Francesca by cuiliqing

VIEWS: 10 PAGES: 187

									The Chemical
Evolution of the
Galaxy and Dwarf
Spheroidals of the
Local Group

 Saas-Fee, March 4-10, 2007
 Chemical Evolution of
the Galaxy and dSphs of
    the Local Group
Outline of the lectures
Lecture I: Basic principles of chemical
evolution, main ingredients (star formation
history, nucleosynthesis and gas flows)
Lecture II: Supernova progenitors, basic
equations, analytical and numerical solutions
Lecture III: Detailed chemical evolution
models for the Milky Way
Lecture IV: Model results for the formation
and evolution of the Milky Way
Outline of the lectures
Lecture V: SFR and Hubble sequence, the
effects of the time-delay model on different
galaxies
Lecture VI: Chemical properties and models
of chemical evolution of dSphs
Lecture VII: Comparison between the
evolution of dSphs and the Milky Way.
Interpretation of the alpha/Fe and s- and r-
process el./Fe ratios
How to model galactic
 chemical evolution
Initial conditions (open or closed-box;
chemical composition of the gas)
Birthrate function (SFRxIMF)
Stellar yields (how elements are
produced and restored into the ISM)
Gas flows (infall, outflow, radial flow)
Equations containing all of of this...
     Initial Conditions
a) Start from a gas cloud already present at
t=0 (monolithic model). No flows allowed
(closed-box)
b) Assume that the gas accumulates either
fastly or slowly and the system suffers
outflows (open model)
c) We assume that the gas at t=o is
primordial (no metals)
d) We assume that the gas at t=o is pre-
enriched by Pop III stars
Star Formation History
We define the stellar birthrate function as:
B(m,t) =SFRxIMF
The SFR is the star formation rate (how many
solar masses go into stars per unit time)
The IMF is the initial stellar mass function
describing the distribution of stars as a
function of stellar mass
Parametrization of the
        SFR
The most common parametrization is the
Schmidt (1959) law where the SFR is
proportional to some power (k=2) of the gas
density

Kennicutt (1998) suggested k=1.5 from
studying star forming galaxies, but also a law
depending of the rotation angular speed of
gas
 Other parameters such as gas temperature,
viscosity and magnetic field are usually
ignored
Kennicutt’s (1998) SFR
Kennicutt’s law
SF induced by spiral
   density waves
SFR accounting for
    feedback
The IMF
How to derive the IMF
The current mass distribution of local
MS stars per unit area, n(m), is called
Present Day Mass Function (PDMF)
For stars in the range 0.1-1 Msun, with
lifetimes > the age of the Galaxy, tG, we
can write:
How to derive the IMF
If the IMF is assumed to be constant in
time, we can write:
How to derive the IMF
For stars with lifetimes << tG (m> 2
Msun) we can see only the stars born
after
Therefore, we can write:
How to derive the IMF
If we assume the IMF is constant in time
we can write:



Having assumed that the SFR did not
change during the time interval
corresponding to stellar lifetimes
How to derive the IMF
We cannot apply the previous
approximations to stars in the range 1-2
Msun
Therefore, the IMF is this mass range
will depend on b(tG):
Constraints on the SFH
    from the IMF
In order to obtain a good fit of the two
branches of the IMF in the solar vicinity
one needs to assume (Scalo 1986):
             The IMF
Upper panel:
different IMFs
Lower panel:
normalization of the
multi-slope IMFs to
the Salpeter IMF
Figure from Boissier
& Prantzos (1999)
How to derive the local
         SFR
An IMF should be assumed and then
one should integrate the PDMF in time

Timmes et al. (1995), by adopting the
Miller & Scalo (1979) IMF , obtained:
       The Infall law
The infall rate can simply be constant in
space and time
Or described by an exponential law:
     The outflow law
The rate of gas loss from a galaxy
through a galactic wind can be
expressed as:
  The Yield per stellar
      generation
The yield per stellar generation of a single
chemical element, can be defined as (Tinsley
1980):




Where p_im is the stellar yield and the
instantaneous recycling approximation has
been assumed
Instantaneous Recycling
     Approximation
 The I.R.A. states that all the stars with
 masses < 1 Msun live forever (and this
 is true) but also that the stars with
 masses > 1 Msun die instantaneously
 (and this is not true)
 I.R.A. affects mainly the chemical
 elements produced on long timescales
 (e.g. N and Fe)
 The returned fraction
We define returned fraction the amount of
mass ejected into the ISM by an entire stellar
generation



Instantaneous recycling approximation (IRA)
is assumed , namely stellar lifetimes of stars
with M> 1 Msun are neglected
       Stellar Yields
We call stellar yield the newly produced
and ejected mass of a given chemical
element by a star of mass m

