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For this model. a convolved synthetic map for the emission in [OI] 26300 is given in Fig. 8.. | For this model, a convolved synthetic map for the emission in [OI] $\lambda$ 6300 is given in Fig. \ref{Fig_emissmaps_innertrunc}. |
After rescaling the derived jet width to AU. we found an almost constant width in the range of the observed values. | After rescaling the derived jet width to AU, we found an almost constant width in the range of the observed values. |
Note that our model does not provide results farther out than 100 AU due to a small Ro. | Note that our model does not provide results farther out than 100 AU due to a small $R_0$. |
We studied the jet widths derived from synthetic emission maps in different forbidden lines as the full-width half-maximum of the emission. | We studied the jet widths derived from synthetic emission maps in different forbidden lines as the full-width half-maximum of the emission. |
We found that the untruncated model ADO of Vlahakisetal.(2000). cannot account for the small jet widths found in recent optical images taken with HST and AO. | We found that the untruncated model ADO of \citet{VTS00} cannot account for the small jet widths found in recent optical images taken with HST and AO. |
The density normalization is not important for the resulting measured jet width as long as we are far from the critical regime. | The density normalization is not important for the resulting measured jet width as long as we are far from the critical regime. |
We investigated different effects for reducing the deriving jet width: by imposing an outer radius of the launching region of the underlying accreting disk and thus also of the outflow on the observable structure of the jet and by imposing an inner radius of the underlying aceretion disk due to interactions with the stellar magnetosphere. | We investigated different effects for reducing the deriving jet width: by imposing an outer radius of the launching region of the underlying accreting disk and thus also of the outflow on the observable structure of the jet and by imposing an inner radius of the underlying accretion disk due to interactions with the stellar magnetosphere. |
We created synthetic images based on our simulations of truncated disk winds (Stuteetal.2008) as well as new simulations and found that the extracted jet widths in the truncated models decrease for models SCla-lIg. compared to those of the untruncated model ADO. as naively expected. | We created synthetic images based on our simulations of truncated disk winds \citep{STV08} as well as new simulations and found that the extracted jet widths in the truncated models decrease for models SC1a–1g, compared to those of the untruncated model ADO, as naively expected. |
In the present paradigm. Jets are emitted only by the inner part of the disk. | In the present paradigm, jets are emitted only by the inner part of the disk. |
Hence in the other parts the disk can be described by a standard aceretion disk (SAD). in the inner parts by a jet-emitting disk JED). | Hence in the other parts the disk can be described by a standard accretion disk (SAD), in the inner parts by a jet-emitting disk (JED). |
Andersonetal.(2003) showed that one can estimate the launching region às This transition was. constrained observationally with measured jet rotation velocities and using the equation above and radi of the order of 0.I-] AU we used in several theoretical studies as e.g. Combet&Ferreira(2008). | \citet{ALK03} showed that one can estimate the launching region as This transition was constrained observationally with measured jet rotation velocities and using the equation above and radii of the order of 0.1–1 AU are used in several theoretical studies as e.g. \citet{CoF08}. |
. Our results can be used to infer the “real” value of the truncation. radius Ay; in the observed sample of jets and interpret 1t as the transition. radius of the JED to the SAD. assuming the specific model of VOO applies. | Our results can be used to infer the “real” value of the truncation radius $R_{\rm trunc}$ in the observed sample of jets and interpret it as the transition radius of the JED to the SAD, assuming the specific model of V00 applies. |
At the lower boundary in our simulations. the truncation radii are given in Table 1.. | At the lower boundary in our simulations, the truncation radii are given in Table \ref{tbl_models}. |
They vary from 5.375 Ro in model SCla to 0.575 Ro m model SCIg. However. these radii are set atz=6R, (the lower boundary). not in the equatorial plane. | They vary from 5.375 $R_0$ in model SC1a to 0.575 $R_0$ in model SC1g. However, these radii are set at $z = 6\,R_0$ (the lower boundary), not in the equatorial plane. |
