Spectral Energy Distribution of galaxies

Transcription

1 Spectral Energy Distribution of galaxies Paola Santini PhD in Astronomy, Astrophysics and Space Science A.A

2 Key points lecture 1 Multiwavalength astronomy: 1. Gives a complete view of the galaxy as a whole 2. Allows to study galaxy evolution with redshift 3. Allows to estimate redshift and physical parameters Bimodality in galaxy properties (holds on average!): o red colours high luminosity large stellar mass typically low SF activity dense environments elliptical/s0 morphology supported by random motions o blue colours low luminosity low stellar mass on going star formation low density environments spiral/irregular morphology supported by rotation Galaxies follow empirical scaling relations connecting the luminosity to the stellar motions (both related to the depth of the DM potential well): o spirals follow the TF o ellipticals follow the FJ, and more in general the FP LF and MF are useful tools to study galaxy populations in a statistical sense 1<z<3 : peak of star formation activity of the Universe

3 Lecture 2 Range: ~ 300 >22000 Å Source of emission/absorption: stars Stellar spectra Initial Mass Function Star Formation History Models of stellar population synthesis Spectral evolution and main spectral features Colours of stellar populations Physical parameters estimation from the spectral continuum

4 Galaxy integrated spectra Integrated spectra: the entire galaxy is observed simultaneously, integrated emission of the spectra of billions of stars (+ gas and dust) Elliptical Spiral

5 Galaxy integrated spectra Integrated spectra: the entire galaxy is observed simultaneously, integrated emission of the spectra of billions of stars (+ gas and dust) Galaxy emission is complex, but we can reconstruct it from stellar emission, which is rather well known. We just need to sum up all the different stellar populations born at different epochs stellar population synthesis Ingredients to build a galaxy spectrum: 1. Stars 2. Initial Mass Function 3. Star Formation History

6 Stars HR diagram: every star occupies a position determined by its mass, and moves along this diagram depending on its age Colour! Absolute magnitude!

7 Evolutionary track Stars

8 Stars λ peak (Å) Stellar mass distribution is important!!

9 Stars Temperature (which depends on the mass) is important for 2 reasons: 1. it shifts the peak of black body emission Black body emission: radiation emitted by a body in thermal equilibrium (i.e. at constant temperature) B(ν,T) = 2hν 3 1 c 2 hν e B(λ,T) = 2hc 2 1 λ 5 hc e kt 1 ktλ 1 [ Wm 2 sr 1 Hz 1 ] [ Wm 2 sr 1 cm 1 ] Wien s law λ max T = 0.29cmK 2. it determines the atom ionization level at surface Photons with energy > E ion are absorbed by neutral atoms and ionize them. Hydrogen: E ion =13.6eV (λ~912å, T eff ~ K): if T eff >~ K: all H is ionized (no absorption); if T eff <~ K: H is neutral, so all photons at energy > E ion are absorbed and no light is emitted at λ<912å (H is very abundant, large cross section) Helium: E ion =54.4eV (λ~228å, T eff ~10 5 K): stars are not hot enough to completely ionize He He absorption at λ<~228å heavier elements: partially ionized absorption lines

10 Atomic transitions and series limit Hydrogen energy levels Transition (absorption or emission): change in energy level (ΔE=hν) of the electron (e.g. Hα 3 2 : the electron falls from n=3 to n=2 in the atom). Limit: limit to reach the continuum; if a photon has more energy than this threshold (shorter wavelength) it can ionize the atom (Lyman limit: 912Å; Balmer limit: 3646Å).

