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This last approach relies ou je assuniption made for the size «ot the cutting region. ut in the range of sya Iuninosiics that we measured us effect is overwheluect 1xw the internal dispersion found oei our sample and the uncertainty affecting the slope of ιο correlation between the Pan aid the san huninosity surface densities (Calzettietal.2007:Diaz-Sautosct 2008.. see also Alouso-ILercoroetal. 20063). | This last approach relies on the assumption made for the size of the emitting region, but in the range of $\mu$ m luminosities that we measured this effect is overwhelmed by the internal dispersion found in our sample and the uncertainty affecting the slope of the correlation between the $\alpha$ and the $\mu$ m luminosity surface densities \citealt{Cal:07,Diaz:08}, see also \citealt{Alonso:06}) ). |
This relation is also consistent with the Pao aud the dy ni Iunuinosities that were recently measured in a subuilimieter lensed ealaxy at zo2.5 by Papewichetal.(209). | This relation is also consistent with the $\alpha$ and the $\mu$ m luminosities that were recently measured in a submillimeter lensed galaxy at $z \sim 2.5$ by \citet{Papo:09}. |
.. It suggests that it could still be applied to hiebh-redshiT SOUECOS CVC though it was derived froi observations of the neurbv Universe. | It suggests that it could still be applied to high-redshift sources even though it was derived from observations of the nearby Universe. |
Given the good agreenent between these cdiffereut relationships. we thus inferred the total IR luuinosity and star formation rate ο the COSMOS. MIPS sources using the linear relation otained by Elbaz et al. ( | Given the good agreement between these different relationships, we thus inferred the total IR luminosity and star formation rate of the COSMOS MIPS sources using the linear relation obtained by Elbaz et al. ( |
2011. submitted) as well as our refeqisfrl..— baring dn ndud that the other possible methods would have led ο very simular results. | 2011, submitted) as well as our \\ref{eq:sfr1}, baring in mind that the other possible methods would have led to very similar results. |
Dased on these deteriuuations o| Lg aud SER. we computed the contributions of the «ifferent populations of MIPS-selected sources discussed. in 33 to the IR Iuuunositv density of the Universe aid the cosmüc star formation rate density iun three recshift bins: 1.5«:L9. 1.9<2.<23 and 2.3<Dox2.5 | Based on these determinations of $L_{\rm IR}$ and SFR, we computed the contributions of the different populations of MIPS-selected sources discussed in 3 to the IR luminosity density of the Universe and the cosmic star formation rate density in three redshift bins: $1.5<z<1.9$, $1.9<z<2.3$ and $2.3<z<2.7$. |
To reach this goal we derived Iuuinositv fuuctio win a wav simular to the nethod described 1 55.1. and we iutegrated the LES above the hunumositv correspouding to the 21 jan flux iuit of our survev PF,linLOS παν) at the median of cach redshift bij | To reach this goal we derived luminosity functions in a way similar to the method described in 5.1, and we integrated the LFs above the luminosity corresponding to the 24 $\mu$ m flux limit of our survey $F_{24\mu m} = 0.08$ mJy) at the median of each redshift bin. |
Suc| we only aim at determining iow the various galaxy sul-Saluples deteeted with MIPS contribute to the otal SER cleusity. we did not assuue any extra]xdation o ‘the ALIPS-selected source population o the zn end of the luminosity fuucjon. | Since we only aim at determining how the various galaxy sub-samples $detected$ with MIPS contributed to the total SFR density, we did not assume any extrapolation of the MIPS-selected source population to the faint end of the luminosity function. |
Our results are shown in rofügrhogsFR..whichalsoi Hustratesthec colutionofthe total (Μεoptica ηςleh osten 1h Reeletuesoct ptoz 2.55aid the cosmuc star formation lisory compiled b Ilopkius&Deacon(2006). | Our results are shown in \\ref{fig:rho_SFR}, which also illustrates the evolution of the total IR luminosity density inferred by \citet{Rodi:10} up to $z \sim 2.5$ and the cosmic star formation history compiled by \citet{Hopkins:06}. |
The 1306 sources witli optical counterparts are obviously not taken iuto account rere bu we already saw that they coutrinte ouly a sinall raction of the total uuubcr of galaxies ideutified in t1ο COSMOS field. | The 1306 sources with no optical counterparts are obviously not taken into account here but we already saw that they contribute only a small fraction of the total number of galaxies identified in the COSMOS field. |