Those can be calculated by extrapolating the field line. t.e. with. @uune=aretan(Ryunel-—_6/6)G(PBug: and G. taken from- the analytical solution of VOO. | Those can be calculated by extrapolating the field line, i.e. with $\theta_{\rm trunc} = \arctan ( R_{\rm trunc} |_{z = 6} / 6 )$ and $G$ taken from the analytical solution of V00. |
This gives the following results: | This gives the following results: |
this looks essentially the same as the window function is a eood indication that there are no periodicities in the cala which result from “PW Pic iself. | this looks essentially the same as the window function is a good indication that there are no periodicities in the data which result from TW Pic itself. |
Phe third panel is another "dirty power spectrum of the time series. but this time with the mean value of the data removed. prior to caleulaion of the Fourier transform. | The third panel is another `dirty' power spectrum of the time series, but this time with the mean value of the data removed prior to calculation of the Fourier transform. |
TEus cllectively removes the +irst order window function [rom the data allowing anv wcoals signals to be seen more clearIv. | This effectively removes the `first order' window function from the data allowing any weak signals to be seen more clearly. |
Note that the vertical sca eof this power spectrum is 100 imes greater than in the second panel. | Note that the vertical scale of this power spectrum is 100 times greater than in the second panel. |
Some residual structure ap»ears to be present in this power spectrum. but close inspecjon again reveals that all the peaks are at window function [recuencies. | Some residual structure appears to be present in this power spectrum, but close inspection again reveals that all the peaks are at window function frequencies. |
The bottom panel shows the CLEANed. power spectrum. with the same vertical scale as he hird panel. | The bottom panel shows the ed power spectrum, with the same vertical scale as the third panel. |
The two largest spikes in the CLEANed. [)0WeOr spectrum. near (o ⋅n . ⇀↗≻⋅⋅↱≻↓∪∐∠⊳⋜⊔⋅∢⋅⋜∐∖∖⋎↓⊔∠⇂∪∖∖⊽⇂⊔⊔≼∼⋅jon pequencies and so are unlikely to represent real signals. ( | The two largest spikes in the ed power spectrum, near to $3.5 \times 10^{-4}$ Hz, are at window function frequencies and so are unlikely to represent real signals. ( |
In act. their frequencies correspond. tofadf the orbital peloc of the satellite and | In fact, their frequencies correspond to the orbital period of the satellite and |
of Equation (35). | of Equation (35). |
Fortunately, this transition is greatly simplified by the very simple equation of state implied by the condition Ry=cf. given by with w=—1/3. as we discussed earlier. | Fortunately, this transition is greatly simplified by the very simple equation of state implied by the condition $R_{\rm h}=ct$, given by with $w=-1/3$, as we discussed earlier. |
For a universe with density p and pressure p=wp. the linear relativistic version of Equation (35) is Therefore, for an Aj=cf universe, the dynamical equation tor ó, is We need to emphasize several important features of this equation. | For a universe with density $\rho$ and pressure $p=w\rho$, the linear relativistic version of Equation (35) is Therefore, for an $R_{\rm h}=ct$ universe, the dynamical equation for $\delta_\kappa$ is We need to emphasize several important features of this equation. |
First of all, the active mass in this universe is proportional to e+3p=0. and therefore the gravitational term normally appearing in the standard model is absent (see Equation 35). | First of all, the active mass in this universe is proportional to $\rho+3p=0$, and therefore the gravitational term normally appearing in the standard model is absent (see Equation 35). |
But this does not mean that 6, cannot grow. | But this does not mean that $\delta_\kappa$ cannot grow. |
Instead, because p<0, the (usually dissipative) pressure term in Equation (35) here becomes an agent of growth. | Instead, because $p<0$, the (usually dissipative) pressure term in Equation (35) here becomes an agent of growth. |
Moreover, there is no Jeans length scale. | Moreover, there is no Jeans length scale. |
In its place is the gravitational radius, which we can see most easily by writing Equation (42) in the form where Note, in particular, that both the gravitational radius A, and the fluctuation scale 2 vary with ¢ in exactly the same way, so A, is therefore a constant in time. | In its place is the gravitational radius, which we can see most easily by writing Equation (42) in the form where Note, in particular, that both the gravitational radius $R_{\rm h}$ and the fluctuation scale $\lambda$ vary with $t$ in exactly the same way, so $\Delta_\kappa$ is therefore a constant in time. |