11 Stars λ peak (Å) Stellar mass distribution is important!!

12 Stars Luminosity! Temperature! Duration of MS phase! Mass!

13 Stars high M! high T! short age! Lyman break Balmer break Note the displacement of the BB peak. No emission at λ<228å (neutral He absorption) low M! low T! long age! No emission at λ<912å for stars with T<T ionh due to absorption from neutral H)

14 Stars low M! low T! long age! Low T allows the formation of molecules many metals absorption lines. Solar like star. The BB spectrum can be recognized. Several metals absorption lines. 4000Å break : caused by the absorption of Ca, Mg, etc. Note the displacement of the BB peak. No or few metal lines. Balmer absorption lines due to Hydrogen excitation. Balmer limit or Balmer break. high M! high T! short age! T>T ionh : very little H absorption. Absorption lines due to HeI and HeII excitation Stellar spectra also depend on metal abundance

15 Galaxy integrated spectra Integrated spectra: the entire galaxy is observed simultaneously, integrated emission of the spectra of billions of stars (+ gas and dust) Galaxy emission is complex, but we can reconstruct it from stellar emission, which is known rather well. We just need to sum up all the different stellar populations born at different epochs stellar population synthesis Ingredients to build a galaxy spectrum: 1. Stars Stellar mass is important!! (it determines L and T) 2. Initial Mass Function 3. Star Formation History Weigth function for stellar spectra

16 Stellar Initial Mass Function (IMF) Mass distribution of stars at birth φ(m) = d(n /V ) dm = dn dm When building a galaxy spectrum, the IMF represents a weight function for the components of different stellar mass of a stellar population (i.e. a group of stars born at the same time) Salpeter (1955) φ(m) m α α = 2.35 The total mass is dominated by low M stars Assumption: the IMF is universal (the same at all epochs and in all environments). Possible dependence of the IMF on time and/or metallicity and/or environment is an open problems of modern astrophysics.

17 Stellar Initial Mass Function (IMF) Though in principle it s just a matter of counting stars, it is not easily measured: need to observe young environments, otherwise giant stars will be dead need to observe a single population: SF regions (young) are dusty, observations are complex and only feasible for the largest stars globular clusters (most of the stars at the same stage of stellar evolution) are old and lack massive stars

18 Galaxy integrated spectra Integrated spectra: the entire galaxy is observed simultaneously, integrated emission of the spectra of billions of stars (+ gas and dust) Galaxy emission is complex, but we can reconstruct it from stellar emission, which is known rather well. We just need to sum up all the different stellar populations born at different epochs stellar population synthesis Ingredients to build a galaxy spectrum: 1. Stars Stellar mass is important!! (it determines L and T) 2. Initial Mass Function Weight function 3. Star Formation History

19 Star Formation History (SFH) Average mass of stars created per unit time as a function of time A gas cloud fragments and forms a first population of stars, whose mass is distributed according to the IMF. As time goes by, the initial stellar mix changes (more massive stars die within few Myr). At the same time, new stellar populations are created in subsequent star formation events. The evolution of the star formation rate (stellar mass/unit time) with time is described by the SFH. Therefore, the SFH regulates the stellar mix with time. The SFH is arbitrary, and to some extent random. We can only attempt simple parameterizations. Three simple cases: Instantaneous burst Constant star formation rate Exponential law SFH(t) = δ(t t 0 ) SFH(t) = const SFH(t) = τ 1 t /τ e (arbitrary normalization) (exponentially declining SFH are a good approximation)

20 Star Formation History (SFH) Real SFH may be much more complicated SFH according to hydrodynamical simulations SFH(t) = δ(t t 0 ) SFH(t) = const SFH(t) = τ 1 t /τ e Simha et al., 2014

21 Galaxy integrated spectra Integrated spectra: the entire galaxy is observed simultaneously, integrated emission of the spectra of billions of stars (+ gas and dust) Galaxy emission is complex, but we can reconstruct it from stellar emission, which is known rather well. We just need to sum up all the different stellar populations born at different epochs stellar population synthesis Ingredients to build a galaxy spectrum: 1. Stars Stellar mass is important!! (it determines L and T) 2. Initial Mass Function Weight function 3. Star Formation History Regulates the evolution of the stellar mix with time Let s put things together!

