Transcript What can emission lines tell us? Grażyna Stasińska 2006

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What can emission lines
tell us?
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lecture 4
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Grażyna Stasińska
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Some pending questions
and some strategies to solve them
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aperture correction
dereddening
underlying stellar absorption
escape of ionizing radiation
dust
temperature fluctuations
chemical inhomogeneities
the role of shocks
Aperture correction
When the studied objects are more extended than the observing beam
• aperture correction is needed if the observing beams are not the same for all
wavelenghts. For example
• combining ground optical and UV spectra from IUE or HST
• combining FIR measurements with optical measurements
Aperture correction are usually done
• using line ratios that have a known intrinsic value (e.g. HeII 1640 / He II 4686)
• using ratios of apertures
• Such procedures bear uncertainties
• they do not take into account the ionization stratification of the nebulae
The best way to do this (collaboration Stasinska, Morisset, Simon-Diaz ...2006)
• build a photoionization model reproducing the observed H surface brightness distr.
• compute the intensities through each observing slit
• compare the observed intensities with the model intensity through appropriate slit
Correction for dust extinction
The method:
• The “logarithmic extinction at H”, C, is derived from the observed Ha/H ratio by
comparing it to the theoretical one for case B recombination
assuming an extinction law f(l)
C = [ log (FHa / FH)B - log (FHa / FH)obs ] / (fa- f)
• Emission line ratios are then dereddened using the formula
log (Fl1 / Fl2)corrB = log (Fl1 / Fl2)obs + C (fl1- fl2)
Problems:
• The “extinction law” is not universal
• The intrinsic Ha/H ratio may be different from the theoretical case B
(collisional excitation, case C)
• If some dust is mixed with the ionized gas and strongly contributes to the extinction,
no “extinction law” applies
the “extinction law” is not universal
• the canonical extinction law corresponds to
RV=AV / E(B-V) = 3.2
• in Orion RV = 5.5
• towards the Galactic bulge, RV ~ 2.5 (eg Stasinska
et al 1994)
• larger values of RV are found for lines of sight
crossing molecular clouds where dust grains are
expected to be larger
Extinction laws corresponding to
various values of RV=AV / E(B-V)
.1989
Histogram of RV for 95 galactic O stars
Patriarchi et al 2001
checks on the reddening correction
Before dereddening check that conditions for case B are likely satisfied
• If not, consider building a photoionization model with a code that treats the H atom
correctly, and redden the resulting the emission lines to fit the observed Balmer
decrement.
If case B is relevant for the object under study
• check that, after reddening correction , Hg / H is close to the case B value
• If not, [OIII]4363/5007 is likely to be in error by the same amount as Hg/H differs
from the case B value
• If many Balmer lines are measured with good accuracy
• rather than using an “extinction law”
• fit the observed Ha/H Hg/H  to the theoretical case B values
• this method is valid (in first approximation) also if dust is mixed with the HII gas
• but, of course, it does not allow to derive C(H)
underlying stellar absorption
Some nebular spectra may contain a lot of light from stars
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•
•
•
giant HII regions
small size PNe
nuclei of galaxies
entire galaxies
This light may contaminate the emission lines
This is an important problem
• if the emission lines are faint
• eg faint nuclei of galaxies
• if one desires a high accuracy in line measurements
• eg determination of the pregalactic helium abundance
underlying stellar absorption
Possible ways out
• correct for reddening and stellar absorption at the same time
• in order to obtain the correct Balmer decrement (eg Izotov et al)
• observe with good spectral resolution
• stellar lines are generally broader than emission lines
• use « template » continuum spectra
• this was common practise in the last decade to study faint nuclear
emission regions
• do a model for the stellar light and subtract it from the observed spectrum
• this is now routinely done by many groups studying galaxy spectra and
using stellar population synthesis techniques
• Cid Fernandes et al, Tremonti et al etc...
escape of ionizing radiation
most nebular studies assume that the nebulae are ionization-bounded
• modelling of planetary nebulae
• estimate of T* via the Zanstra method
• estimation of star formation rates in galaxies etc ...
