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Because of their low radiative trausition rates. these ecnussion lines are quite weak: the metal fine-structure lues are expected to be considerably stronger. even at T=200 Ik. once the metallicity rises above 1A | Because of their low radiative transition rates, these emission lines are quite weak; the metal fine-structure lines are expected to be considerably stronger, even at $T = 200$ K, once the metallicity rises above $10^{-4}~Z_{\odot}$. |
Using individual netallicities for C. ο. Si. and Fe. we found an nmuportaut distinction from hne sinele-uctallicity case. | Using individual metallicities for C, O, Si, and Fe, we found an important distinction from the single-metallicity case. |
These results sugeest that. for lMassive stars chanced iu a-process elements (O. Si aud Fe). the critical metallicity for fragmentation is much ligher than xevious works have suggested. | These results suggest that, for massive stars enhanced in $\alpha$ -process elements (O, Si and Fe), the critical metallicity for fragmentation is much higher than previous works have suggested. |
Iaüeher uetallicitics are also required if the metalenriched gas has lower deusities han ney. | Higher metallicities are also required if the metal-enriched gas has lower densities than $n_{\rm cr}$. |
Redshitted fine-structure lines frou metal-euriched xuxnordial eas at 200 I& could be observable from high-redshift halos. | Redshifted fine-structure lines from metal-enriched primordial gas at 200 K could be observable from high-redshift halos. |
Iu addition to the weak rotational lines of ((28 gu aud 17 jan) aud 1157.71 jin. the mos inportaut coolants are likely to be 663.18 µια, 33.8 pan. aud ((25.99 jan and 35.35 pau). | In addition to the weak rotational lines of (28 $\mu$ m and 17 $\mu$ m) and 157.74 $\mu$ m, the most important coolants are likely to be 63.18 $\mu$ m, 34.8 $\mu$ m, and (25.99 $\mu$ m and 35.35 $\mu$ m). |
With redshifting. these lines would appear in the FIR sub-uuu. aud even the band. if (1|:)z1020. | With redshifting, these lines would appear in the FIR sub-mm, and even the mm-band, if $(1+z) \approx 10-20$. |
We are erateful to Aparna Veukatesan. JasonΕΕ Jason Cilenuu. and Dau Lester for useful discussions regarding this project.aud Phil Maloney. for help with the nuuercal techniques. | We are grateful to Aparna Venkatesan, JasonTumlinson, Jason Glenn, and Dan Lester for useful discussions regarding this project,and Phil Maloney for help with the numerical techniques. |
This research was supported bv astrophysics theory erants from NASA (NACH5-7262) aud NSF (AST02-06012). | This research was supported by astrophysics theory grants from NASA (NAG5-7262) and NSF (AST02-06042). |
resonances. approaching. in the limit of large clusters. bulk optical properties. | resonances, approaching, in the limit of large clusters, bulk optical properties. |
This would bring the absorption spectra of such improved mixtures of PAHs towards better agreement with observations. | This would bring the absorption spectra of such improved mixtures of PAHs towards better agreement with observations. |
This will need to be confirmed by future calculations. | This will need to be confirmed by future calculations. |
Mixtures of PAHs can accurately account for the UV bump and non-linear far-UV rise in extinction, and the PAH charge state is linked to the relative intensity of the two features. | Mixtures of PAHs can accurately account for the UV bump and linear UV rise in extinction, and the PAH charge state is linked to the relative intensity of the two features. |
This provides a physical explanation and quantitative relationship of the 2-dimensional variations in ISECs described by FMO7. | This provides a physical explanation and quantitative relationship of the dimensional variations in ISECs described by . |
. In this framework. it 15 therefore not surprising that the bump and the non-linear far-UV rise appear to be unrelated from an observational point of view: this reflects the variations of the spectral properties of PAHs in different charge states. | In this framework, it is therefore not surprising that the bump and the linear UV rise appear to be unrelated from an observational point of view: this reflects the variations of the spectral properties of PAHs in different charge states. |
Table [ lists the line luminosity densities as observed by summing up the data after extinction correction aud conleteuness correction. combined with the star forming rates for [fy=TO PABMye. d. | Table 4 lists the line luminosity densities as observed by summing up the data after extinction correction and completeness correction, combined with the star forming rates for $H_{\rm 0}=70$ $^{-1}$ $^{-1}$. |
