Astrophysics of Galaxies 2019-2020 - Lecture V: The Spectral Energy Distribution of Galaxies Star Formation Rate, Ionized ISM
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Lecture V: The Spectral Energy Distribution of Galaxies Star Formation Rate, Ionized ISM Astrophysics of Galaxies 2019-2020 Stefano Zibetti - INAF Osservatorio Astrofisico di Arcetri Lecture V
Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V Star Formation Rate (SFR) ✤ SFR= dM*/dt ✤ galaxies evolve, and not only passively ✤ rely on short-lived stars to measure SFR ✤ if they are “observed” (directly or indirectly), they must have formed less than ∆t time ago
Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V Why short-lived stars trace SFR
Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V How can we “count” short lived stars? ✤ Mass-tsurv relation ✤ Mass-Teff relation ✤ ⇒ follow the UV photons! ✤ t-M* “uncertainty principle”: the more “instantaneous” the SFR measurement, the largest the extrapolation in mass
Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V The fate of UV photons ✤ The highest-energy UV photons (E>13.6 eV, i.e. 912Å) ionize the atoms in the birth clouds of stars, cascade down the recombination ladders n=7 n=6 ✤ Hydrogen series (mainly Balmer, Paschen, n=5 Brackett; Lyman is resonant, not optically n=4 thin) n=3 ✤ other nebular lines (e.g. [OII]) are excited by hot free electrons n=2 ✤ Lower energy UV photons escape and produce UV luminosity, possibly attenuated by dust n=1 ✤ All can heat up dust and excite ISM molecular emission (like PAHs)
are expected: then cooling in the O zone is dominated by collisional excitation of fine Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V structure lines in the ground level of O++ , while the absence of fine structure lines in the ground level of O+ forces the temperature to rise in the outer zones (Stasińska 1980a, Garnett 1992). Hydrogen recombination lines – For a given T⋆ , Te is generally lower at higher metallicity. – For a given metallicity, Te is generally lower for lower T⋆ . – For a given T⋆ and given metallicity, Te increases with density in regions where n is larger than a critical density for collisional deexcitation of the most important cooling lines (around 5 102 – 103 cm−3 ). 1.3. Line intensities ✤ Ionized In ISM conditions is optically prevailing in PNethin and to H iinon-resonant linesemission regions the observed (not solines manyare atoms/ions optically excited thin, above except for resonance the ground lines state), i.e.such NOT as H Lyα, C ivλ1550, N vλ1240, Mg iiλ2800, Si ivλ1400, and some helium lines. Also the fine structure IR lines could be optically thick ✤ Intensity in compact ofH iihydrogen recombination regions or giant lines almost H ii regions (however, independent the velocity fields areof T, ~proportional to generally the number such that this of recombinations, is not the case). The hence to most fact that the number of used of the lines ionizing photons if for abundance equilibrium holds determinations and all are optically thinionizing photons makes their absorbed use robust (“case B”) and powerful. ✤ Dust extinction can be estimated based on recombination line ratios: the theoretical ratio is altered by wavelength dependent extinction; once an extinction curve is assumed, one just needs the extinction ratio at two different λ to get the optical depth ✤ Intrinsic line luminosity can be derived ⇒ ionizing flux ✤ SPS models provide the conversion from ionizing flux to SFR ✤ Only stars with M>10M⦿ and life-time
L’estinzione ad unaStefano certa Zibetti -λ può INAF essere OAArcetri - Astrophysicsscritta come of Galaxies - Course A(λ) 2019/2020 =V f(λ) AV dove f(λ) è - Lecture la curva di estinzione, che dipende unicamente dalle proprietà dei grani, mentre AV è l’estinzione in banda V che dipende anche dalla quantità di Extinction from “Balmer decrement” polvere lungo la linea di vista. Spesso si utilizzano le righe dell’idrogeno (di cui si conoscono i rapporti), es. il “decremento di Balmer” Hα/Hβ, per stimare l’estinzione ef f L(H ) (H 0 , Te ) ⇥ H ef f 3 per Te = 104 K L(H⇥) H⇥ (H 0, T ) e ⇥oss ⇥em F (H ) F (H ) = 10(AH⇥ AH )/2.5 F (H⇥) F (H⇥) ⇥oss F⇤ (H ) ⇥3 10AV [f (H⇥) f (H )]/2.5 F⇤ (H⇥) data una curva di estinzione f(λ) (spesso assunta) si ricava AV f(λ) is the assumed extinction curve Sunday, December 16, 12 AV is the only left unknown
Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V Recombination lines across the Hubble sequence: a sequence of SFH intensity Gavazzi et al. (2002)
Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V The Astrophysical Journal, 734:82 (17pp), 2011 June 20 Izotov & Thuan The plethora of NIR recombination lines Figure 1. 3.5 m APO/TripleSpec NIR spectrum of II Zw 40 in five orders. In each panel, the noisy regions of the upper spectrum are omitted. They are caused by insufficient sensitivity or strong telluric absorption. The flux scale on the y-axis corresponds to the upper spectrum. The lower spectrum is downscaled by a factor of 50 as compared to the upper spectrum. It is shown for the whole wavelength range in each order.
Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V
Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V SFR from forbidden lines: [OII] ✤ very useful to extend to higher z but staying in the optical window ✤ forbidden lines in general are not directly coupled to ionizing radiation (see next slides), BUT [OII] is well behaved and can be calibrated empirically. ✤ metal abundance has a relatively small effect on the [OII] calibration, over most of the abundance range of interest (0.05 Z⊙ ≤ Z ≤ 1 Z⊙)
Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V A simple model to interpret galaxy spectra 1605 Attenuation DaCunha+08 “Cirrus” component, heated by diffuse ISRF “Warm” component, directly heated by Direct UV xamples of spectral energy distributions obtained by combining the infrared models of Table 1 with attenuated stellar population hotspectra stars (non-ionizing, λ>912Å) µ µ tot g to the same contributions by dust in stellar birth clouds (1 − f ) and in the ambient ISM (f ) to the total energy L absorbed and reradiated by d n 2.3). (a) Quiescent star-forming galaxy spectrum combined with the ‘cold’ infrared model of Table 1; (b) normal star-forming galaxy spectrum th the ‘standard’ infrared model of Table 1; (c) starburst galaxy spectrum combined with the ‘hot’ infrared model of Table 1 (see text for details rameters of the stellar population models). Each panel shows the unattenuated stellar spectrum (blue line), the emission by dust in stellar birth
Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V SFR from UV continuum ✤ Use SPS models assuming ✤ IMF ✤ SFH: adapt to the galaxy to study ✤ typically constant over the time traced by the66 MEURER, HECKM UV bright stars ✤ shorter if considering a starburst ✤ Integrated spectrum 1500-2500Å dominated by stars M≳5 M⦿ ✤ timescale ~108 yr ✤ substantial extrapolation to total mass ✤ Large sensitivity to dust extinction, typically 0-3 mag, but can be much much more for dusty starburst! ✤ β-A UV relation: empirically calibrated via IRX, assuming all IR is re-radiated flux missing from the UV (Meurer, Heckman & Calzetti 1999) FIG. 1.ÈRatio of FIR to UV Ñux at 1600 A! compared to UV spectral
Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V SFR from IR continuum ✤ Total IR luminosity LTIR is a good proxy for SFR as long as 1994ApJ...429..582C ✤ most of the UV photons are absorbed and re-radiated by dust ✤ dust heating is predominantly done by young hot stars ✤ note that the absorption cross section of dust increases at shorter λ Calzetti, Kinney & Storchi-Bergmann (1994)
Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V Proxies to TIR (and SFR) 1964 Figure 5. Relation between observed 24 µm infrared luminosity and observed M. Galametz et al. Kennicutt+09 Hα luminosity (uncorrected for attenuation) for galaxies in the SINGS and MK06 samples. Symbol shapes are coded by dominant emission-line spectral type: circles for H ii region like spectra, crosses for spectra with strong AGN signatures, and crossed circles for galaxies with composite spectra. The solid line shows a relationship with linear slope, to illustrate the strong nonlinearity in the observed relation. The axis label at the top of the diagram shows the approximate range of SFRs (absent a correction for dust attenuation at Hα) for reference. calibrations as functions of various properties of the galaxies and their SEDs. 4.1. Combinations of Hα and 24 µm Measurements As an introduction, it is instructive to compare the consis- tency of Hα and IR SFR measures before any corrections for attenuation are applied. Figure 5 compares the observed Hα Galametz+13 luminosities of the SINGS and MK06 galaxies with their corre- sponding 24 µm (νfν ) luminosities. When converting fluxes to coded by 12+log(O/H) luminosities, we have used distances listed for SINGS galaxies in Table 1, and from MK06 in Table 2. Both sets of distances assume H = 70 km s−1 Mpc−1 with local flow corrections.
Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V 878 CALZETTI ET AL. Vol. 666 Mid IR: PAHs and hot dust A simple model to interpret galax ✤ Dominated by hot dust and PAH DaCunha+08 emission, excited by soft UV (see e.g. Draine+07) Fig. 3.— (a) LSD at 8 !m, S8 !m;dust , as a function of the extinction-corrected LSD at Pa", SPa" ;corr , for the 220 H ii knots in the 33 galaxies for which photometric measurements have been obtained. LSDs are averaged over 1300 photometric apertures. The 8 !m emission is stellar continuum subtracted (x 3.1). Data points are divided into three metallicity bins: high (red filled triangles), intermediate ( green stars), and low (blue asterisks) oxygen abundance (x 2). Filled black squares mark the PAH luminosity: dependence on local starbursts from the sample of Engelbracht et al. (2005) (x 5.1). The 3 # error bars are indicated for each data point. Open black stars indicate the location of the ✤ nonthermal sources (Seyfert 2s or LINERs; x 4.1), and open black circles indicate extended background sources. The best-fit line through the high-metallicity (red ) data points is shown as a solid line, while the dashed line is the linear fit through the same data points with fixed slope of 1. (b) Histogram of the deviation of the H ii knot data in panel a from the best-fit line through the high-metallicity data (the solid line in panel a). The values of the best-fit coefficients are c ¼ 0:94 " 0:02 and d ¼ metallicity 4:80 " 0:85 (eq. [2]). Three separate histograms are shown, for high- (red ), intermediate- (green), and low-metallicity (blue) data. The intermediate- and low-metallicity histograms have been multiplied by a factor of 2 to make them visible. ✤ 8%m must be decontaminated from 4.2. Uncertainties in the Photometric Measurements tial misregistration of the multiwavelength images. The vari- ance on the image background is derived in each case from the The uncertainties assigned to the photometric values at each original-pixel-size images. The impact of potential background stellar continuum wavelength and for each galaxy are the quadrature combination of four contributions: Poisson noise, variance of the background, under- or oversubtractions varies from galaxy to galaxy and also depends on the relative brightness of the background and the photometric calibration uncertainties, and variations from poten- sources. The effect of potential misregistrations has been evaluated 878 CALZETTI ET AL. Vol. 666 Figure 5. Examples of spectral energy distributions obtained by combining the infrared models of Table 1 with attenuated corresponding to the same contributions by dust in stellar birth clouds (1 − f µ ) and in the ambient ISM (f µ ) to the total energy Ldtot dust (Section 2.3). (a) Quiescent star-forming galaxy spectrum combined with the ‘cold’ infrared model of Table 1; (b) normal sta Calzetti+07 Fig. 3.— (a) LSD at 8 !m, S8 !m;dust , as a function of the extinction-corrected LSD at Pa", SPa" ;corr , for the 220 H ii knotscombined measurements have been obtained. LSDs are averaged over 1300 photometric apertures. Fig. Herrero 4.