there is growing evidence that many nebulae are density bounded
in at least some directions
• planetary nebulae
• giant HII regions
• high redshift galaxies
the effects of dust
• evidence for the presence of dust mixed with ionized gas
• the effect of dust on the ionization structure
• the effect of dust on the nebular thermal balance
• the effect of dust on resonance lines
• relevance of elemental abundances in case of depletion
dust coexists with ionized gas
Evidence from IR spectroscopy
• strong IR continuum due to dust
heated by stellar radiation
• emission features attributed to
silicate or carbon-based particles
• the dust temperature indicates that
grains are not limited to the neutral
outskirts of the nebulae
IR flux distribution in the PN NGC 3918
Harrington et al 1988
Evidence from IR imaging
• dust emission is seen in the ionized region (eg Graham et al 1993)
Evidence from optical spectroscopy
• Refractory elements (Fe, Mg, Si,Ca) are largely depleted in HII regions and PN,
indicating that dust is intimately mixed with ionized gas
the dust-to-gas mass ratio
in ionized nebulae
Its estimate depends on the adopted grain size distribution
Very different values are quoted
• in HII regions
• md/mg = 10-4 - 10-3 Hoare et al 1991
• in planetary nebulae
• md/mg = 10-4 - 10-2 Natta & Panagia 1981, Stasinska & Szczerba 1999
• an extreme case: the dusty PN in the globular cluster M22
• md/mg = 0.4 Borkowski & Harrington 1991
dust and the ionization structure
the optical depth of dust in the EUV can be significant
D = sD (nD/nH) nH R ≈ 0.6 (U/10-3) if (nD/nH) has the local ISM value
absorption of ionizing photons by dust reduces
• the intrinsic H luminosity and the ionization parameter
effect of the wavelength dependence of s
l
Aanestad 1989
• s l peaks at 700A
• dust will absorb
Hionizing photons more
efficiently than He-ionizing
photons
• the ionization level
increases with respect to
dust-free case
• in Orion the He+ zone
merge with the H+ zone
Baldwin et al 1991
effects of the presence of grains
on the thermal balance
consequences of depletion
• coolants such as Mg, Si, Fe (also C to a lesser extent) are partly tied up in grains
• collisional line cooling is therefore reduced
• especially in outer zones where Mg, Si and Fe cooling is most efficient
• and Te is enhanced relative to the dust-free case
gas-grain collisions
• are a cooling factor for the gas
photolectric effect on dust grains
• electrons ejected from grains by photoelectric effect heat the gas Spitzer 1948
r
dust-
and H- heating
thermal gains of the gas due to photolectric effect on dust :
GD = nD
 4p Jn / (hn) an (D)(hn-E°) dn
thermal gains of the gas due to H ionization:
GH = A n(H+) ne 
ratio of dust- to H-heating :
GD / GH = nD
 4p Jn / (hn) an (D)(hn-E°) dn  n(H+) ne  )
GD / GH a  n D  n H  U
heating by photoelectric effect on dust grains becomes relatively important
• when dust-to-gas ratio is high
• when ionization parameter is high
the effects of grains on Te
dust can be important for the thermal balance of the gas
Baldwin, Ferland, Martin et al 1991
Fraction of total heating due
to photoelectric effect
and fraction of total cooling
due to grain-gas collisions
in the Orion nebula
Baldwin et al 1991
the effects of grains on Te
Heating by dust is more efficient when a population of small grains (10A) is present
Dopita & Sutherland 2000
heating contributions from
photoionization of
small grains __
large grains ---__
hydrogen
_________________________________________
Te
as a function of fractional radius
the effects of grains on Te
small grains can give rise to important “temperature fluctuations” in
filamentary or knotty nebulae
Stasinska & Szczerba 2001
___ T
e
___ T
e
_____
_____
filamentary dust free model
ne
ne
filamentary model with small dust grains
if dielectronic recombination is enhanced at high Te, small grains could also
perhaps help solving the recombination line conundrum
the effects of dust on resonance lines
Attenuation of resonance lines
• resonance lines experience important scattering in the nebulae
• can be selectively attenuated by dust absorption compared to other lines
• should not be used for abundance determinations without caution
Attenuation of resonance
lines by dust in NGC 3918
Harrington et al 1988
Departure from case B
• destruction of H Lyman lines by dust absorption
100% conversion of
high-n Lyman lines into H Lya and Balmer lines (the case B assumption is no more
verified) Cota & Ferland 1988
relevance of elemental abundances
in case of depletion
Mg, Si, Fe, Ni, Ca
• these elements can be almost entirely in the form of grains
• their abundances in the gas phase cannot be easily used as indicators
• of chemical evolution of galaxies
• or nuclear processes in PN progenitors
C
• can be heavily depleted by carbon-based grains (graphite, PAHs...)
O
• can be slightly depleted (20% for Orion, estimated from depletion pattern of metals,
Esteban et al 1998)
He, Ne, Ar
• rare gases, do not combine into grains
Temperature fluctuations
Were postulated by Peimbert (1967) to explain discrepancies
betwen Te from various diagnostics
Peimbert’s formalism
T0 (N i ) 
 T N N dV
 N N dV
e
i
i
e
e
Do temperature fluctuations exist?