The completeness correction is done by calculating the area under the Schechter function according to: A comparison between the SER derived from the sium of the observed. galaxy line eiissious aud the area under the Schechter function using the parameters marked in the plots shows that for the case of 2=0.25. the completcucss correction is negligible. while for the +=1.2 sample it ainounts to a factor 22. | The completeness correction is done by calculating the area under the Schechter function according to: A comparison between the SFR derived from the sum of the observed galaxy line emissions and the area under the Schechter function using the parameters marked in the plots shows that for the case of $z=0.25$, the completeness correction is negligible, while for the $z=1.2$ sample it amounts to a factor $>2$. |
Suiall uncertainties iu the slope paraieters à may lead to large errors in the total Iuuunosity aud SER. | Small uncertainties in the slope parameters $\alpha$ may lead to large errors in the total luminosity and SFR. |
For l.L which is the mean value in our curves. à variation of the slope paramcter a by £0.1 would alter the SFR bx col». | For $\alpha \sim 1.4$ , which is the mean value in our curves, a variation of the slope parameter $\alpha$ by $\pm 0.1$ would alter the SFR by $\pm \sim 15\%$. |
The SER values (Table. 1) show an increase by about a factor of 10 between 2=0.25 and +=1.2. | The SFR values (Table 4) show an increase by about a factor of 10 between $z=0.25$ and $z=1.2$. |
This is in goodl agreement with the slope of the SFR data published by Ποσο et al. ( | This is in good agreement with the slope of the SFR data published by Hogg et al. ( |
1998). who used the Uj] line for all redshifts from 0.2 to 1.2. aud with the slope derived from UV fiux densities (Lally et al. | 1998), who used the ] line for all redshifts from 0.2 to 1.2, and with the slope derived from UV flux densities (Lilly et al. |
1996). | 1996). |
As can be secu. the values at 2=(h88 and 1.2 are rather sensitive to the extinction correction applied. | As can be seen, the values at $z=0.88$ and 1.2 are rather sensitive to the extinction correction applied. |
The dispersion of “Ε.Τ in the reddening relation (Sect. | The dispersion of $\pm0.1$ in the reddening relation (Sect. |
3.5) vields a systematic error bar of EI0 for both SFR values. | 3.5) yields a systematic error bar of $\pm40\%$ for both SFR values. |
The SFR densities derived from the ΠΠ line follow the trend of the other data within the errors. | The SFR densities derived from the ] line follow the trend of the other data within the errors. |
Using a large sample of galaxies aud averaged line ratios. this line vields satisfactory results. | Using a large sample of galaxies and averaged line ratios, this line yields satisfactory results. |
Large scale structure expresses itself im stronglv varviug uunber counts roni field to feld aud from redshift to redshüft. | Large scale structure expresses itself in strongly varying number counts from field to field and from redshift to redshift. |
This effect can nicely be seen in the Mgl diagrams for the CADIS ealaxies (Fried et al. | This effect can nicely be seen in the $M_{\rm B} - z$ diagrams for the CADIS galaxies (Fried et al., |
2001). | 2001). |
Since the redshift iutervals of our FP windows correspoud roughly to the typical size of a LSS cell AAIpec). an even more distinct effect is expected here. | Since the redshift intervals of our FP windows correspond roughly to the typical size of a LSS cell Mpc), an even more distinct effect is expected here. |
Tudeed. the line Imuinositv density varies from field to field by a factor ofabout 2 (ratio of highest to lowest values for ideutical redshift intervals). | Indeed, the line luminosity density varies from field to field by a factor of about 2 (ratio of highest to lowest values for identical redshift intervals). |
An average over only three fields leaves an uncertainty of the order1054. | An average over only three fields leaves an uncertainty of the order. |
. For optically selected. galaxy. samples the SFR cau also ο estimated by means of the UV chussion. | For optically selected galaxy samples the SFR can also be estimated by means of the UV emission. |
While the chussion lines from regious measure the rate of nassive stars born less than a few willion vears ago. the integrated UV flux from short-lived stars indicates the star ornation rate of a somewhat older star population. | While the emission lines from regions measure the rate of massive stars born less than a few million years ago, the integrated UV flux from short-lived stars indicates the star formation rate of a somewhat older star population. |
