of the photon.

If the mean energy required to create a signal carrier is W, then the number of signal carriers is

N = E/W

And the uncertainty in enegy is given by variance on N so

σ2 = FN F is the Fano Factor

So for the energy resolution and resolving power we have

ΔE = 2.36 (F x E x W)1/2 and R = A x E1/2

The resolving power increases as the square root of the energy

For Xe proportional counters W = 25 eV and F = 0.2

For Si CCD W = 3.6 eV ad F = 0.1

ENERGY DISPERSIVE SPECTROSCOPYENERGY DISPERSIVE SPECTROSCOPY

ENERGY DISPERSIVE SPECTROSCOPYENERGY DISPERSIVE SPECTROSCOPY

PERSEUS WITH MECS

PERSEUS WITH EPIC

A comparison of spectral resolution between a gas proportional counter (MECS on board SAX) and a Si CCD (EPIC-PN) on board XMM. We will return later on residuals in the fits …

DISPERSIVE SPECTROMETERSDISPERSIVE SPECTROMETERSAll convert into dispersion angle and hence into focal plane position in an X-ray imaging detector

•BRAGG CRYSTAL SECTROMETERS (EINSTEIN, SPECTRUM X-GAMMA): Resolving power up to 2700 but disadvantages of multiplicity of cristals, low throughput, no spatially resolved spectroscopy

n x = 2d x sin

•TRANSMISSION GRATINGS (EINSTEIN, EXOSAT, CHANDRA)

m x = p x sin

where m is the order of diffraction and p the grating period

•REFLECTION GRATINGS (XMM)

m x = p (cos - cos)

The resolving power for gratings is given by , assuming a focal lenght f and a position X relative to the optical axis in the focal plane

X = f tan f sin X = f so

is constant

m

p

RX

X

1

DISPERSIVE SPECTROMETERSDISPERSIVE SPECTROMETERSThe disadvantages are the very low effective area (140 cm2 at the peak for RGS), the limited bandpass (0.35-2.5 keV for RGS) and very limited spatially resolved spectroscopy. But the resolving power is R = 500 corresponding to an energy resolution of 2 eV at 1 keV ( 70 eV at 1 keV for CCD)

M87 RGS Spectra Sakeliou et al. (2002)

ABUNDANCE MEASUREMENTSABUNDANCE MEASUREMENTSThe plasma in the ICM is modeled according to the CORONAL model. The basic assumptions are:

•Elastic Coloumb collision time scale shorter than the age or cooling time, so particles (electrons in particular) have a Maxwell-Boltzmann distribution at temperature T

• Collisional (de)excitation processes slower than radiative decays (low densities) so all ions start from the ground state

•Radiation field on the gas is insignificant (no stimulate emission)

•Gas is optically thin, no radiation transport

•Time scales for ionization and recombination considerably less than any relevant hydrodinamic time scale, so the plasma is in collisional equilibrium

So you use a plasma code like MEKAL or APEC and you have the abundance of your favourite element …

ABUNDANCE MEASUREMENTSABUNDANCE MEASUREMENTS“The analysis is very straightforward compared to the analysis of stellar or H II region spectra, where the effects of radiative transfer, dust and ionization balance are very complex and difficult to remove accurately. In many respects, it is easier to derive the abundances from X- ray spectra of the ICM in clusters than in any other field of astrophysics”

Mushotzky et al. (1996)

But there are some subtle tricks, when you consider:

•Resonance scattering

•Temperature structure

RESONANCE SCATTERINGRESONANCE SCATTERINGAbsorption of a line photon followed by immediate reemission in an other direction.

Although for densities and temperatures typical of clusters the gas is optically thin to Thomson scattering for the continuum, it can be optically thin in the resonance X-ray lines (in particular in the denser core) and in particular in the Fe He emission line at 6.7 keV (the most important emission line !)

No longer optically thin, the surface brightness is distort and you underestimate abundances in the core of the cluster.

You can measure it by making the ratio of a supposed optically thick (Fe He ) and optically thin (Fe He at 7.90 keV) lines and see if there are deviation for the prediction of the plasma code.