 (Te  T0(Ni )) Ni Ne dV
2
t 2 (Ni ) 
T02  Ne Ni dV
• see reviews by Peimbert 1995, 2001, Mathis 1997, Stasinska 1998, Esteban 1998,
2001
• Numerous studies point towards t2 ~ 0.04
• But little direct evidence is seen
example of indirect evidence for t2 ≠ 0
In planetary nebulae Te from Balmer discontinuity is smaller than
T[OIII] 4363/5007 (Liu & Danziger 1993)
t2 ~ 0.04 is a
representative value
If Te fluctuations exist they affect abundance
determinations
e.g. abundance derived in M8 (Peimbert et al 1993) using Taylor series expansion of the
line emissivites for various values of t2
t2=0
t2=0.02
t2=0.04
t2=0.06
He
11.02
11.01
11.00
10.99
C
8.21
8.31
8.48
8.77
N
7.57
7.66
7.77
7.88
O
8.50
8.60
8.71
8.84
Abundances derived from optical forbidden lines with respect to H are
underestimated when ignoring t2
Abundance ratios like N/O or C/O are less affected
Abundances derived from recombination or FIR lines are not affected
visualisation of the Peimbert formalism
on a two-zone toy model
V1
n1
N1
T1
V2
n2
N2
T2
volume
electron density
ionic density
temperature
f = N2n2V2 / N1n1V1
in this case the values of T0 and t2 are simply
T1  fT2
T0 
1 f
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__t0
__0.02
__0.04
__0.06
(T1  T0 )  f(T2  T0 )
2
t 
(1 f )To 2
2
2
Even in such a simple model , the
temperature distribution requires
3 parameters to be defined (T0, t2, f),
not 2 (T0, t2 )
2
variations of T1 and T2 with f for fixed t2 and fixed T0
• f >> 1 may represent a photoionized
nebula with small shock-heated
regions of very high T1
• f << 1 may represent a nebula with
high metallicity clumps of low T2
effect of t2 on derived abundances
if t2 is not accounted for
• O++5007 is underestimated
because T4363/5007 overestimates
the temperature characteristic of
the [OIII]5007 emission
• the bias depends on T0 and f
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__t 0
__ 0.02
__ 0.04
__ 0.06
2
t
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__t 0
__ 0.02
does the Peimbert formalism give correct abundances
?
__ 0.04
__ 0.06
2
• t2obs ≠ t2
• the result depends on f
Even if the computed t2 is not equal
to the true one, does the Peimbert
formalism lead to accurate
abundances?
• not quite
• the bias depends on T0 and f
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frequent misuse of the Peimbert’s formalism
• from expansion of the emission coefficient in Taylor series, and
integrating over the observed volumes, one obtains:
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• from which T0 and t2 are obtained
• but, except if O is entirely in the form of O++ in the nebula,
• t2(H+) ≠ t2(O++)
• T0(H+) ≠ T0(O++)
Visualisation of energy requirements
The simplest example:
for t2 = 0.04 and T0 = 10000K, f =1 implies T1 = 12000K and T2 = 8000K
black log of heating rate in arbitrary units
red: log of cooling rate in the O++ zone
By shifting the heating curve up and down one
understands how Te varies with energy input
t2 = 0.04 requires D log G = 0.3, ie a factor 2
difference in heating rates between regions 1
and 2 !
What fluctuates?
T0 ( N i ) 


Te N i N e dV
N i N e dV
 (Te  T0(Ni )) Ni Ne dV
2
t 2 (Ni ) 
T02  Ne Ni dV
Te ?
• Natural gradients in photoionized nebulae are small
• except at high metallicities
• (Stasinska 1980, Garnett 1992, Kingdon & Ferland 1995, Perez 1997)
Ne ?
• In high density clumps collisional dexcitation increases Te with respect to the
ambient medium Kholtygin 1998, Mathis et al. 1998
• (this is not sufficient to explain t2 ~ 0.04)
• densities above 105cm-3 boost [OIII] 4363/5007 (Viegas & Clegg 1994)
• but there is no evidence of such high densities in the O++ zones
What fluctuates?
Ni ?
• Te is lower in C-rich zones (Torres-Peimbert et al 1990)
• The O++ discrepancy between CEL and ORL requires the existence of O-rich zones
(Stasinska 1998, Liu et al 2000, 2001, Péquignot 2001)
is photoionization the only heating source in
photoionized nebulae?