But the growth rate of 6, depends critically on whether 2 is less than or greater than A. | But the growth rate of $\delta_\kappa$ depends critically on whether $\lambda$ is less than or greater than $R_{\rm h}$. |
A simple solution to Equation (43) is the power law where evidently so that Thus, for small fluctuations Cl<< Ry). | A simple solution to Equation (43) is the power law where evidently so that Thus, for small fluctuations $\lambda<<R_{\rm h}$ ), |
1993))). | ). |
The values of the που thus derived for all the regions are listed in Table 2. | The values of the $I_{TRGB}$ thus derived for all the regions are listed in Table 2. |
The distance modulus is eiveu by where fyrgeip is the dereddened Z-baud magnitude of the TRGB. | The distance modulus is given by where $I_{0, TRGB}$ is the dereddened $I$ -band magnitude of the TRGB. |
BC, is the bolometric correction to the I magnitude which depeuds on color as follows: where (V.Dorpep is the dereddened color of the TRGB. | $BC_{I}$ is the bolometric correction to the I magnitude which depends on color as follows: where $(V-I)_{0, TRGB}$ is the dereddened color of the TRGB. |
The bolometric maenitude of the TRGD. ΑΙτης. ds given as a function of mictallicity [Fe/TI] by: Metallicity can be estimated from the (V.£) color at the absolute Z-baud magnitude of M;=3.5 eiven bv Lee.Freediman.&Madore(1993) (see also Saviane (2000))) as follows: We have enploved an iterative procedure iu which an initial guess at the distance is used to estimate the metallicity which is iu turn used to refine the distauce uutil the solution converges. which occurs after only a few iterations. | The bolometric magnitude of the TRGB, $M_{bol,TRGB}$, is given as a function of metallicity [Fe/H] by: Metallicity can be estimated from the $(V-I)$ color at the absolute $I$ -band magnitude of $M_I = -3.5$ given by \citet{lee93}
(see also \citet{sav00}) ) as follows: We have employed an iterative procedure in which an initial guess at the distance is used to estimate the metallicity which is in turn used to refine the distance until the solution converges, which occurs after only a few iterations. |
It is important to note that the regions used for this study. are located in various euvironnienuts iuchiding voung to old stellar populations: thus. the broad ROBs seen in the CALDs are actually à uüxture of itermediate-age to old populations. as well as a range of metallicities. | It is important to note that the regions used for this study are located in various environments including young to old stellar populations; thus, the broad RGBs seen in the CMDs are actually a mixture of intermediate-age to old populations, as well as a range of metallicities. |
Tf we simply use the mean color |WLy, απ of the cutire apparent RGB iu this case. the resulting metallicity will be an underestimate. because there are vounecr populations with bluer color on the blue side of the ROB. | If we simply use the mean color $(V-I)_{0, -3.5}$ ] of the entire apparent RGB in this case, the resulting metallicity will be an underestimate, because there are younger populations with bluer color on the blue side of the RGB. |
For this reason we tried to use the median value of the color of the stars along the RGB to reduce the effect of iuteriiecdiate-age populations. | For this reason we tried to use the median value of the color of the stars along the RGB to reduce the effect of intermediate-age populations. |
As a check of our method. we Lave also derived the mean metallicity using the slope of the RGB as calibrated by Sarajedinietal.(2000).. obtainins verv simular results to those from the median color of the RGB stars. | As a check of our method, we have also derived the mean metallicity using the slope of the RGB as calibrated by \citet{sar00}, obtaining very similar results to those from the median color of the RGB stars. |
The mean metallicities resulting from this procedure are listed in Table 2. | The mean metallicities resulting from this procedure are listed in Table 2. |
The mean metallicity ranges frou Fe/H]| = 0.6 to 0.9 dex. | The mean metallicity ranges from [Fe/H] $\approx$ –0.6 to –0.9 dex. |
Figure 5 displays the mean uctallicity versus the deprojected radial distance of the regions (filled circles). | Figure 5 displays the mean metallicity versus the deprojected radial distance of the regions (filled circles). |
Iu Figure 5 there is clearly a weeative radial eracicut of the metallicity. | In Figure 5 there is clearly a negative radial gradient of the metallicity. |