Stellar yields depend upon the mass
and the chemical composition of the
parent star
Primary and Secondary
      elements
We define primary element an element
produced directly from H and He
A typical primary element is carbon or
oxygen which originate from the 3- alpha
reaction
We define secondary element an element
produced starting from metals already
present in the star at birth (e.g. Nitrogen
produced in the CNO cycle)
  Simple Model and
 Secondary Elements
The solution of the Simple model of
chemical evolution for a secondary
element Xs formed from a seed
element Z
Xs is proportional to Z^(2)
Xs/Z goes like Z
         Primary versus
           secondary
Figure from Pettini et al.
(2002)
Small dots are
extragalactic HII regions
Red triangles are
Damped Lyman-alpha
systems (DLA)
Dashed lines mark the
solution of the simple
model for a primary and
a secondary element
        Stellar Yields
Low and intermediate mass stars (0.8-8
Msun): produce He, N, C and heavy s-
process elements. They die as C-O white
dwarfs, when single, and can die as Type Ia
SNe when binaries

Massive stars (M>8-10 Msun): they produce
mainly alpha-elements, some Fe, light s-
process elements and r-process elements
and explode as core-collapse SNe
        Stellar Yields
Yields for Fe in
massive stars
(Woosley & Weaver
1995; Thielemann et
al. 1996; Nomoto et
al. 1997; Rauscher
et al. 2002, Limongi
& Chieffi 2003)
        Stellar Yields
Mg yields from
massive stars
Big differences
among different
studies
Mg yields are too
low to reproduce the
Mg abundances in
stars
          Stellar Yields
Oxygen yields from
massive stars
Different studies agree
on O yields
Oxygen increases
continuously with stellar
mass from 10 to 40
Msun
Not clear what happens
for M>40 Msun
        Stellar Yields
New yield from
Nomoto et al. (2007)
for Oxygen in
massive stars

They are computed
for 4 different
metallicities
          Stellar Yields
Yields of Fe from
massive stars from
Nomoto et al. (2007)

The yields are
computed for 4
different metallicities
thermo
      Supernova taxonomy
                                                    core
nuclear                    H                        collapse

           I                                   II
          Si                       light curve
                                                         line
                                                      profile
                      IIb
  Ia
               He              phase       phase



                               IIL IIP                 


          Ic    Ib                     luminosity     IIn
 energy
                                       faint II
  Hypernovae    = GRBs ?
Basic SN types
   max.          +10 months

          Ia




          Ib


          II
SN type             Ia              II+Ib/c
Progenitor      WD in binary     single or binary
                 system (M <     massive star (>
                    8MO)              8MO)
Mechanism      thermo-nuclear     core-collapse


total energy     ~1051 ergs        ~1053 ergs
Remnant            none          neutron star (or
                                       BH)
Ejecta            1.4 MO            1 - 30 MO
composition         Fe          O, Mg, Si, Ne, Ca
age            0.03 – 10 Gyr        < 30 Myr
  SFR and galaxy type




Kennicutt (1998)
Type Ia SN progenitors
Single-degenerate scenario Whelan &
Iben 1974): a binary system with a C-O
white dwarf plus a normal star. When
the star becomes RG it starts accreting
mass onto the WD
When the WD reaches the
Chandrasekhar mass it explodes by C-
deflagration as Type Ia supernova
Type Ia SN progenitors
Double-Degenerate scenario (Iben &
Tutukov, 1984): two C-O WDs merge after
loosing angular momentum due to
gravitational wave radiation
When the two WDs of 0.7 Msun merge, the
Chandrasekhar mass is reached and C-
deflagration occurs
The nucleosynthesis is the same in the two
scenarios
Single-Degenerate
     scenario
DD-scenario
  The clocks for the
 explosions of SNe Ia
Single-Degenerate model: the clock to the
explosion is given by the lifetime of the
secondary star, m2. The minimum time for
the appearence of the first Type Ia SN is
tSNIa= 30Myr (the lifetime of a 8 Msun star)
Double-Degenerate model: the clock is given
by the lifetime of the secondary plus the
gravitational time-delay. tSNIa= 35 Myr +
Delta_grav= 40 Myr
The maximum timescale is 10 Gyr in the SD
 and several Hubble times in the DD
        Type Ia SN
      nucleosynthesis
A Chandrasekhar
mass (1.44 Msun)
explodes by C-
deflagration
C-deflagration
produces 0.6 Msun
of Fe plus traces of
other elements from
C to Si
Tycho SNR (type Ia)