RESONANT SCATTERING IN PERSEUS ?RESONANT SCATTERING IN PERSEUS ?

If there is resonant scattering, you underestimate abundance in the prominent line at 6.7 keV you cannot reproduce the line at 7.8 keV and you see residuals in the fit

RESONANT SCATTERING IN PERSEUS ?RESONANT SCATTERING IN PERSEUS ?

Assuming optically thin emission

Using only K

Correcting the apparent abundances

Molendi et al. 1998

But the excess could be due also to overabundance of Ni respect to solar ratios (Dupke & Arnaud 2001). The Ni He is at 7.80 keV and you cannot distinguish with the MECS resolution

If you fit the XMM spectra with a MEKAL model you find again an excess …

RESONANT SCATTERING IN PERSEUS ?RESONANT SCATTERING IN PERSEUS ?

Fe He 7.90 keV

Ni He 7.80 keV

As we can see leaving the abundance of Ni free, the excess is due to this element and not to an anomalosly high Fe He line. There is no resonant scattering

XMM OBSERVATION OF M87XMM OBSERVATION OF M87

• DRAMATIC DECREASE OF ABUDANCES WITHIN 1 ARCMIN OF THE CORE (WITH 1T MODELS …) RESONANT LINE SCATTERING (BOHRINGER ET AL. 2001) CONTINUUM OPACITY (MATHEWS ET AL. 2001)

BOHRINGER ET AL. 2001

XMM OBSERVATION OF M87XMM OBSERVATION OF M87

Fitting multi T spectrum with single temperature models give underestimated abundances (“Fe bias” Buote 2000)

Molendi & Gastaldello 2001

WHAT WE KNOW …WHAT WE KNOW …• Intracluster gas is metal rich abundances 0.3 solar with little

dispersion up to redshift z = 0.4 (Allen & Fabian 98)

• Metals are made in stars in the galaxies and then expelled particularly via supernovae

• Iron mass highly correlated with optical light from ellipticals and S0 (Arnaud et al. 92)

WHAT WE KNOW …WHAT WE KNOW …

• Abundance gradients in cool core clusters, while in not relaxed clusters flat abundance profiles (De Grandi & Molendi 2001)

… … AND WHAT WE DON’T KNOWAND WHAT WE DON’T KNOW

• THE ENRICHMENT PROCESSES How metals leave the galaxies to pollute the ICM

• THE DIFFERENT CONTRIBUTION OF THE TWO DIFFERENT TYPES OF SNe (SNII AND SNIa) AS A FUNCTION OF THE POSITION IN THE CLUSTER

SNIa vs. SNIISNIa vs. SNII

•Progenitors have short lifetimes

106 - 107 yr

•yields in solar masses (Nomoto 1997)

O 1.80 Mg 0.11 Si 0.21 Fe 0.1

•Progenitors evolve over 109 yr

•yields in solar masses O 0.14 Mg 8e-3 Si 0.15 Fe 0.7

ENRICHMENT MECHANISMENRICHMENT MECHANISM

• protogalactic winds (Larson & Dinerstein 1975, Kauffmann & Charlot 1998)

• ram pressure stripping (Gunn & Gott 1972, Toniazzo & Schindler 2001)

• galaxy-galaxy interactions

• later galactic winds from cD driven by SNIa or AGN and confined by ICM pressure and gravitational potential (Renzini 1993, Dupke & White 2000)

GALACTIC WINDSGALACTIC WINDS

Supernova driven superwinds develop when the supernova rate per unit volume is high enough that supernova remnants overlap before they cool. These superbubbles return large amount of kinetic energy (thermalized via shock) along with metal enrich ejecta creating a hot high pressure gas. Breaking out the disk expands at high velocity ( 1000 km/s). But you do not observe the wind but its interaction …

H CONTINUUM

WHITE CROSSES ARE 50 YOUNG SNR DETECTED IN RADIO

Strickland 2002

RAM PRESSURE STRIPPINGRAM PRESSURE STRIPPING

Ram-pressure stripping occurs when a galaxy is moving through the ICM approahing the cluster center. At a certain point the ram pressure ICMvREL

2 is so strong that gas is stripped from the galaxy gravitational potential. It is increased in the case of cluster mergings …