For aree redshifts (2>> 2.8) where the observation of euission Ines is extremely difficult and the Lo line is often buried o duterual dust extinction. Steidel et al. ( | For large redshifts $z > 2.8$ ) where the observation of emission lines is extremely difficult and the $\alpha$ line is often buried by internal dust extinction, Steidel et al. ( |
1999) used the UV flux a mum to derive star formation rates at hiel redshifts. | 1999) used the UV flux at nm to derive star formation rates at high redshifts. |
The conversion factor between UV flux and SER is still a matter of debate aud is sensitive to the metallicity acl IME used in the model calculation (Clazebrook et al. | The conversion factor between UV flux and SFR is still a matter of debate and is sensitive to the metallicity and IMF used in the model calculation (Glazebrook et al. |
1999). but also to the intrinsic dust extinction adopted. | 1999), but also to the intrinsic dust extinction adopted. |
We used for the extinction in the UV. continua the relation eiven bv Calzetti et al. ( | We used for the extinction in the UV continuum the relation given by Calzetti et al. ( |
1091: 7=O.57p!. where rp" and n are the optical depths for starlielit coutiuuuni enüssion and for lue emission from the regions. respectively. | 1994): ${\tau_{\rm B}}^c = 0.5 {\tau_{\rm B}}^l$, where ${\tau_{\rm B}}^c$ and ${\tau_{\rm B}}^l$ are the optical depths for starlight continuum emission and for line emission from the regions, respectively. |
For the redshifts discussed above. the nuu (rest fune) contin corresponds to observed waveleusths inbetween nn aud 616nun. | For the redshifts discussed above, the nm (rest frame) continuum corresponds to observed wavelengths inbetween nm and nm. |
Except for the lowest redshift bin (2= 125). the UV flux density at uum can thus be casily determined fron the CADIS filter data. | Except for the lowest redshift bin $z=0.25$ ), the UV flux density at nm can thus be easily determined from the CADIS multi-filter data. |
In the case of 2=25. we use the nuu filter data if available. corresponding to a rest wavelength of Gunn instead of 280n1nn. | In the case of $z=0.25$, we use the nm filter data if available, corresponding to a rest wavelength of nm instead of nm. |
To account for this1 wavelength difference. a correction factor of (280/316)? (corresponding to a flat spectrmu £,= constant) ds applied to the UV fux densities. | To account for this wavelength difference, a correction factor of $^2$ (corresponding to a flat spectrum $F_{\nu}=constant$ ) is applied to the UV flux densities. |
The relation between UV. coutimaiun at nni and enission line fluxes is shown iu 111. | The relation between UV continuum at nm and emission line fluxes is shown in 11. |
Iu this eraph. the [Ou] fluxes are scaled wp by a factor 1.1 to account for | In this graph, the ] fluxes are scaled up by a factor 1.1 to account for |
star HD 38529. for which their slopes. metallicities and y values indicate that this object should be taken with caution. | star HD 38529, for which their slopes, metallicities and $\chi^2$ values indicate that this object should be taken with caution. |
The derived parameters of this star are based only on 3 Fell lines (and 24 Fel lines). | The derived parameters of this star are based only on 3 FeII lines (and 24 FeI lines). |
In literature. some values have been rounded in T,;;. log g and & to within 50 K or 10 K. 0.05 dex and 0.1 km/s. respectively (see.forexample.Gonza-lez1997.1998:Gonzalezetal. 1999). | In literature, some values have been rounded in $_{eff}$, log g and $\xi$ to within 50 K or 10 K, 0.05 dex and 0.1 km/s, respectively \citep[see, for example,][]{gonzalez97,gonzalez98,gonzalez99}. |
There is a good agreement between parameters and previous works from literature within the errors. which is logic taking into account that use a very similar method. | There is a good agreement between parameters and previous works from literature within the errors, which is logic taking into account that use a very similar method. |
However the Figure shows that there ts a dispersion in the values of the parameters (particularly /ogg and £) and probably a sistematic tendence for the metallicity (0.01 dex below literature values. see next discussion). | However the Figure shows that there is a dispersion in the values of the parameters (particularly $log g$ and $\xi$ ) and probably a sistematic tendence for the metallicity $\sim$ 0.01 dex below literature values, see next discussion). |