color code: red .30-.95 keV, green .95-2.65 keV, blue 2.65-7.00 keV


Chandra X-ray images
           Type II SNe
Type II SNe arise from
the core collapse of
massive stars (M=8-40
Msun) and produce
mainly alpha-elements
(O, Mg, Si, Ca...) and
some Fe
Stars more massive can
end up as Type Ib/c
SNe
       Summary of
     Nucleosynthesis
During the Big Bang light elements are
formed,
Spallation process in the ISM produces
6Li, Be and B
Supernovae II produce alpha-elements
(O, Ne, Mg, S, S, Ca), some Fe, light s-
and r-process elements
      Summary of
    Nucleosynthesis
Type Ia SNe produce mainly Fe and Fe-
peak elements plus some traces of
elements from C to Si
Low and intermediate mass stars
produce
Deuterium is only destroyed to produce
3He which is also mainly destroyed
   The Simple model
The Simple Model of galactic chemical
evolution
One-zone, closed -box model (no infall
or outflow)
IMF constant in time
Instantaneous recycling approximation
Instantaneous mixing approximation
Solution of the Simple
        Model
If we assume that Xi is the abundance
by mass of an element i, we have:




where
Simple model with
     outflow
Simple model with infall
Abundance ratios and
   Simple Models
Under the assumption of the I.R.A.
it is always true that the ratio of two
abundances is equal to the ratio of the
two corresponding yields:
 Models with no I.R.A.
When the I.R.A. Is relaxed then is NOT
more true that the ratio between the
abundances of two different elements is
equal to the ratio of the corresponding
yield!!
Basic Equations
Definitions of variables
dGi/dt is the rate of time variation of the
gas fraction in the form of an element i
Xi(t) is the abundance by mass of a
given element i
Qmi is a term containing all the
information about stellar evolution and
nucleosynthesis
Definition of variables
A =0.05-0.09 is the fraction in the IMF of
binary systems of that particular type to give
rise to Type Ia SNe. B=1-A
Tau_m is the lifetime of a star of mass m
f(mu) is the distribution function of the mass
ratio in binary systems
A(t) and W(t) are the accretion and outflow
rate, respectively
The Milky Way
The Milky Way
 The formation of the
      Milky Way
Eggen, Lynden-Bell & Sandage (1962)
suggested a rapid collapse lasting 300
Myr for the formation of the Galaxy
Searle & Zinn (1978) proposed a central
collapse but also that the outer halo
formed by merging of large fragments
taking place over a timescale > 1Gyr
Different approaches in
   modelling the MW
Serial approach: halo, thick and thin
disk form as a continuous process (e.g.
Matteucci & Francois 1989)
Parallel approach: the different galactic
component evolve at different rates but
they are inter-connected (e. G. Pardi,
Ferrini & Matteucci 1995)
Different approaches in
   modelling the MW
Two-infall approach: halo and disk form
out of two different infall episodes (e.g.
Chiappini, Matteucci & Gratton 1997;
Alibes, Labay & Canal 2001)
Stochastic approach: mixing not
efficient especially in the early halo
phases (e.g. Tsujimoto et al. 1999;
Argast et al. 2000; Oey 2000)
   A scenario for the
formation of the Galaxy
The two-infall model
of Chiappini,
Matteucci & Gratton
(1997) predicts two
main episodes of
gas accretion
During the first one
the halo and bulge
formed, the second
gave rise to the disk
 The two-infall model
The two-infall model has been adopted
also in other studies such as Chang et
al.(1999) and Alibes et al. (2001)
In particular, Chang et al. applied the
two-infall scheme to the thick and thin
disk
Alibes et al. adopted the same scheme
as Chiappini et al. (1997)
Gas Infall at the present
           time
Another scenario
         The creation of the
         Milky way
         Hera, flowed when
         she realized she
         had been giving milk
         to Heracles and
         thrust him away her
         breast
  Recipes for the two-
      infall model
SFR- Kennicutt’s law with a dependence on
the surface gas density (exponent k=1.5) plus
a dependence on the total surface mass
density (feedback). Threshold of 7 solar
masses per pc squared
IMF, Scalo (1986) normalized over a mass
range of 0.1-100 solar masses
Exponential infall law with different
timescales for inner halo (1-2 Gyr) and disk
(inside-out formation with 7 Gyr at the S.N.)
Recipes for the model
Type Ia SNe- Single degenerate model
(WD+RG or MS star), recipe from Greggio &
Renzini (1983) and Matteucci & Recchi
(2001)
Minimum time for explosion 35 Myr (lifetime
of a 8 solar masses star), confirmed by recent
findings (Mannucci et al. 2005, 2006)
Time for restoring the bulk of Fe in the S.N. is
1 Gyr (depends on the assumed SFR)
        Solar Vicinity
We study first the solar vicinity, namely
the local ring at 8 kpc from the galactic
center
Then we study the properties of the
entire disk from 4 to 22 Kpc
Stellar Lifetimes
The star formation rate
  (threshold effects)
   Stellar abundances
[X/Fe]= log(X/Fe)_star-log(X/Fe)_sun is the
abundance of an element X relative to iron
and to the Sun
The most recent accurate solar abundances
are from Asplund et al. (2005)
Previous abundances from Anders &
Grevesse (1989) and Grevesse & Sauval
(1998)
The main difference is in the O abundance,
now lower
   Predicted SN rates
Type II SN rate
(blue) follows the
SFR
Type Ia SN rate
(red) increases
smoothly (small
peak at 1 Gyr)
    Time-delay model
Blue line= only Type