—Same The d ¼ and et al. #6:88 8 !m as (2006) Fig. 3, but for the (blackisasterisks; emission LSD inatthe stellar continuum !m, 2433 about(xthe S24with galaxies x 5.2). Thesubtracted values of (x In the !m . for parameters ‘standard’ addition which the3.1). photometric parameters (c, d )are Data points ofmark infrared to the same data points asmodel of Table Fig. 3, panel 1; (c) a also reports the starburst galaxy LIRGs from the sample spectrum of Alonso- combined with the ‘hot’ infrared model of in the horizontal axis of panel b are given in eq. (3) and are c ¼ 1:23 " 0:03 and thethestellar population models). Each panel shows the unattenuated stellar spectrum (blue line), the emis divided into three metallicity bins: high (red filled triangles), intermediate ( green stars), low"(blue 0:97.asterisks) oxygen abundance 2). Filled black squares local starbursts from the sample of Engelbracht et al. (2005) (x 5.1). The 3 # error bars are indicated for each data point. Open clouds (green black stars line), indicate the emission the location of the by dust in the ambient ISM (red line) and the total emission from the galaxy, corresponding to the s nonthermal sources (Seyfert 2s or LINERs; x 4.1), and open black circles indicate extended background sources. The best-fit line through the high-metallicity (red ) data spectrum points is shown as a solid line, while the dashed line is the linear fit through the same data points with fixed slope of 1. (b) Histogram and theof total of the deviation the H iiinfrared knot data emission (black line).
Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V FIR-radio correlation de Jong+1985; Helou+1985: GHz vs. 60µm or FIR ✤ LFIR∝SFR, FIR-radio correlation ⇒ LGHz∝SFR : why?? ✤ GHz is synchrotron, i.e. relativistic CR electrons + magnetic field ✤ CR electrons are produced in SN explosions ✤ “Electron Calorimeter” theory: synchrotron cooling timescale shorter than the escape time: tsynch
otropically, whereas the UV light escapes along preferred di- Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V ctions (such as out of the plane of a disk galaxy). Future works ll be able to address point 1 through more detailed model- Bolometric SFR indicators: UV-IR g of UV/optical / IR SEDs and will be able to ascertain to what tent averaging over inclination angles effectively addresses int 2. In practice, we estimate the SFR using a calibration de- ved from the PEGASE stellar population models (see Fioc & ✤ Combine emerging UV flux with flux re-emitted in the IR by the dust occa-Volmerange 1997 for a description of an earlier version the model), assuming a 100 Myr old stellar population with nstant SFR and a Kroupa (2001) IMF: ✤ Bell+05: ¼ 9:8 ; 10$11 (LIR þ 2:2LUV ); ð1Þ M# yr$1 here LIR is the total IR luminosity and LUV ¼ 1:5#l#;2800 is a ugh estimate of the total integrated 1216–3000 8 UV lumi- sity, derived using the 2800 8 rest-frame luminosity l#;2800 om COMBO-17. The factor of 1.5 in the 2800 8–to–total UV nversion accounts for the UV spectral shape of a 100 Myr old pulation with constant SFR. The SFR is derived assuming at LIR reflects the bolometric luminosity of young, completely scured populations and that LUV reflects the contribution of
Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V SFR indicators ✤ Look up most updated references! ✤ Main review Kennicutt ARAA 1998 ✤ Calzetti+07: mid-IR ✤ Calzetti+10: monochromatic FIR
Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V Properties of the ionized gas ✤ See Osterbrock (1989)! Also review by Stasinska (2002, arXiv:astro- ph/0207500) - Dopita & Sutherland “Astrophysics of the Diffuse Universe” ✤ Key quantities: ✤ T e ✤ n e ✤ abundances ✤ Key processes: ✤ photon excitation and ionization ✤ collisional excitation and de-excitation (electrons on atoms/ions) ✤ recombination ✤ radiative de-excitation