22 Stellar population synthesis Summing up the spectra of co eval stellar populations with different age Spectrum of a simple/ single stellar population (SSP) f λ SSP (t,z) = IMF (weight) (assumed universal) M max M min φ(m) f λ (m,t,z)dm Stellar spectrum (M, t and Z are fundamental parameters) Spectrum depends on metallicity: the surface temperature depends on the abundance of metals Integrated spectrum of the galaxy (composite stellar population) f λ gal (t,z) = t 0 f λ SSP (t t',z)sfh(t')dt' Rate at which new stellar populations are formed

23 Spectral evolution: instantaneous burst SFH(t) = δ(t t 0 ) 504Å (24.6eV): HeI ionization 228Å (54.4eV): HeII ionization 912Å (13.6eV): Lyman limit neutral H absorption 3646Å (3.4eV): Balmer limit maximum at age Gyr Red giants Immediately after burst Spectrum is dominated by high M stars (despite their low number) very bright and hot (short λ) emission at λ<912å (H is completely ionized and does not absorb photons) After a few Myr UV decreases: most massive stars disappear rapidly 912 Å cut: neutral H absorption (H not fully ionized because of lower T) near IR increases: red giants emission Increasing age UV decreases rapidly (more rapid evolution of high M stars) opt/near IR decreases slowly (low M stars die at a slower rate) As a consequence, spectra get redder with age White dwarfs Bruzual & Charlot 1993

24 Spectral evolution: constant SFH SFH(t) = const UV: after some time it does not change anymore: massive stars die but they are continuously replaced by new ones. If SF is prolonged in time UV emission (Attention! We are not considering dust yet) Therefore, if the galaxy emits in the UV, it means that it is forming stars since at least several Myr. UV emission is a tracer of SFR Opt/near IR: increases because evolved stars accumulate (the lifetime of low M ones is larger that the age of the Universe). Slow increase because they are not very luminous. near IR emission is a measure of the total stellar mass (dominated by low M stars remember the IMF) Bruzual & Charlot 1993

25 Spectral evolution: exponential SFH SFH(t) = τ 1 t /τ e SFH decreases, but does not stop (~ intermediate case). UV: increases for some time (massive stars are accumulated), then decreases (massive stars disappear faster than are created), but slower than the burst case Opt/near IR: increases slowly, then decreases (but, as always, slower than the UV and slower than the burst case) Bruzual & Charlot 1993

26 Lyman break (spectral discontinuity) Spectral features It is the end of the Lyman series: 912Å ionization energy. Can be observed in presence of neutral Hydrogen (T eff <~ K). Bruzual & Charlot 1993 Rapid increase at ~10Myr: most massive stars evolve off the MS and ionizing photons are no more produced. Decrease at ~1Gyr: low mass galaxies evolve off the MS and UV photons are generated again by white dwarfs. This rapid evolution is washed out by the dispersion of stellar ages when the SF is prolonged in time. Thus, a large amplitude of the Lyman break indicates that SF occurred recently, and on a short time scale. However, the interpretation of the break is complicated by the sharp absorption possibly produced by B ν (912 A ) = F ν (λ)dλ F ν (λ)dλ Instantaneous burst Constant SFH neutral H in the ISM or IGM. Bruzual & Charlot 1993 τ

27 4000Å break (spectral discontinuity) Spectral features It is caused by several features of absorption of metals (Ca, Mg, etc.) in the stellar atmospheres. It increases with age and metallicity (age indicator, but trends with metallicity). Gorgas et al D 4000 = F ν (λ)dλ F ν (λ)dλ Poggianti & Barbaro 1997 Bruzual & Charlot 1993