II SNe to produce
Fe
Red line= only Type
Ia SNe to produce
Fe
Black line: Type II
SNe produce 1/3 of
Fe and Type Ia SNe
produce 2/3 of Fe
Specific prediction by
 the two-infall model
The adoption of a
threshold in the gas
density for the SFR
creates a gap in the
SFR
This gap occurs
between the halo-thick
disk and the thin-disk
phase
It is observed in the
data
G-dwarf distribution
 (Chiappini et al.)
Different timescales for
     disk formation
         G-dwarf
distribution(Alibes et al.)
  G-dwarf distribution
 Chiappini et al. (1997) , Alibes et al.
(2001) and Kotoneva et al. (2002)
concluded that a good fit to the G-dwarf
metallicity distribution can be obtained
only with a time scale of disk formation
at the solar distance of 7-8 Gyr
Evolution of the element
      abundances
Chiappini et al. follow the evolution in space
and time of 35 chemical species (H, D, He, Li,
C, N, O, Ne, Mg, Si, S, Ca, Ti, K, Fe, Mn, Cr,
Ni, Co, Sc, Zn, Cu, Ba, Eu, Y, La, Sr plus
other isotopes)
They solve a system of 35 equations where
SFR, IMF, nucleosynthesis and gas accretion
are taken into account
Yields from massive stars WW95, from low-
intermediate stars van den Hoeck+
Groenewegen 1997, from Type Ia SNe
Iwamoto et al. 1999
Results from Francois et
        al. 2004
Results from Francois et
        al. 2004
Results from Francois et
        al. 2004
Corrected Yields
Corrected Yields
Corrected Yields
   Suggestions for the
        Yields
Yields from Woosley
& Weaver 1995
(WW95), Iwamoto et
al. (1999)
Major corrections for
Fe-peak elements
O, Fe, Si and Ca are
ok. Mg should be
increased
Inhomogeneous Model
Argast et al. (2000) computed 3-D
hydrodynamical calculations following
the evolution of SN remnants
No mixing was assumed for [Fe/H] > -
3.0 dex, complete mixing for [Fe/H]> -
2.0 dex
They predicted a too large spread for
[Mg/Fe] and [O/Fe] vs. [Fe/H]
Results from Alibes et
          al.
Alibes et al. (2001)
adopted the two-
infall model
Metallicity-
dependent yields
from WW95 and
Van den Hoeck &
Groenewegen
(1997)
Results from Chiappini
         et al.
Evolution of Carbon and
Nitrogen as predicted
by the two-infall model
of Chiappini, Matteucci
& Gratton (1997)
The green line in the N
plot is an euristic model
with primary N from
massive stars
Last data on Nitrogen
From Ballero et al.
(2005)
It shows new data (filled
circles and triangles) at
low metallicity
endorsing the
suggestion that N
should be primary in
massive star
Stellar rotation can
produce such N
(Meynet & Maeder
2002)
 Last data on N and C
Primary nitrogen from
rotating very metal poor
massive stars
Models from Chiappini
et al. (2006) (dashed
lines)
Large squares from
Israelian et al. 04;
asterisks from Spite et
al. 05; pentagons from
Nissen 04
      s- and r-process
          elements
Data from Francois
et al.(2006) with
UVES on VLT
Models Cescutti et
al. (2006): red line,
best model, with
Ba_s from 1-3 solar
masses (Busso et
al.01) and Ba_r from
10-30 solar masses
     Old Prescriptions
Travaglio et
al.(1999) assumed
Ba_r from 8-10 solar
masses
The new data show
a source of Ba_r
from more massive
stars is required
      s- and r- process
          elements
Data from Francois
et al. (2006)
Models from
Cescutti et al.
(2006): red line, best
model with Eu only
r-process from 10-
30 solar masses
      s- and r- process
          elements
Lanthanum- Data from:
Francois et al. (2006)
(filled red squares),
Cowan & al.(2005)
(blue hexagons), Venn
et al.(2004) (blue
triangles), Pompeia et
al.(2003) (green
hexagons)
Models from Cescutti,
Matteucci, Francois &
Chiappini (2006): same
origin as Ba
 Abundance Gradients
The abundances of heavy elements
decrease with galactocentric distance
in the disk
Gradients of different elements are slightly
different (depend on their nucleosynthesis
and timescales of production)
Gradients are measured from HII regions,
PNe, B stars, open clusters and Cepheids
How does the gradient
       form?
If one assumes the disk to form inside-
out, namely that first collapses the gas
which forms the inner parts and then the
gas which forms the outer parts
Namely, if one assumes a timescale for
the formation of the disk increasing with
galactocentric distance, the gradients
are well reproduced
 Abundance gradients
Predicted and observed
abundance gradients
from Chiappini,
Matteucci & Romano
(2001)
Data from HII regions,
PNe and B stars, red
dot is the Sun
The gradients steepen
with time (from blue to
red)
 Abundance gradients
Predictions from
Boissier & Prantzos
(1999), no threshold
density in the SF
They predict the
gradient to flatten in
time
The difference is
due to the effect of
the threshold
 Abundance Gradients
New data on Cepheids
from Andrievsky &
al.(02,04) (open blue
circles)
Red triangles-OB stars
from Daflon & Cunha
(2004)
Blue filled hexagons,
Cepheids from Yong et
al.(2006), blue open
triangles from Young et
al. 05, cian data from
Carraro et al.(2004)
Different halo densities
Only Cepheids data
from Andrievsky
Blue dot-dashed
line: model with halo
density decreasing
outwards
Red continuous line
(BM):model with
halo constant
density
 Abundance Gradients
Blue filled hexagons
from open Cepheids
(Yong et al. 2006)
Cian data from open
clusters (Carraro et al.
2004), open triangles
are open clusters
Black data from
Cepheids (Andrievsky
et al., 2002,04)
Dashed lines=prediction
for 4.5 Gyr ago
Abundance Gradients