ijl If the Stefano density is sufficiently high, some collisional deexcitation may occur and cooling Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V is reduced. In the two-level approach one has: Lijl = n e n(X j )ne qijlhνijl (1/(1 + ne (q12 + q21 )/A21 ). (1.20) Ionized gas diagnostics coll i So, in a first approximation, one can write that the electron temperature is determined by ! ijl G=L= Lcoll , (1.21) ijl where G is given by Eq. (1.18) and Lijl coll by Eq. (1.20). The following properties of the electron temperature are a consequence of the above equations: – Te is expected to be usually rather uniform in nebulae, its variations are mostly ✤ Ionized ISM is optically thin to non-resonant lines (not so many determined by the mean energy of the absorbed stellar photons, and by the populations atoms/ions excited of the above main cooling ions. the ground It is only state) (over solar) that large Te gradients at high metallicities are expected: then cooling in the O++ zone is dominated by collisional excitation of fine structure lines in the ground level of O++ , while the absence of fine structure lines in the ✤ Intensity of recombination ground level of O+ forceslines almost to the temperature independent rise in the outer of T,(Stasińska 1980a, zones Garnett proportional 1992).number of recombinations, hence of ionizing to the – For a given T⋆ , Te is generally lower at higher metallicity. photons if equilibrium holdsTe is generally lower for lower T⋆ . – For a given metallicity, – For a given T⋆ and given metallicity, Te increases with density in regions where n is larger than a critical density for collisional deexcitation of the most important cooling ✤ Intensity of collisionally lines (around 5 102 – excited 103 cm−3 ).lines depends on T and n 1.3. Line intensities ✤ Resonant lines haveprevailing In conditions very complicated behavior: in PNe and H ii regions the observed emission lines are optically thin, except for resonance lines such as H Lyα, C ivλ1550, N vλ1240, Mg iiλ2800, Si ivλ1400, and some helium lines. Also the fine structure IR lines could be optically thick in compact H ii regions or giant H ii regions (however, the velocity fields are generally
Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V Gas (metallicity) diagnostics ✤ Use collisionally excited lines of atoms/ions of given element ✤ O is the most popular, in fact estimates of gas metallicity are often given as 12+log(O/H), with solar [O/H] being 8.69 ✤ Forbidden lines: opposite to photon excited levels, collisionally excited levels do not require to follow the quantum selection rules relative to the ground state ✤ Dependence on density and temperature, not only element/ion abundance! ✤ “Direct” measurements require many lines, not always available for moderate SNR spectra ✤ Revert to indirect methods or “strong line methods”
Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V 5.1.4 Two- and three level At Key concepts of collisionally excited lines As a result of very low density and weak atoms of any particular element and ioni (see e.g. Dopita & Sutherland’s 5.1.4 Two- and three level Atomsbook) excited states are populated Page 27 through col As a result of very low density and weak radiation fields, (From theDopita vast majority of the & Sutherland, Astrophysic atoms of any particular element and ionization state reside in the ground state and the excited states are populated through collisions with electrons. ✤ Excitation rate depend on cross (From Dopita & Sutherland, Astrophysics of the Diffuse Universe) section Cross Section: ✤ Cross sections depends on the h2 12 12 ( E ) , 12 21 energy of the exciting electron 8 me E g1 (see figure) Energy distribution of the electrons ✤ Energy of the electrons depends(with mean density n e ): on temperature, following the 2ne 1/ 2 E Maxwell-Boltzmann distributionf (E) 1/ 2 3/ 2 E exp dE (kT ) Collisional Rate R 12kT ne N1 12 C12 N1 Collisional Rate R 12 ne N1 12 C12 N1 : 2E 2 4 1/ 2 E 2E R12 ne N1 12 ( E ) f ( E )dE ne N1 TR121/ 2 ne12N1exp 12 (12E ) f ( E )dE ne N E12 me kme3 g1 E12 kT me v
Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V Page 28 Line emission of the two-level Atom Low density limit: R12 C12 N1 N 2 A21 equilibrium set by radiative de-excitation rate 1/ 2 4 12 2 1 1/ 2 12 E12 N2 ne N1 ne N1 A T 21 exp A21 kme3 g1 kT Emission line flux: 1/ 2 4 2 2 1/ 2 12 E12 F12 E12 A21 N 2 n i e E12T exp kme3 g1 kT where i =N i /n e is the relative abundance of the ion i considered. For low temperatures the exponential term dominates because few electrons have energy above the threshold for collisional excitation, therefore the line rapidly fades with decreasing temperature. At high temperatures the T -1/2 term controls the cooling rate, so the line fades slowly with temperature. The maximum line flux is reached at T=E12 / k . To measure abundances i one needs to know n e and Te . IMPRS Astrophysics Introductory Course (credits to Ralf Bender et al.) Fall 2009
Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V Page 29 Line emission of the two-level Atom High density limit : Boltzmann equilibrium (LTE) N2 g2 E12 exp N1 g1 kT Emission line flux: g2 E12 F12 E12 A21 N 2 i E12 A21 ne exp g1 kT Note that the line flux scales with n e and it tends to a constant value at high temperatures. Critical density where A 21 R12 : 1/ 2 1/ 2 4 A21 g 2T 2 ncrit 3 (see derivation in next slide!) 12 kme In a log n e -log F plot the slope changes from 2 to 1 at this density. IMPRS Astrophysics Introductory Course (credits to Ralf Bender et al.) Fall 2009
2 e 1 e 1 3 21 A21 km g1 Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V e kT N 2 limit High density g 2 : Boltzmann E12 equilibrium (LTE) Emission line flux:exp Low density N1 g1 kT N2 g2 E212 4 1/ 2 E12 exp Emission F12 E12 A21 N 2 n 2 line flux: E T 1/ 2 12 exp N1 g1 i e kTkme3 12 g1 kT g2 E12 Emission where /n12line i =N iF e is Eflux: 12 A21 N 2 the i E12 A21 ne relative abundance = exp of the ion g1 i considered. For low temperatures the exponential term dominates because few F12 Ehave Note that the line flux g scales 2 with forexp n E kT and 12 it ten electrons 12 A21 N 2 above energy i E12 A21threshold the ne collisional e excitatio the line at rapidly high fades temperatures. g1 kT with decreasing temperature. At high tempe High density Note -1/2 that the line flux scales with n the T termCritical controlsdensity where the cooling rate,Aso e and 21 theRline 12 it slowly : fades tends wt at high temperatures. The maximum line flux is reached 1/ 2 2 at T=E 4 / 1/k2. A21 g 2T 12 ncrit To measure abundances i one needs to3know n e and Te . Critical density where12 A 21 kmRe12 : In a log 1/ n 2 -log F plot 1/ 2 slope changes from 2 the Course 4 A gT IMPRS Astrophysics Introductory e 2
Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V Page 30 Three-level Atoms, Ions with E 23 E12 Low density limit: C12 N1 N 2 A21 , C13 N1 N 3 A31 Examples: [OII]3726,3729 F31 E31 A31 N 3 E31C13 31 E23 g3 [SII]6731,6716 Flux ratio: exp F21 E21 A21 N 2 E21C12 21 kT g2 [ArIV]4740,4711 High density limit: Boltzman ratio N 3 / N 2 g3 / g 2 Only if F31 A31 g3 ~1 Flux ratio: DENSITY INDICATOR E23/kT F21 A21 g 2 is small! IMPRS Astrophysics Introductory Course (credits to Ralf Bender et al.) Fall 2009
Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V Key points for ne determinations ✤ Same “ground” level of the same ion: remove the i dependence ✤ E32
Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V Page 28 Line emission of the two-level Atom Focus onlimit:Temperature Low density R C N N A 12 12 1 1/ 2 2 21 4 12 2 1 1/ 2 12 E12 N2 ne N1 ne N1 A T 21 exp Page 29 A21 kme3 g1 kT Line emission of the two-level Atom ✤ LowEmission density:line flux: High density limit : Boltzmann 4 1/ 2 equilibrium (LTE) 2 2 1/ 2 12 E12 FN 122 E12 g 2A21 N 2 En i e 12 E12T exp exp kme3 g1 kT N1 g1 kT where i =N i /n e is the relative abundance of the ion i considered. ✤ High density: Emission line flux: For low temperatures the exponential term dominates because few electrons have energy above thegthreshold 2 E12collisional excitation, therefore for F12 E12 A21 N 2 i E12 A21 ne exp the line rapidly fades with decreasingg1 kT temperature. At high temperatures Note the thatterm T -1/2 the controls line fluxthescales withrate, cooling n e so andthe it line tendsfades to aslowly constant value with temperature. at high The temperatures. maximum line flux is reached at T=E12 / k . Critical To density measure where A 21 abundances i oneR : needs 12 to know n e and Te . 1/ 2 1/ 2 4 A21 g 2T 2 nIMPRS crit Astrophysics Introductory Course 3 Fall 2009 12 kme
Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V Page 31 Three-level Atoms, Ions with E12 E 23 4363 (2321) Page 28 Line emission of the two-level Atom Low density limit: R12 C12 N1 N 2 A21 N2 ne N1 12 ne N1 2 4 1/ 2 A211T 1/ 2 12 5007/4959 exp E12 A21 kme3 g1 kT Emission line flux: Examples: 1/ 2 F12 E12 A21 N 2 n2 i e 2 4 E12T 1/ 2 12 exp E12 [OIII]4363,5007/4959 kme3 g1 kT Low [NII]5755,6583/6548 where idensities: F32of the ion =N i /n e is the relative abundance E32i considered. C13 A32 [SIII]6312,9532/9069 N ForC N 3 ( A32the exponential low temperatures A31 ) termF21dominates E21because C12 A few 1 13 31 electrons have energy above the threshold for collisional excitation, therefore F32 E32 A32 13 E32 Nthe1C N A N A Flux ratio: exp line 12 rapidly fades 3 32 with decreasing 2 21 C temperature. At high E temperatures 13 23 13 F21 E21 A31 kT -1/2 exp 12 Nthe1 T Nterm 2 controls the cooling rate, so the line fades slowly with temperature. N3 1 C kT The maximum line flux is reached at T=E12/ k . 12 12 Normalization To measure abundances i one needs to know n e and Te . TEMPERATURE INDICATOR Equilibrium: same total transition rates from/to levels IMPRS Astrophysics Introductory Course Fall 2009 IMPRS Astrophysics Introductory Course (credits to Ralf Bender et al.) Fall 2009
Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V Key points for Te determinations ✤ Same “ground” level of the same ion: remove the i dependence ✤ E32~E21: ✤ C13
Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V “Strong line methods” ✤ ne, T, abundance and extinction require 4 independent line ratios (sensitive to these quantities!) ✤ only 4 strong lines typically well measured (Pagel 1979): Hα, Hβ, [OII] λ3727 and [OIII] λ5007 ⇒ 3 independent line ratios ✤ Hβ/Hα Balmer decrement ⇒ extinction ✤ empirical dependence of T (and ne) on metallicity ✤ Alternatively: simultaneous fit of several lines based on photoionization models (e.g. Tremonti et al. 2004)
Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V R23 method G THE R23 DEGENERACY R23 y on the= (½O ii" k3727 þ ½O iii" kk4959; 5007)/H! line with metallicity; R23 gives both a low metallicity estima (for ✤ a branches, Two discussion see, e.g., KK04). Additional line ratio high and low Z ✤ Requires calibration using s ( Izotov et al. 2004, 2006a; Kniazev et al. 2003, 20 direct methods limited emission-line surveys because they are intrinsi ✤ needs a rough estimate of metallicity to choose the van Zee 2000). For the purpose of investigating the u branch! other Z-sensitive metallicity galaxy sample described in Kewley et al. (20 line ratios must be axy sample. adopted (see Kewley & Ellison 2008) n [ N ii]/ H " or [ N ii]/[O ii] metallicity calibration becau ed on [N ii]/H" or [N ii]/[O ii] and the calibration base ranch of R . For example, an [ N ii]/H" metallicity c
1981PASP...93....5 Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V (1981) BPT diagram(s) power law PNe ✤ Insensitive to reddening and absolute flux HII regions shock calibration [OIII]/Hβ ✤ Insight into heating mechanisms [NII]/Hα
Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V The physics of BPT diagrams ✤ Metallicity [O/H] ✤ Shape of the spectrum (“hardness”) ✤ Ionization parameter of the radiation field: density of ionizing photons (Lyman H + IE → H+ + e− IE = 13.6 eV N + IE → N+ + e− IE = 14.5341 eV continuum) over electron density O + IE → O+ + e− IE = 13.6181 eV O+ + IE → O++ + e− IE = 35.1185 eV 0-age SSP [O/H] ➘ 4Myr continuous u Sánchez+2015, CALIFA
Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V The physics of BPT diagrams ✤ Metallicity [O/H] ✤ Electron density: increases rate of ✤ Dust extinction collisional excitation 0-age SSP 0-age SSP, fixed Ne [O/H] ➘ [O/H] ➘ Sánchez+2015, CALIFA
Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V Sánchez+2015, CALIFA
galaxies fall into one of two categories: (1) galaxi Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V < 0.73/[log([O I]/Hα) + 0.59] + 1.33. (3) Seyfert region in either the [S II]/Hα or the [O I]/Hα the LINER region in the remaining ([O I]/Hα or [S te galaxies lie between the Ka03 and Ke01 classifi- BPT diagrams he [N II]/Hα versus [O III]/Hβ diagram: /Hα) − 0.05] + 1.3 < log([O III]/Hβ), (4) or (2) galaxies that lie in the composite region (belo in the [N II]/Hα diagram but that lie above the Ke the [S II]/Hα or the [O I]/Hα diagram. According to this scheme, our 85 224-galaxy /Hα) − 0.47] + 1.19 > log([O III]/Hβ). (5) 63 893 (75 per cent) star-forming galaxies, 24 alaxies lie above theline ✤ Emission Ke01classification classification line of on the galaxies: the role Seyferts, 6005of (7 the AGN! per cent) LINERs, and 5870 (7 Hα,✤and [O I]/Hα diagnostic diagrams and above the posites. The remaining galaxies are ambiguous T-Z sequence for HII galaxies in line on the [S II]/Hα and [O I]/Hα diagrams, that is, [NII] BPT 8 per cent). ✤ Radiation “hardness” and shocks diagnostics to the right of the line 1.5 (a) (c) AGN Seyfert Seyfert 1.0 LOG ([OIII]/Hβ) 0.5 0.0 LINER HII HII LINER -0.5 -1.0 Comp -2.0 -1.5 -1.0 -0.5 0.0 0.5 -1.0 -0.5 0.0 0.5 -2.0 -1.5 -1.0 -0.5 0.0 LOG ([NII]/Hα) LOG ([SII]/Hα) LOG ([OI]/Hα) Kewley et al. (2006) [also Kauffmann et al. 2003]
Stefano Zibetti - INAF OAArcetri - Astrophysics of Galaxies - Course 2019/2020 - Lecture V Emission lines from AGN ✤ Type I: ✤ Broad permitted lines directly from BLRs (Balmer lines, MgII, CIV, Lyα) ✤ FWHM>2000 km/s ✤ Narrow forbidden lines from NLRs ([OIII], [NII], [OII] etc.) ✤ Type II: ✤ narrow permitted and forbidden lines, FWHM ~ several 100 km/s ✤ Possible mixing with emission from HII regions and diffuse ionized gas
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