The values of AT,;; and Alog g also seems to slightly decrease with T,;; and log g. respectively. | The values of $\Delta$ $_{eff}$ and $\Delta$ log g also seems to slightly decrease with $_{eff}$ and log g, respectively. |
The median of the differences for the fundamental parameters compared to literature are 24 K. 0.06 dex. 0.03 dex and 0.08 km/s. corresponding to Το. /ogg. |Fe/H and &. respectively. | The median of the differences for the fundamental parameters compared to literature are 24 K, 0.06 dex, 0.03 dex and 0.08 km/s, corresponding to $_{eff}$, $log g$, [Fe/H] and $\xi$, respectively. |
The higher differences in the parameters are 118 K (16 Cyg B). 0.30 dex (HD 27442). 0.16 dex (16 Cvg B) and 0.31 km/s (47 Uma). corresponding to Το. ους. |Fe/H and &. respectively. | The higher differences in the parameters are 118 K (16 Cyg B), 0.30 dex (HD 27442), 0.16 dex (16 Cyg B) and 0.31 km/s (47 Uma), corresponding to $_{eff}$, $log g$, [Fe/H] and $\xi$ , respectively. |
Now we discuss the possible origi of the dispersions and probable tendences observed in the Figure 3.. | Now we discuss the possible origin of the dispersions and probable tendences observed in the Figure \ref{fig1}. |
We tested FUNDPAR modifying significatively the values of the weights In the function v and verified that the difference in [Fe/H] with literature and the dispersion in the other parameters changes very slightly. | We tested FUNDPAR modifying significatively the values of the weights in the function $\chi^2$ and verified that the difference in [Fe/H] with literature and the dispersion in the other parameters changes very slightly. |
Then the weights do not seem to be the cause. | Then the weights do not seem to be the cause. |
We use a method similar to literature. however it is not totally identical. | We use a method similar to literature, however it is not totally identical. |
FUNDPAR use different Kurucz model atmospheres (Castelli&Kuruez2003) than those used in literature (most of them are previous to the creation of the ODFNEW models). | FUNDPAR use different Kurucz model atmospheres \citep{castelli-kurucz03}
than those used in literature (most of them are previous to the creation of the ODFNEW models). |
The code use model atmospheres derived by ATLAS using ODFNEW opacities and solar abundances from Grevesse&Sauval(1998) instead of Anders&Grevesse (1989). | The code use model atmospheres derived by ATLAS using ODFNEW opacities and solar abundances from \citet{grevesse-sauval98} instead of \citet{anders-grevesse89}. |
. The new models present differences compared to older Kuruez models (Kurucz1990) such as the solar abundances. the replacement of TiO and H:0 molecular lines. some HI quasi-molecular absorptions. etc. | The new models present differences compared to older Kurucz models \citep{kurucz90} such as the solar abundances, the replacement of TiO and $_2$ 0 molecular lines, some HI quasi-molecular absorptions, etc. |
taken into account in the NEWODF opacities. | taken into account in the NEWODF opacities. |
Preliminar improvements are the U-B and u-b color indices forK. all color indices for cooler stars. and the better modeling for the upper layers of cool and giant stars (Castelli&Kurucz2003). | Preliminar improvements are the U-B and u-b color indices for, all color indices for cooler stars, and the better modeling for the upper layers of cool and giant stars \citep{castelli-kurucz03}. |
. In this example. the model atmospheres are computed with convection. (mixing-lenght parameter=1.25). overshooting (Wz1) and damping in the Unnsold approximation but multiplied by a factor as suggested by Blackwelletal.(1995). | In this example, the model atmospheres are computed with convection (mixing-lenght $=$ 1.25), overshooting (W=1) and damping in the Ünnsold approximation but multiplied by a factor as suggested by \citet{blackwell95}. |
. The use of the fourth condition within the function y is (possibly) another difference with previous studies. | The use of the fourth condition within the function $\chi^2$ is (possibly) another difference with previous studies. |
Literature works surely take this into account. however it is not totally clear for us how. | Literature works surely take this into account, however it is not totally clear for us how. |
Finally. we use the MOOG 2009 of the program instead of the 2002 version used in literature. although we expect almost the same abundance values from both versions. | Finally, we use the MOOG 2009 of the program instead of the 2002 version used in literature, although we expect almost the same abundance values from both versions. |