Blue filled hexagons
from Andrievsky &
al.(02,04)
Red squares are the
average values
For Barium there
are not yet enough
data to compare
    The Galactic Bulge
A model for the Bulge
(green line, Ballero et
al. 2006)
Yields from Francois et
al. (04), SF efficiency of
20 Gyr^(-1), timescale
of accretion 0.1 Gyr
Data from Zoccali et al.
06, Fulbright et al. 06,
Origlia &Rich (04, 05)
    The Galactic Bulge
Model (red, Ballero et
al. 2006)
Predicts large Mg to Fe
for a large Fe interval
Turning point at larger
than solar Fe. Mg flatter
than O
Data from Zoccali et al.
06; Fulbright et al. 06,
Origlia & Rich (04, 05)
   The Galactic Bulge
Metallicity
Distribution of Bulge
stars, data from
Zoccali et al. (2003)
and Fulbright et al.
(2006) (dot-dashed)
Models from Ballero
et al. 06, with
different SF eff.
   The Galactic Bulge
Models with different
IMF
The best IMF for the
Bulge is flatter than
in the S.N: and
flatter than Salpeter
Best IMF: x=0.95 for
M> 1 solar mass
and x=0.33 below
Bulge vs. Thick and Thin
       Disk Stars
Zoccali et al. (2006)
compared new high
resolution data for the
Bulge (green dots and
red crosses) with data
for thick disk (yellow
triangles) and thin disk
(blue crosses)
The Bulge stars are
systematically more
overabundant in O
  Other Bulge Models
Molla, Ferrini &
Gozzi (2000): the
Bulge formed by
collapse but with a
more prolonged star
formation history
They failed in
reproduding [Mg/Fe]
   Other Bulge Models
Immeli et al. (2004)
computed dynamical
simulations for the
formation of the Bulge
They studied the
efficiency of energy
dissipation and different
SF histories
Model B assumes an
early and fast SFR
Comparison with data
Comparison between
the B, D and F models
of Immeli et al. (2004)
with data from Zoccali
et al. (2006)
The best model predicts
a very fast Bulge
formation
However, Immeli’s
models have a fixed
delay for Type Ia SNe
   Conclusions on the
         Bulge
The best model for the Bulge suggests that it
formed by means of a strong starburst
The efficiency of SF was 20 times higher than
in the rest of the Galaxy
The IMF was very flat, as it is suggested for
starbursts
The timescale for the Bulge formation was 0.1
Gyr and not longer than 0.5 Gyr
   Conclusions on the
       Milky Way
The Disk at the solar ring formed on a time
scale not shorter than 7 Gyr
The whole Disk formed inside-out with
timescales of the order of 2 Gyr in the inner
regions and 10 Gyr in the outer regions
The inner halo formed on a timescale not
longer than 2 Gyr
Gradients from Cepheids are flatter at large
Rg than gradients from other indicators
Dwarf Spheroidals of the
      Local Group
    SF and Hubble
Sequence from Sandage
SF and HS from
  Kennicutt
Models for the Hubble
     Sequence
Type Ia SN rate in
    galaxies
Timescales for Type Ia
   SNe enrichment
The typical timescale for the Type Ia SN
enrichment is the maximum in the Type
Ia SN rate (Matteucci & Recchi 2001)
It depends on the star formation history
of a specific galaxy, IMF and stellar
lifetimes
Typical timescales for
         SNIa
In ellipticals and bulges the timescale
for the maximum enrichment from Type
Ia SNe is 0.3-0.5 Gyr
In the solar vicinity there is a first peak
at 1 Gyr, then it decreases slightly (gap
in the SF) and increases again till 3 Gyr
In irregulars the peak is for a time > 4
Gyr
Time-delay model in
 different galaxies
Interpretation of time-
     delay model
Galaxies with intense SF (ellipticals and
bulges) show overabundance of alpha-
elements for a large [Fe/H] range
Galaxies with slow SF (irregulars) show
instead low [alpha/Fe] ratios at low
[Fe/H]
The SFR determines the shape of the
[alpha/Fe] vs. [Fe/H] relations
     Identifying high-z
          objects
Lyman-break galaxy
cB58, data from
Pettini et al. 2002
The model
predictions are for
an elliptical galaxy
of 10^(10) Msun
(Matteucci & Pipino
2002)
 Dating high-z objects
The Lyman-break
galaxy cB58
Predicted
abundance ratios
versus redshift
The estimated age
is 35 Myr
Conclusions on high-z
      objects
Comparison between data and
abundance ratios of high-z objects
suggests:
DLA are probably dwarf irregulars or at
most external parts of disks
Lyman-break galaxies are probably
small ellipticals in the phase of galactic
wind
  How do dSphs form?
CDM models for galaxy formation predict
dSph systems (10^7 Msun) to be the first to
form stars (all stars should form < 1Gyr)
Then heating and gas loss due to
reionization must have halted soon SF
Observationally, all dSph satellites of the MW
contain old stars indistinguishable from those
of Galactic globular clusters and they have
experienced SF for long periods (>2 Gyr,
Grebel & Gallagher, 04)
Chemical Evolution of
 Dwarf Spheroidals
Lanfranchi & Matteucci (2003, 2004)
proposed a model which assumes the SF as
derived by the CMDs
Initial baryonic masses 5x10^(8)Msun
A strong galactic wind occurs when the gas
thermal energy equates the gas potential
energy. DM ten times LM but diffuse (M/L
today of the order of 100)
The wind rate is assumed to be several times
the SFR
Standard Model of LM03
LM03 computed a standard models for dwarf
spheroidals
They assumed 1 long star formation episode
(8 Gyr), a low star formation efficiency <1
Gyr^(-1)
They assumed that galactic winds are
triggered by SN explosions at rates > 5 times
the SFR . The final mass is 10^(7)Msun
The IMF is that of Salpeter (1955)
       Galactic winds
LM03 included the energetics from SNe and
stellar winds to study the occurrence of
galactic winds, the condition for the wind
being:


Dark matter halos 10 times more massive
than the initial luminous mass (5x10^(8)
Msun) but not very concentrated (see later)
The binding energy of
         gas
The binding energy of
         gas
Binding energy of gas
S is the ratio
between the
effective radius of
the galaxy and the
radius of the dark
matter core
 We assume S=0.10
in dSphs
DM in Dwarf Spheroidals
Mass to light ratios vs.
Galaxy absolute V
magnitude (Gilmore et
al. 2006)
The solid curve shows
the relation expected if
all the dSphs contain
about 4x10^(7) Msun of
DM interior to their
stellar distributions
No galaxy has a DM
halo < 5x10^(7)Msun
          DM in dSphs
Mass to light ratios
in dSphs from
Mateo et al. (1998)
In the bottom panel
the visual absolute
magnitude has been
corrected for stellar
evolution effects
The Sgr point is an
upper limit
       Galactic Winds
The energy feedback from SNe and stellar
winds in LM03 is:
SNe II inject 0.03 Eo (Eo is the initial blast
wave energy of 10^(51) erg )
SNe Ia inject Eo since they explode when the
gas is already hot and with low density
(Recchi et al. 2001)
Stellar winds inject 0.03 Ew (Ew is 10^(49)
erg)
 Gas Infall and Galaxy
      Formation
LM03 assumed that each galaxy forms
by infall of gas of primordial chemical
composition