TONIAZZO & SCHINDLER 2001

M86

Gavazzi et al. 2001

ENRICHMENT MECHANISM(S) INFERRED ENRICHMENT MECHANISM(S) INFERRED FROMFROM

• CHEMICAL MIX OF ELEMENTS SNII have high /Fe abundance ratios

SNIa have low /Fe abundance ratios

• SPATIAL DISTRIBUTION

creating abundance gradients ram pressure stripping later SNIa wind from cD

homogenising abundances cluster mergers protogalactic winds

THE SUPERNOVAE DEBATETHE SUPERNOVAE DEBATE

• CONTROVERSIAL RESULT FROM ASCA SNII DOMINATED (Mushotzky 1996)

50/50 SNIa/SNII expecially in the core (Fukazawa et al. 1998, Dupke & White 2000)

uncertain SN model yelds (Gibson et al. 1997)

• SNIa Fe MASS FRACTION

SNIISNIa

SNIa

totalFe

SNIaFe

FeFe

Fe

M

Mf

1,

,

Where is the relative frequency of SNIa (Ishimaru & Arimoto 1997)

SOME RESULTS FROM XMM …SOME RESULTS FROM XMM …

A496 Tamura et al. 2001

UNIFORM ENRICHMENT BY SNII AND DOMINANCE OF SNIa IN THE CENTER DUE TO THE cD

SOME RESULTS FROM XMM …SOME RESULTS FROM XMM …

M87 Gastaldello & Molendi 2002

CONSTRAINTS ON SNIa MODELSCONSTRAINTS ON SNIa MODELS

Dupke & White 2000 Gastaldello & Molendi 2002

WHAT THEORETICIANS SAY …WHAT THEORETICIANS SAY …

The lack of evolution is not unexpected. If we try to model the feedback into the IGM by young stars in term of supernova explosion energy output, UV photon output and mass ejection form supernova explosion and reproducing the UV background as a function of redshift and the star formation with redshift (“Madau plot”) we do not expect evolution up to z 1 (Cen & Ostriker 1999).

WHAT THEORETICIANS SAY …WHAT THEORETICIANS SAY …

Renzini 2000,1997,1993

The MFe/LB is independent of cluster richness, hence of custer richness and optical luminosity

•Clusters did not lose irons (barions) nor acquired fresh barions and the conversion of barions into stars and galaxies procedeed with the same efficiency and IMF in all clusters (also because cluster are nearly solar)

•The absence of variations argues that iron is ejected via galactic winds. Ram pressure effects become much stronger with increasing cluster richness and ICM temperature

WHAT THEORETICIANS SAY …WHAT THEORETICIANS SAY …

You can use chemical evolution models for ellipticals, integrate over a cluster luminosity function (say K-band) at each given cosmic time for clusters of different mass (temperature).

Pipino et al.2002

SOME HINTS ON COLD FRONTS ?SOME HINTS ON COLD FRONTS ?

Dupke & White 2003

If cold fronts are related to the survival of dense cluster cores during the process of merging there should be also a chemical discontinuity. In the seminal paper on A496 (Dupke & White 2003) we do not see it but because is a cold front in a relaxed cluster … (see P. Mazzotta lesson)

SUMMARYSUMMARY

•FROM NEW DATA WE ARE BEGINNING TO ADDRESS IN GREAT DETAIL THE ISSUE OF THE VARIOUS CONTRIBUTION OF SNe TO ICM

•WE STILL HAVE NO CLEAR-CUT EVIDENCE OF THE ENRICHMENT PROCESS OR THE DIFFERENT CONTRIBUTION OF VARIOUS ENRICHMENT PROCESSES

•NO CLEAR EXPLANATION OF ABUNDANCE GRADIENTS (ALTHOUGH WE HAVE SOME IDEA)

•WE NEED A MODEL FOR CLUSTERS AND CLUSTERS FORMATION WHICH DEAL WITH BOTH THE HOT (ICM) AND COLD (GALAXIES) PHASE TOGETHER (P. Rosati). METALS ARE THE MOST DIRECT LINK BETWEEN THE TWO.