These differences. at least in part. produce the slightly disimilar values showed in the Figure 3.. which suggest that FUNDPAR use probably a different metallicity scale than used in literature. | These differences, at least in part, produce the slightly disimilar values showed in the Figure \ref{fig1}, which suggest that FUNDPAR use probably a different metallicity scale than used in literature. |
A complete comparison. require the exact knowledge of all involved details used in. the literature calculation. | A complete comparison require the exact knowledge of all involved details used in the literature calculation. |
Our intention is to clearly present all the asumptions used in FUNDPAR. in the model atmospheres and within the code. | Our intention is to clearly present all the asumptions used in FUNDPAR, in the model atmospheres and within the code. |
In The Figure 4 we show the histogram distributions of the uncertainties derived in Το. /ogg. Fe/H| and €. | In The Figure \ref{histog.uncer} we show the histogram distributions of the uncertainties derived in $_{eff}$, $log g$, [Fe/H] and $\xi$. |
The densely and slightly shaded histograms correspond to uncertaities derived following Gonzalez&Vanture(1998) and using the y function. respectively. | The densely and slightly shaded histograms correspond to uncertaities derived following \citet{gonzalez-vanture98}
and using the $\chi^2$ function, respectively. |
Some distributions present a peak in à common uncertainty value. such as ~0.05 dex in the |Fe/H] distribution and ~30 K in the distribution of effective temperature. | Some distributions present a peak in a common uncertainty value, such as $\sim$ 0.05 dex in the [Fe/H] distribution and $\sim$ 30 K in the distribution of effective temperature. |
We see that the errors derived using both methods are comparable. | We see that the errors derived using both methods are comparable. |
The metallicities presented in the Table 4 correspond to a group of exoplanet host stars. | The metallicities presented in the Table \ref{tabla.pars} correspond to a group of exoplanet host stars. |
The median of the group ts 0.17 dex with a dispersion of 0.22 dex. | The median of the group is 0.17 dex with a dispersion of 0.22 dex. |
In the Figure 5. we present the histogram of the metallicity distribution. | In the Figure \ref{histog.feh} we present the histogram of the metallicity distribution. |
Then. as an example of practical use ofFUNDPAR. we verified the metal-rich nature of main sequence stars with low mass companions. a fact known from the literature (see.forexam-ple.Gonzalez1997;Santosetal. 2000). | Then, as an example of practical use of, we verified the metal-rich nature of main sequence stars with low mass companions, a fact known from the literature \citep[see, for example,][]{gonzalez97,santos00}. |
. We have implemented a fortran algorithm available from the web that estimate fundamental parameters of solar type stars. requirmg only the measure of Fe equivalent widths. | We have implemented a fortran algorithm available from the web that estimate fundamental parameters of solar type stars, requiring only the measure of Fe equivalent widths. |
The final solution should verify the three conditions of the standard method: (re. ionization equilibrium). independence of the metallicity with the excitation potential (Le. excitation equilibrium) and. with respect to the equivalent widths. | The final solution should verify the three conditions of the standard method: (i.e. ionization equilibrium), independence of the metallicity with the excitation potential (i.e. excitation equilibrium) and with respect to the equivalent widths. |
We also add another condition: the input netallicity used in themodel atmosphere should be similar to the resulting metallicity from the equivalent widths. | We also add another condition: the input metallicity used in themodel atmosphere should be similar to the resulting metallicity from the equivalent widths. |
We taken into account these conditions in one variable called v. adopting an expression which include the weightswy. | We taken into account these conditions in one variable called $\chi^2$, adopting an expression which include the weights. |
....w4. FUNDPAR use Kuruez model atmospheres with the NEWODF opacities (Castelli&Kuruez2003).. solar-scaled abundances from Grevesse&Sauval(1998) and the 2009 version of the MOOG program. | FUNDPAR use Kurucz model atmospheres with the NEWODF opacities \citep{castelli-kurucz03}, solar-scaled abundances from \citet{grevesse-sauval98} and the 2009 version of the MOOG program. |