The mean uctallicity data ave fit by [Fe/H] —05ΕΟΟ,L55|x0.02] ([Fe/H| =COO8,010.061. using the metallicity obtained with the RGD slope uethod) for all the data. where Ry, is given in feriis of spe (10=0.27 kpe is assumed). | The mean metallicity data are fit by [Fe/H] $= -0.05[\pm0.01] R_{dp} - 0.55[\pm0.02]$ ([Fe/H] $= -0.04[\pm0.02] R_{dp} - 0.51[\pm0.06]$, using the metallicity obtained with the RGB slope method) for all the data, where $R_{dp}$ is given in terms of kpc $1^\prime=0.27$ kpc is assumed). |
If we exclude the two innermost regious where the erowding is severe. we obtaiu a fit. |Fe/II] = O.07|4Q.01R,,0.5ΕΟΟ οσο =μεςµηνΟΦΗ using the metallicity obtained with the RGB slope ΕΟΟ)method). similar to values found iu our Galaxys disk using open clusters aud Πο] saute (d|Fe/H|/dR=0.050+0.008. 1) 1979). | If we exclude the two innermost regions where the crowding is severe, we obtain a fit, [Fe/H] $= -0.07[\pm0.01] R_{dp} - 0.48[\pm0.04]$ ([Fe/H] $= -0.08[\pm0.03] R_{dp} - 0.34[\pm0.10]$ using the metallicity obtained with the RGB slope method), similar to values found in our Galaxy's disk using open clusters and field giants $d[Fe/H]/dR = - 0.050 \pm 0.008$ $^{-1}$ ) \citep{jan79}. |
. Iu Fiewre 5 the metallicity of the red eiauts is compare’ with that of ΠΠ regious in MO. | In Figure 5 the metallicity of the red giants is compared with that of HII regions in M33. |
The metallicity of the ΠΠ regions was converted from [O/II| values given in the literature (Iswitter&Aller1981:AIcCall.Rybska.IIuchra1991). using the following relation taken frou King(2000) ΙΟΡο=015ΗΕΠ|0.019. | The metallicity of the HII regions was converted from [O/H] values given in the literature \citep{kwi81, mcc85, vil88, zar94} using the following relation taken from \citet{kin00}: $[O/Fe]=-0.184[Fe/H]+0.019$ . |
The deprojected radius for the TWD regious was calculated as above. | The deprojected radius for the HII regions was calculated as above. |
The metallicity of the ΠΤΙ regions is fit by |Fe/II| =0.12]2:0.02]|0.33[250.07]. which is somewhat steeper hau that for Ru,the field red giants. | The metallicity of the HII regions is fit by [Fe/H] $= -0.12[\pm0.02] R_{dp} + 0.33[\pm0.07]$ , which is somewhat steeper than that for the field red giants. |
The relation for tle WI reeious slows a trend that is similar to that of the field red elauts. but with a larger scatter. | The relation for the HII regions shows a trend that is similar to that of the field red giants, but with a larger scatter. |
It is natural that the ποσα metallicity of the field red giants is lower than that of the WIT regions. because the red giauts are wich older han the IIII regions. | It is natural that the mean metallicity of the field red giants is lower than that of the HII regions, because the red giants are much older than the HII regions. |
Then we derive the distance modulus of each region using the information eiven above. | Then we derive the distance modulus of each region using the information given above. |
Table 2 lists the xumuueters related to the TRGD icthod: the observed f-hand magnitude of the TRGD Uren). the extinction corrected L-band magnitude of the TROB (της). the nean color of the TROB άτυπο]. the mean color measured at M,=3.5 [AVPy.as]. the mean netallicity (|FeTI) of the RGB. the absolute magnitude ofthe TROB (Wyrgep). aud the distance modulus (iMy]. | Table 2 lists the parameters related to the TRGB method; the observed $I$ -band magnitude of the TRGB $I_{TRGB}$ ), the extinction corrected $I$ -band magnitude of the TRGB $I_{0,TRGB}$ ), the mean color of the TRGB $(V-I)_{0,TRGB}$ ], the mean color measured at $M_I = -3.5$ $(V-I)_{0, -3.5}$ ], the mean metallicity ([Fe/H]) of the RGB, the absolute magnitude of the TRGB $M_{I, TRGB}$ ), and the distance modulus $(m-M)_{0}$ ]. |
Figure 6 displavs Ipper versus [Fe/THI] for the tecre regions iu M33. | Figure 6 displays $I_{TRGB}$ versus [Fe/H] for the ten regions in M33. |
The value of τραμ varies little. wit uo obvious net metallicity dependence. | The value of $I_{TRGB}$ varies little, with no obvious net metallicity dependence. |
The mean vali of τομ for the ten regions is Zprge;p=20.58c0.0 showing a remarkably small dispersion. | The mean value of $I_{TRGB}$ for the ten regions is $I_{TRGB} =
20.88\pm0.04$, showing a remarkably small dispersion. |
The average value of the distance moduli for all of the fields is calculated to be (avMoyrere=2L81cx0.0[Grandoni) Uri (systematic). | The average value of the distance moduli for all of the fields is calculated to be $(m-M)_{0,TRGB}=24.81\pm0.04$ $^{+0.15}_{-0.11}$ (systematic). |
The errors for the cistauce modulus are based ou the error budget listed in. Table 3. | The errors for the distance modulus are based on the error budget listed in Table 3. |