28 Balmer break Spectral features It is the end of the Balmer series: 3646Å energy needed for an electron to escape from level 2. It is maximum at age 0.3 1Gyr (spectral type A).. Several spectral indices have been defined to measure the Balmer break. These are sensitive to the abundance of certain elements and to the age of the stellar population. Some indices are measured in terms of an equivalent width (see next lecture) EW = Bruzual & Charlot 1993 Poggianti & Barbaro 1997 line F cont F ν (λ) dλ F cont

29 Equivalent Width Measure of the intensity of the line, independently on its profile, compared to the continuum. EW: width (in λ or ν) of the nearby continuum flux that contains the same power as the line. Large EW pronounced line compared to the continuum. For emission, the EW can be very large if continuum is small. EW = λ 2 λ 1 F c F λ F c dλ EW = λ 2 λ 1 F λ F c F c dλ F c F λ

30 Balmer break Spectral features It is the end of the Balmer series: 3646Å energy needed for an electron to escape from level 2. It is maximum at age 0.3 1Gyr (spectral type A).. Several spectral indices have been defined to measure the Balmer break. These are sensitive to the abundance of certain elements and to the age of the stellar population. Some indices are measured in terms of an equivalent width (see next lecture) EW = Bruzual & Charlot 1993 Poggianti & Barbaro 1997 line F cont F ν (λ) dλ F cont

31 Age metallicity degeneracy Age: spectra become redder as the galaxy ages, since more massive stars evolve off the MS. Metallicity: spectra become redder as the metal content increases, due to a combination of line blanketing effects (spectral light decrease due to unresolved absorption lines) and higher opacities in the stellar atmospheres, which cause the effective temperature to decrease (metals block the escape of photons and surface temperature is lower). The spectra of metal rich galaxies are redder Same age metal poor The spectra of old galaxies are redder Same metallicity young metal rich old

32 Age metallicity degeneracy young & metal rich old & metal poor Chavez & Bertone 2011

33 Galaxy spectra morphological type Ellipticals: redder objects, reproduced by synthetic spectra with long age and short τ: no on going SF, old stellar populations Spirals: bluer objects, reproduced by synthetic spectra with shorter age and longer τ: SF prolonged in time, young + old stellar populations Elliptical Spirals Emission lines produced by gas in the ISM (see next lecture), not by stars Bruzual & Charlot 1993

34 Galaxy spectra morphological type Ellipticals: redder objects, reproduced by synthetic spectra with long age and short τ: no on going SF, old stellar populations Spirals: bluer objects, reproduced by synthetic spectra with shorter age and longer τ: SF prolonged in time, young + old stellar populations Bruzual & Charlot 1993 Spectral Energy Distribution (SED) fitting can give strong indications on the SFH

35 Ellipticals: redder objects, reproduced by synthetic spectra with long age and short τ: no on going SF, old stellar populations Galaxy SFH Spirals: bluer objects, reproduced by synthetic spectra with shorter age and longer τ: SF prolonged in time, young + old stellar populations

36 Galaxy stellar populations - colours Colours are easier to measure than spectra and give an indication of the stellar populations. Elliptical The optical colour is often quantified as B V (B: ~4400Å, V:~5500Å) It is often useful to consider two bands sampling a spectral feature at a given redshift. Spiral Ellipticals: B V > 1 ( 22<MV< 18) Spirals: B V 1 (bulge B V>1) ( 21<MV< 17) Irregulars: B V < 0.8 ( 18<MV< 10) (Attention! We are not considering dust yet)

37 Galaxy stellar populations - colours Colours are easier to measure than spectra and give an indication of the stellar populations. As stellar populations age, the luminosity decreases because of the disappearance of bright main sequence stars and supergiants, and the decrease is more rapid for bluer bands (more massive and hotter stars evolve faster). As a consequence, galaxies become redder with time. Sharp colour increase during the first ~1 2 Gyr massive stars evolve off the main sequence very rapidly. Slow colour increase with time after ~1 2 Gyr old stellar populations are dominated by low mass galaxies, which evolve slowly.