Furtheruxxe. Inost oftrese uuildeutified sources have low μι. fluxes close to the seusitivitv inut of the COSMOS ÀIPS data. | Furthermore, most of these unidentified sources have low $\mu$ m fluxes close to the sensitivity limit of the COSMOS MIPS data. |
This implies that their contribution to the IR huninosity density is even less han their coutribution O the ummber of uuc-IR selected sources, aud therefore the lack of identification for these sources ean not severely bias our results. | This implies that their contribution to the IR luminosity density is even less than their contribution to the number of mid-IR selected sources, and therefore the lack of identification for these sources can not severely bias our results. |
The top panel of refüe:LE illustrates the hhuuimositv functions associated with the MIPS ποιαον identi&ed with the four color selection techniques descried iu 33. compared to the total Sjiu hnuuiuositv function observed at L3?«i2.5 in the COSMOS fickl. | The top panel of \\ref{fig:LF} illustrates the luminosity functions associated with the MIPS sources identified with the four color selection techniques described in 3, compared to the total $8\mu$ m luminosity function observed at $1.7<z<2.3$ in the COSMOS field. |
The later is also compared with the sg u LFs obtained by Caputictal.(2007) and Rodighieroctal.(2010) from the GOODS aud SWIRE suvevs. which show relatively good aereeineut with our estiuates, | The later is also compared with the $8\mu$ m LFs obtained by \citet{Caputi:07} and \citet{Rodi:10} from the GOODS and SWIRE surveys, which show relatively good agreement with our estimates. |
Except for the selection of Optically-Faiit Πιο] ealaxies. our analysis seenis to indicate that the fraction of MIPS lieh+vedshift ealaxies selected withthe B:NK. ΟΛΠΕΝ iuxLIRAC peaker criteria does not critically depend on the Span luminosity itself. | Except for the selection of Optically-Faint IR-bright galaxies, our analysis seems to indicate that the fraction of MIPS high-redshift galaxies selected with the $BzK$, BM/BX and IRAC peaker criteria does not critically depend on the $8\mu$ m luminosity itself. |
To better illustrate this result we rc]xoduced the evolution of these fractions with Γκ iu the bottom paucl of the figure. | To better illustrate this result we reproduced the evolution of these fractions with $L_{8\mu m}$ in the bottom panel of the figure. |
To increase the statistics we extended the redshift biu up to 1.5«<2a. which corresponds to the largest redshitt liice where the efficiencies of our differcit color selections Cai be simultaneously compared with one another. | To increase the statistics we extended the redshift bin up to $1.5<z<2.8$, which corresponds to the largest redshift range where the efficiencies of our different color selections can be simultaneously compared with one another. |
We see again that the fractions of MIPS galaxies selected as DK aud BAL/BN sources or as IRAC yealsers are fairly constait over the rauge of sya huninosities probed by the MIPS observations. | We see again that the fractions of MIPS galaxies selected as $BzK$ and BM/BX sources or as IRAC peakers are fairly constant over the range of $\mu$ m luminosities probed by the MIPS observations. |
They correspond to the average fractions we had already «lerived based on their ΙΟ clesity and their redslüft «listribution. | They correspond to the average fractions we had already derived based on their number density and their redshift distribution. |
While the abseuce of correlation between L5, aud the fraction of DA sources or IRAC peakers is not necessarily unexpected. the result tji we fud for the DM/DX ealaxies niav deserve further explanations. | While the absence of correlation between $L_{8\mu m}$ and the fraction of $BzK$ sources or IRAC peakers is not necessarily unexpected, the result that we find for the BM/BX galaxies may deserve further explanations. |
Lhudeed we attributed the uas observed in the DM/DX selection ο the effect of dust exinction (see 220). | Indeed we attributed the bias observed in the BM/BX selection to the effect of dust extinction (see 2c). |
(ναι 1ο elobal trei that exists between galaxy bolometric unumosities aud the IR/UV luminosity ratio (e.g..Dell.Ww WO3).. one could have exvected some broad correlation ETWEE1 the sy i luuinesity aud the fraction of MIPS [umigh-vedshift sources άσσος, by the ΕΛΙΕΣ criterion. | Given the global trend that exists between galaxy bolometric luminosities and the IR/UV luminosity ratio \citep[e.g.,][]{Bell:03}, one could have expected some broad correlation between the $\mu$ m luminosity and the fraction of MIPS high-redshift sources missed by the BM/BX criterion. |