The calibration of the TRGD is based. ou Galactic elobular clusters with [Fe/II] = 2.1 to 0.7 dex. vet the derived mean metallicities of the four iuner regions iu our saluple ([Fo/II] = 0.61 to 0.68 dex) are slightly larger than the upper boundary of the calibration range. | The calibration of the TRGB is based on Galactic globular clusters with [Fe/H] = –2.1 to –0.7 dex, yet the derived mean metallicities of the four inner regions in our sample ([Fe/H] = –0.61 to –0.68 dex) are slightly larger than the upper boundary of the calibration range. |
If we use only the six other regions. excluding these four iuner ones, we obtain an average distance modulus of (11) .j(svstematic). | If we use only the six other regions, excluding these four inner ones, we obtain an average distance modulus of $(m-M)_{0,TRGB}=24.83\pm0.06$ $^{+0.15}_{-0.11}$ (systematic). |
un If the theoretical calibration given by Salavis&Cassisi(1998). is adopted CMrpge;p=3.953|LL.IT|OLLALHp? and [ALT]2.39.2:0|OLG87|Do,ss].36.851,"Do,as)?opGas3s| μὴν— the average distance modulus will be GaMorpep=2190940.0 statistical). which is 0.2 mae fainter than that derived usine the enipirical calibration of Lee.Freediman.&Madore(1993). | If the theoretical calibration given by \citet{sal98} is adopted $M_{I,TRGB} = -3.953 + 0.437[M/H] + 0.147[M/H]^2$ and $[M/H]=-39.270+64.687[(V-I)_{0,-3.5}] -36.351[(V-I)_{0,-3.5}]^2
+6.838 [(V-I)_{0,-3.5}]^3 $ ), the average distance modulus will be $(m-M)_{0,TRGB}=24.99\pm0.04$ (statistical), which is 0.2 mag fainter than that derived using the empirical calibration of \citet{lee93}. |
.. However. it las beeu found that the distance obtained with the theoretical calibration is not consistent with other results as shown by Dolphinctal.(2001) in IC 1613. | However, it has been found that the distance obtained with the theoretical calibration is not consistent with other results as shown by \citet{dol01} in IC 1613. |
So we prefer to use the empirical calibration ratler than the theoretical oue iu this study. | So we prefer to use the empirical calibration rather than the theoretical one in this study. |
Finally we adopt GnADytrop=2L81£0.0Lrandou) |irit (svstematic) as the TRGB distance modulus to N23. | Finally we adopt $(m-M)_{0,TRGB}=24.81\pm0.04$ $^{+0.15}_{-0.11}$ (systematic) as the TRGB distance modulus to M33. |
We lave deteriined the distance to M33 using the red chunip as well. | We have determined the distance to M33 using the red clump as well. |
Red clumps are clearly fouud in the CMDz of all the regions. as seen in Fieure 3. | Red clumps are clearly found in the CMDs of all the regions, as seen in Figure 3. |
It is miportaut to note that implicit in the subsequent analysis is the asstuuption that the metallicity determined from the RGB | It is important to note that implicit in the subsequent analysis is the assumption that the metallicity determined from the RGB |
The study of the chemical evolution of nearby spiral galaxies is very important to improve our knowledge about the main ingredients used in chemical evolution models and to test the basic assumptions made for modelling our Galaxy. | The study of the chemical evolution of nearby spiral galaxies is very important to improve our knowledge about the main ingredients used in chemical evolution models and to test the basic assumptions made for modelling our Galaxy. |
M31 and M33 are other spiral members of the Local Group of galaxies and during recent vears many observational studies have been made to investigate the chemical and dynamical properties of these neighbouring systems. | M31 and M33 are other spiral members of the Local Group of galaxies and during recent years many observational studies have been made to investigate the chemical and dynamical properties of these neighbouring systems. |
New surveys (Braun et al. | New surveys (Braun et al. |
2009. Magrini et al. | 2009, Magrini et al. |
2007.2008) contributed to the analysis of different stellar populations and provided more accurate data to constrain the chemical evolution models. | 2007,2008) contributed to the analysis of different stellar populations and provided more accurate data to constrain the chemical evolution models. |
The disks of M31 and M33 have many similarities with the Milky Way disk but some observational constraints like the present day gas distribution can only be explained by assuming different star formation histories for these galaxies. | The disks of M31 and M33 have many similarities with the Milky Way disk but some observational constraints like the present day gas distribution can only be explained by assuming different star formation histories for these galaxies. |
The SFR is one of the most important parameters regulating the chemical evolution of galaxies (Kennicutt 1998. Matteucci 2001. Bossier et al. | The SFR is one of the most important parameters regulating the chemical evolution of galaxies (Kennicutt 1998, Matteucci 2001, Boissier et al. |