38 elliptical Bimodal colour distribution irregular Baldry et al (ApJ) Bright spiral Faint u - r! u - r!

39 Galaxy stellar populations - colours Colour is a measure of the age of stellar population Colour ~ age! Baldry et al (ApJ) Red sequence Bright, red, old galaxies. Small dispersion: stars have ~ the same age Blue cloud Remember!! Dust reddening (see lecture 5) is complicating things Absolute mag ~ stellar mass! Faint, blue galaxies. Large dispersion: wide distribution of stellar ages

40 Red sequence in clusters Coma cluster Clusters are dominated by elliptical galaxies Renzini et al Elliptical galaxies follow a tight colour magnitude relation red sequence Thickness of the red sequence is due to the different metallicities Used to search for clusters

41 Galaxy integrated spectra Integrated spectra: the entire galaxy is observed simultaneously, integrated emission of the spectra of billions of stars (+ gas and dust) Elliptical Spiral

42 Galaxy integrated spectra Integrated spectra: the entire galaxy is observed simultaneously, integrated emission of the spectra of billions of stars (+ gas and dust) Elliptical Note: this spectrum is slightly shifted in λ (z>0) G band: collection of spectral absorption features, predominantly due to iron, produces a depression around 4100Å

43 Galaxy integrated spectra Integrated spectra: the entire galaxy is observed simultaneously, integrated emission of the spectra of billions of stars (+ gas and dust) Spiral Emission lines due to gas in the ISM heated by OB type stars (see lecture 4) Note: this spectrum is slightly shifted in λ (z>0)

44 Physical parameter estimation from stellar continuum Stellar mass Measured from the overall normalization of the red spectrum (total mass is dominated by low M, low T, long lived stars which emit at near IR λ; need to take into account red giants due to SF) Age Measured from the amplitude of the break (Balmer break and 4000Å break depends on the age) Attention: age metallicity degeneracy! Bruzual & Charlot 1993 SFR (averaged over yr) UV emission indicates on going SF since several Myr

45 Physical parameter estimation from stellar continuum U 360nm B 420nm V 520nm R 650nm I 800nm J 1250nm K 2200nm Multicolour surveys allow us to estimate photometric redshifts and physical parameters by fitting synthetic spectra (SSP + assumed IMF and SFH) Limited by model degeneracy z=1 mag! 3x10 10 <M = 6 x < 9x10 10

46 There s more than stars Galaxies are not made only of stars. They also contain (lots of) gas and dust, and are surrounded by the Intergalactic Medium (IGM) Gas: nebular emission lines HI absorption λ < 912Å Dust: reddens the optical spectrum re emits in the thermal IR IGM: at high z, absorbs part of radiation at λ < 1216Å

47 Key points Galaxy spectra result from the integrated emission of all its components. In this lecture we have considered the contribution of stars: continuum + absorption lines Range: ~ 300 >22000 Å Source of emission/absorption: stars Galaxy stellar spectra depend on the IMF and the SFH Important features: o o o o o o o Lyman break: neutral Hydrogen absorption Balmer break: Balmer limit 4000Å break: caused by metal absorption UV emission indicative of on going SFR Opt/near IR emission indicative of the stellar mass Spectra become redder with age (and with metal abundance) Age/metallicity degeneracy Ellipticals: old stellar populations, red colours; spirals: young + old stellar populations, SF prolonged in time, blue colours

48 Useful bibliography Mo, van den Bosch and White, Galaxy Formation and Evolution, pp Bruzual & Charlot, Spectral evolution of stellar populations using isochrone synthesis, 1993, Astrophysical Journal, vol. 405 Walcher, Groves, Budavari, Dale, Fitting the integrated Spectral Energy Distribution of Galaxies, 2011, Astrophysics and Space Science, vol. 331 (sections 1, 2.1) Madau & Dickinson, Cosmic Star Formation History, 2014, Annual Review of Astronomy and Astrophysics, vol 52 (section 3)