Different details could be selected. such às the mixing-lenght parameter. the overshooting and the damping of the lines. for instance. | Different details could be selected, such as the mixing-lenght parameter, the overshooting and the damping of the lines, for instance. |
We have planed a new version that include the option of use the WIDTH9 program instead of MOOG deriving abundances from equivalent widths. | We have planed a new version that include the option of use the WIDTH9 program instead of MOOG deriving abundances from equivalent widths. |
The code include the derivation of the uncertainty in the 4 parameters following the eriterta of Gonzalez Vanture (1998) and another uncertainty estimation using the y function. | The code include the derivation of the uncertainty in the 4 parameters following the criteria of Gonzalez Vanture (1998) and another uncertainty estimation using the $\chi^2$ function. |
We verified the metal-rich natureof a group of exoplanet host stars. | We verified the metal-rich natureof a group of exoplanet host stars. |
The parameters derived with are in agreement with previous works in literature. | The parameters derived with are in agreement with previous works in literature. |
that derived by Lubow Pringle (1993). which describes the full set of wave modes of a vertically isothermal cisc. | that derived by Lubow Pringle (1993), which describes the full set of wave modes of a vertically isothermal disc. |
Oeilvic Lubow (1999) slightly extended that work to allow for the possibility that Q.z©. and found that. for a disc undergoing isothermal perturbations (= 1). where w=iuo mtm is the azimuthal mode number and n ds the vertical mode number. | Ogilvie Lubow (1999) slightly extended that work to allow for the possibility that $\Omega_z\ne\Omega$, and found that, for a disc undergoing isothermal perturbations $\gamma=1$ ), where $\hat\omega=\omega-m\Omega$, $m$ is the azimuthal mode number and $n$ is the vertical mode number. |
For the case m=I. n—0. which corresponds to a tilt or warp. we find Unlike equation (7)). this dispersion relation 15 valid when the quantities ||ΩΩ]. απΩΣ! and ALE ave of order unity. | For the case $m=1$, $n=0$, which corresponds to a tilt or warp, we find Unlike equation \ref{dispersion}) ), this dispersion relation is valid when the quantities $|1-\Omega_z^2/\Omega^2|$, $|1-\kappa^2/\Omega^2|$ and $kH$ are of order unity. |
The two equations clearly agree well in the limit of low frequency. je/Q]«1. and agree exactly when w=O (see Fig. | The two equations clearly agree well in the limit of low frequency, $|\omega/\Omega|\ll1$, and agree exactly when $\omega=0$ (see Fig. |
1). | 1). |
The case w=0 is of some interest. as it relates to warped discs that are independent of time. | The case $\omega=0$ is of some interest, as it relates to warped discs that are independent of time. |
When both nodal ancl apsidal precession are present. and are either both prograde or both retrograde. the dispersion relation [or w=0 indicates that the spatial structure of the warp is oscillatory in character. its radial wavenumber being given by An oscillatory solution of this kind may be interpreted as a bending wave having zero phase velocity but non-zero eroup velocity. | When both nodal and apsidal precession are present, and are either both prograde or both retrograde, the dispersion relation for $\omega=0$ indicates that the spatial structure of the warp is oscillatory in character, its radial wavenumber being given by An oscillatory solution of this kind may be interpreted as a bending wave having zero phase velocity but non-zero group velocity. |
Conversely. when the nodal and apsical precession are in opposite senses. a steady warp is spatially evanescent. | Conversely, when the nodal and apsidal precession are in opposite senses, a steady warp is spatially evanescent. |
This situation is qualitatively similar to the case when the viscosity is large and the warp satisfies a diffusion equation. | This situation is qualitatively similar to the case when the viscosity is large and the warp satisfies a diffusion equation. |
Lt applies when the orbital. epievclic and vertical frequencies are derived from an axisvmmetric Newtonian gravitational potential that satisfies Laplace's equation in the mid-plane of the disc. llowever. this need. not. hold in the metric of a black hole. | It applies when the orbital, epicyclic and vertical frequencies are derived from an axisymmetric Newtonian gravitational potential that satisfies Laplace's equation in the mid-plane of the disc, since this implies However, this need not hold in the metric of a black hole. |
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