The formation occurs on a short
timescale of 0.5 Gyr
Standard Model of LM03
Standard Model: SF
lasts for 8 Gyr,
strong wind
removes all the gas
Different SF eff.
and wind eff. are
tested, from 0.005
to 5 Gyr^(-1) for SF
and from (6 to 15)
xSF for the winds
  Abundance patterns
It is evident that the [alpha/Fe] ratios in
dSphs show a steeper decline with
[Fe/H] than in the stars in the Milky Way
This is the effect of the time-delay
model, namely of a low SF efficiency
coupled with a strong galactic wind
After the wind SF continues for a while
   Individual galaxies
Then LM03,04 computed the evolution
of 6 dSphs: Carina, Sextan, Draco,
Sculptor, Sagittarius and Ursa Minor

They assumed the SF histories as
measured by the Color-Magnitude
diagrams (Mateo, 1998;Dolphin 2002;
Hernandez et al. 2000; Rizzi et al. 2003)
Star Formation Histories
        in LM03
SF histories of dSphs
 (Mateo et al. 1998)
Individual galaxies
   Dwarf Spheroidals :
         Carina
Model Lanfranchi &
Matteucci (04,06)
SF history from Rizzi
et al. 03. Four bursts
of 2 Gyr, SF
efficiency 0.15
Gyr^(-1) < 1- 2
Gyr^(-1) (S.N.),
wind=7xSFR
Salpeter IMF
  Predicted C and N in
         Carina
Predicted evolution
of C and N for
Carina’s best model
The continuous line
is for secondary N in
massive stars
The dashed line
assumes primary N
from massive stars
Metallicity distribution
       in Carina
Data from Koch et
al. (2005)
Best model from
Lanfranchi & al.
(2006)
This model well
reproduces also the
[alpha/Fe] ratios in
Carina
   Dwarf Spheroidals:
         Draco
Model and data for
Draco
SF history, 1 burst of
4 Gyr, SF efficiency
of 0.03 Gyr^(-1)
Wind=6xSFR
Salpeter IMF
   Draco’s metallicity
      distribution
Predicted metallicity
distribution for
Draco compared
with the predicted
metallicity
distribution for the
Solar Vicinity
   Dwarf Spheroidals:
       Sextans
Best Model: 1 burst
of 8 Gyr
SF efficiency 0.08
Gyr^(-1)
Wind=9xSFR
Salpeter IMF
  Sextans: metallicity
      distribution
Predicted metallicity
distribution for
Sextans by LM04

The predicted G-
dwarf metallicity
distribution for Solar
Vicinity stars is
shown for
comparison
Dwarf Spheroidals: Ursa
        Minor
Best Model: 1 burst
of 3 Gyr
SF efficiency 0.2
Gyr^(-1)
Wind=10xSFR
Salpeter IMF
Ursa Minor’s metallicity
     distribution
Predicted metallicity
distribution for Ursa
Minor by LM04

The predicted G-
dwarf metallicity
distribution for the
solar vicinity is
shown for
comparison
   Dwarf spheroidals:
      Sagittarius
Best model:one long
episode of SF of
duration 13 Gyr
(Dolphin et al 2002)
SF eff. Like the S.N.,
but very strong wind
9XSFR
Metallicity distribution
    in Sagittarius
Predicted metallicity
distribution by LM04 for
Sagittarius: continuous
line (Salpeter IMF),
dashed line (Scalo IMF)
The predicted G-dwarf
metallicity distribution
for the solar vicinity is
shown by the dotted
line
   Dwarf Spheroidals:
       Sculptor
Model and data for
Sculptor
SF efficiency 0.05-
0.5 Gyr^(-1), wind
7 XSFR
One long SF episide
lasting 7 Gyr
Salpeter IMF
 Sculptor’s metallicity
     distribution
Predicted metallicity
distribution in
Sculptor (LM04)