2003) together with the initial mass function (IMF). | 2003) together with the initial mass function (IMF). |
Another important mechanism is the “inside-out” disk formation that is very important to reproduce the radial abundance gradients (see Colavitti et al. | Another important mechanism is the "inside-out" disk formation that is very important to reproduce the radial abundance gradients (see Colavitti et al. |
2008 for the most recent paper on the subject). | 2008 for the most recent paper on the subject). |
À faster formation of the inner disk relative to the outer disk was originally proposed. by Matteucci Francoois (1989)and supported in the following years by Boissier Prantzos (1999) and Chiappini et al. ( | A faster formation of the inner disk relative to the outer disk was originally proposed by Matteucci Françoois (1989)and supported in the following years by Boissier Prantzos (1999) and Chiappini et al. ( |
2001). | 2001). |
The chemical evolution of M31 in comparison with that of the Milky Way has been already discussed by Renda et al. ( | The chemical evolution of M31 in comparison with that of the Milky Way has been already discussed by Renda et al. ( |
2005) and Yin et al. ( | 2005) and Yin et al. ( |
2009). | 2009). |
Renda et al. ( | Renda et al. ( |
2005) concluded that while the evolution of the MW and M31 share similar properties. differences in the formation history of these two galaxies are required to explain the observations 1n detail. | 2005) concluded that while the evolution of the MW and M31 share similar properties, differences in the formation history of these two galaxies are required to explain the observations in detail. |
In particular. they found that the observed higher metallicity in the M31 halo can be explained by either (1) a higher halo star formation efficiency. or (11) a larger reservoir of infalling halo gas with a longer halo formation phase. | In particular, they found that the observed higher metallicity in the M31 halo can be explained by either (i) a higher halo star formation efficiency, or (ii) a larger reservoir of infalling halo gas with a longer halo formation phase. |
These two different pictures would lead to (a) a higher [O/Fe] at low metallicities. or (b) younger stellar populations in the M31 halo. respectively. | These two different pictures would lead to (a) a higher [O/Fe] at low metallicities, or (b) younger stellar populations in the M31 halo, respectively. |
Both pictures result in a more massive stellar halo in M31. which suggests a possible correlation between the halo metallicity and its stellar mass. | Both pictures result in a more massive stellar halo in M31, which suggests a possible correlation between the halo metallicity and its stellar mass. |
Yin et al. ( | Yin et al. ( |
2009) concluded that M31 must have been more active in the past than the Milky Way although its current SER 15 lower than in the Milky Way. and that our Galaxy must be a rather quiescent galaxy. atypical of its class (see also Hammer et al. | 2009) concluded that M31 must have been more active in the past than the Milky Way although its current SFR is lower than in the Milky Way, and that our Galaxy must be a rather quiescent galaxy, atypical of its class (see also Hammer et al. |
2007). | 2007). |
They also concluded that the star formation efficiency in M31 must have been higher by a factor of two than in the Galaxy. | They also concluded that the star formation efficiency in M31 must have been higher by a factor of two than in the Galaxy. |
However. by adopting the same SFR as in the Milky Way they failed in reproducing the observed radial profile of the star formatior and of the gas. and suggested that possible dynamical interactions could explain these distributions. | However, by adopting the same SFR as in the Milky Way they failed in reproducing the observed radial profile of the star formation and of the gas, and suggested that possible dynamical interactions could explain these distributions. |
Magrini et al. ( | Magrini et al. ( |
2007) computed the chemical evolution of the disk of M33: they claimed to reproduce the observational features by assuming a continuous almost constant infall of σας. | 2007) computed the chemical evolution of the disk of M33: they claimed to reproduce the observational features by assuming a continuous almost constant infall of gas. |
In this work we present a one-infall chemical evolution model for the Galactic disk based on an updated version of the Chiappini et al. ( | In this work we present a one-infall chemical evolution model for the Galactic disk based on an updated version of the Chiappini et al. ( |
2001) model. | 2001) model. |
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