• In a number of nebulae, classical photoionization models produce T[OIII] lower than
observed
• Giant HII regions: Campbell 1990, Garcia-Vargas et al 1997, Stasinska &
Schaerer 1999, Luridiana et al 1999, Luridiana & Peimbert 2001
• PNe: Peña et al 1998
• Additional energy sources have been proposed:
• Shocks (Peimbert et al 1991)
• Conduction fronts (Maciejewski et al 1996)
the ORL /CEL discrepancy
Expected properties of optical recombination lines (ORLs)
• their emissivity is roughly proportional to Te-1
• they should give correct abundances with respect to H
• even in presence of temperature fluctuations
ORL abundances versus CEL (collisionally excited lines) abundances
• ORL abundances are larger than CEL abundances by important factors
• Wyse 1947, Peimbert et al 1993, Liu et al 1995 (O) , Kaler 1986 (C)
• Esteban et al 1998, Liu et al 2000, 2001 (C,N,O)
C++ 1909 / O++ 5007 versus
C++ 4267 /O++ 5007 in planetary nebulae
Compilation Rola & Stasinska 1994
ORL versus CEL abundances
ionic abundances in the planetary nebula NGC 6153 Liu et al 2000
invoked causes of ORL-CEL discrepancy
Faintness of the ORLs
• biased measurements
• flux calibration is difficult over a large
dynamical range
• they may suffer from blends
Heavy element recombination coefficients
are not reliable
Temperature fluctuations
Density condensations
•no (from high S/N spectroscopy)
•no: [OIII]4931/[ [OIII]4959 agrees with theory: 4
10-4 Mathis & Liu 1999
•no (from echelle spectra )
•have been recomputed with the R-matrix method
Storey 1994
•ORL abundances from numerous transitions are
in agreement
•t2 explaining ORL-CEL discrepancy >> t2
explaining Te[OIII] -Te(BJ)
•IR-CEL abundances are consistent with optical CEL abundances Liu 2000,2001
•no (high order Balmer lines) Liu 2000...
invoked causes of ORL-CEL discrepancy
Faintness of the ORLs
• biased measurements
• flux calibration is difficult over a large
dynamical range
• they may suffer from blends
Heavy element recombination coefficients
are not reliable
•no (from high S/N spectroscopy)
•no: [OIII]4931/[ [OIII]4959 agrees with theory: 4
10-4 Mathis & Liu 1999
•no (from echelle spectra )
•have been recomputed with the R-matrix method
Storey 1994
•ORL abundances from numerous transitions are
in agreement
Recombination coefficients computed so far do not include
dielectronic recombination for n > 10 which is likely to be efficient at Te > 20 kK
Temperature fluctuations
•t2 explaining ORL-CEL discrepancy >> t2
explaining Te[OIII] -Te(BJ)
•IR-CEL abundances are consistent with optical CEL abundances Liu 2000,2001
Density condensations
•no (high order Balmer lines) Liu 2000...
Chemical inhomogeneities: they require super metal rich inclusions with solar C/N/O/Ne
Liu 2000, Tsamis 2003
possible origin for chemical inhomogeneities
in planetary nebulae
•
•
ejecta from the central star ?
photoevaporating planetesimals or planet debris ?
possible origin for chemical inhomogeneities
HII regions
droplets containing matter from supernova ejecta
Tenorio Tagle 1996
giant HII region
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galactic disk
superbubbl
e
t=0
t=1-40 Myr
warm oxygen-rich cloudlets
cold oxygen-rich
cloudlets
cold oxygen-rich droplets
t=100 Myr
galactic
disk
Stasinska
Tenorio Tagle
Rodriguez
Henney
2007
with a
spray
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galactic
disk
super
supernova
shell
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galactic
fountain
galactic
disk
no more
supernova
t=40
-100
Myr
neutral oxygen-rich
droplets
ionized
oxygen-rich
droplets
ionized ISM
fully mixed HII gas
new ionizing stars
The t2 problem
and the ORL/CEL discrepancy
are still a subject of debate
The role of shocks
The effects of shocks on emission line spectra as compared to stellar ionization
1. high densities due to gas compression
2. possible presence of highly ionized species (He++)
3. existence of important warm low-ionization zone (emitting [OI], [SII],
[OII] lines
4. higher Te , [OIII]4363/5007 enhanced
5. very high Te close to the shock, producing X-ray radiation
Effects of X-ray ionization as compared to stellar ionization
2, 3, 4
Effect of gas compression without shock
•
local compression of gas lowers U
•
effects 1, 3 are produced (without the need of shock heating)
NB
•
•
what is often attributed to shocks may actually be only due to compression
photoionization is much more efficient than shocks to ionize gas
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What can emission lines
tell us?
a lot !
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