The predicted G-
dwarf metallicity
distribution for the
solar vicinity is
shown for
comparison
    s- and r- process
   elements in dSphs
Lanfranchi et al. 2006 adopted the
nucleosynthesis prescriptions for the s-
and r- process elements as in the S.N.
They calculated the evolution of the
[s/Fe] and [r/Fe] ratios in dSphs
They predicted that s-process elements,
which are produced on long timescales
are higher for the same [Fe/H] in dSphs
Model and data for
     Carina
Model and data for
      Draco
Model and data for
    Sextans
Model and data for
    Sculptor
Model and data for
   Sagittarius
Sagittarius: more data
Best model is
continuous line. Dotted
lines are different SF
efficiencies
Dashed line is the best
model with no wind
The strong wind
compensate the high
SF efficiency
Data from Bonifacio et
al. 02,04 & Monaco et
al. 05 (open squares)
C and N in Sagittarius:
     predictions
Other Models for dSphs
Carigi, Hernandez & Gilmore (2002)
computed models for 4 dSphs by assuming
SF histories derived by Hernandez et al.
(2000)
They assumed gas infall and computed the
gas thermal energy to study galactic winds
They assumed a Kroupa et al.(1993) IMF
   Carigi et al. (2002)
They assumed only a sudden wind
which devoids the galaxy from gas
instantaneously

They predicted a too high metallicity for
dSphs and not the correct slope for
[alpha/ Fe] ratios
Carigi et al’s predictions
     for Ursa Minor
     Model of Ikuta &
     Arimoto (2002)
They adopted a closed model (no infall,
no outflow)
They suggested a very low SFR such
as that of LM03, 04
They had to invoke external
mechanisms to stop the SF
They assumed different IMFs
Ikuta & Arimoto (2002)
Model of Fenner et al.
        2006
Very similar to the model of LM03, 04 with
galactic winds for Sculptor
They suggest 0.05 Gyr^(-1) as SF efficiency
Their galactic wind is not as strong as the
winds of LM03, 04
They conclude that chemical evolution in
dSphs is inconsistent with SF being truncated
after reionization epoch (z =8)
 Comparison between
   dSphs and MW
Blue line and blue data
refer to Sculptor
Red line and red data
refer to the Milky Way
The effect of the time-
delay model is to shift
towards left the model
for Sculptor with a lower
SF efficiency than in the
MW
Comparison dSph and
        MW
Eu/Fe in Sculptor
and the MW
Model and data for
Sculptor are in blue
Model and data for
the MW are in red
Conclusions on dSphs
By comparing the [alpha/Fe] ratios in
the MW and dSphs one concludes that
they had different SF histories
The [alpha/Fe] ratios in dSphs are
always lower than in the MW at the
same [Fe/H], as a consequence of the
time delay model and strong galactic
wind
Conclusions on dSphs
Very good agreement both for [alpha/Fe] and
[s/Fe] and [r/Fe] ratios is obtained if a less
efficient SF than the S.N. one and a strong
wind are adopted
It is unlikely that the dSphs are the building
blocks of the MW
Interactions between the MW and its satellites
are not excluded but they must have occurred
after the bulk of stars of dSphs was formed
Other spirals
 Results for M101
(Chiappini et al. 03)
Results for M101
Properties of spirals
 (Boissier et al. 01)
Conclusions on Spirals
N132D in LMC
                         [SII] red, [OIII] green, [OI] blue
oxygen rich SN remnant
SN 1998dh     How to search
            Compare images taken at different epochs

            • few days < time interval < 1-2 month

               • 14 < limiting magnitude < 24

                • 0.01 < target redshift   <1

              • 5 arcmin < field of view < 1 deg

                    • B-V < band < R-I
         SN search

target       reference       difference



         -               =
                              SN 2000fc
                              type Ia
                              z = 0.42
                              V=22.4
                              IAUC7537
SN distribution in galactic coordinates
           SN rate with redshift
      Madau, Della Valle & Panagia 1998 On the evolution of the cosmic
         supernova rate
      Sadat et al. 1998 A&A 331, L69 Cosmic star formation and Type Ia/II
          supernova rates at high Z
      Yungelson & Livio 2000 ApJ 528, 108 Supernova Rates: A Cosmic
          History
                                        The History of the Cosmic Supernova
      Kobayashi et al. 2000 ApJ 539, 26 0.3Gy
          Rate Derived from the Evolution of the Host Galaxies
      Sullivan et al. 2000 MNRAS 319, 549 A strategy for finding
                                            1Gy
           gravitationally lensed distant supernovae
      Dahlèn & Fransson 1999τ =3Gy 349 Rates and redshift distributions
                              A&A 350,
          of high-z supernovae
      Calura & Matteucci 2003 ApJ 596, 734
Madau, Della Valle & Panagia 1998 MNRAS 297, L17
Astrophysics: massive star
        evolution
                           GRBs
       NS BH




                     Zampieri et al.
                     (2003) MNRAS
                     338, 711

								
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