Iu the most huninous gal:wies though. the reddening of 1ο does not correlate anvinere with the μοι pd. |PMG Liy (eg.Godaderetal.2002:Reddy.al...2006). | In the most luminous galaxies though, the reddening of the optical light does not correlate anymore with the dust obscuration measured by the excess of $_{\rm IR}$ over $_{\rm UV}$ \citep[e.g.,][]{Goldader:02,Reddy:06}. |
. This can be uierstood. Kx mstauce. 1f ou one hare UVfoplca enissjon originates from spatialσοκ reeious of star Ormation where dust 5erains follow a chuupy distribution. while ou the other haud predomunaitly cones fron) very conrpact are icalbthick regio lcose the ceuter of galaxies. | This can be understood, for instance, if on one hand the UV/optical emission originates from spatially-extended regions of star formation where dust grains follow a rather clumpy distribution, while on the other hand the IR light predominantly comes from very compact and optically-thick regions close the center of galaxies. |
Since the ALIPS observations o| COSMOS are oulv seusitive 1 Very bright cud of the high-redshift galaxy Iuninosiv funclon. we are likely in his huninositv recie where the ΕΕV extinction is not directly correlated with the evel of star-forming activity. hence explaining the lack of trend between the IR uuinosity aud the detectability of galaxies based on their CVOR colors. | Since the MIPS observations of COSMOS are only sensitive to the very bright end of the high-redshift galaxy luminosity function, we are likely in this luminosity regime where the $E(B-V)$ extinction is not directly correlated with the level of star-forming activity, hence explaining the lack of trend between the IR luminosity and the detectability of galaxies based on their $U_nGR$ colors. |
FinaIv. it is also possible that our result at the highest. huninosities.eye (Lu,>+Di 1031.) isB slightly affectec by sonie ACN contribution. which can become significant at bielit 214g u fluxes | Finally, it is also possible that our result at the highest luminosities $_{8\mu m} > 10^{12} L_\odot$ ) is slightly affected by some AGN contribution, which can become significant at bright $\mu$ m fluxes |
and orbital eccentricily (6, = 0.0. 0.2. 0.4. 0.6. 0.3. and 0.95). | and orbital eccentricity $e_o$ = 0.0, 0.2, 0.4, 0.6, 0.8, and 0.95). |
The apoclustron is function ol both semi-major axis ancl eccentricity. S,=a,(1+ερ). | The apoclustron is function of both semi-major axis and eccentricity, $S_o = a_o (1 + e_o)$. |
Therefore and if we assume that the two clusters are point-like masses moving in Ixeplerian orbits. the initial orbital period of the svstem can be approximated by where Af, and M» are the cluster masses. the apoclustron is in pe aud (he masses in AL. (see actual values in Table 1)). | Therefore and if we assume that the two clusters are point-like masses moving in Keplerian orbits, the initial orbital period of the system can be approximated by where $M_1$ and $M_2$ are the cluster masses, the apoclustron is in pc and the masses in $M_{\odot}$ (see actual values in Table \ref{results}) ). |
Stellar masses in the range [0.17. 10.0] M... are drawn from a sSalpeterian IAIF with an average stellar mass of 0.5 M... | Stellar masses in the range [0.17, 10.0] $M_{\odot}$ are drawn from a Salpeterian IMF with an average stellar mass of 0.5 $M_{\odot}$. |
Salpeter (1955) used the observed huminosity [unction lor the Solar Neighborhood and theoretical evolution times to derive an initial mass function CALIF) which may be approximated by a power-law: where n(n) is the number of stars per unit mass interval. | Salpeter (1955) used the observed luminosity function for the Solar Neighborhood and theoretical evolution times to derive an initial mass function (IMF) which may be approximated by a power-law: where $n(m)$ is the number of stars per unit mass interval. |
The value of a is 2.35 [or masses between 0.4 and 10.0 Δι. | The value of $\alpha$ is 2.35 for masses between 0.4 and 10.0 $M_{\odot}$. |
The IME used (a single power-law) and the mass range of stars ([0.17. 10.0] M.) are equivalent to (he realistic canonical two-part power-law over the mass range (0.08. 10.0] M... as described in (e.g.) INvoupa. (2008). | The IMF used (a single power-law) and the mass range of stars ([0.17, 10.0] $M_{\odot}$ ) are equivalent to the realistic canonical two-part power-law over the mass range [0.08, 10.0] $M_{\odot}$ as described in (e.g.) Kroupa (2008). |
Neglecting stars more massive (han 10 AL. is a good approximation since such stars are rare (see. e.g.. the ΕΒΕ’ cluster mass relation of Weidner et al. | Neglecting stars more massive than 10 $M_{\odot}$ is a good approximation since such stars are rare (see, e.g., the $m_{max}$ -star cluster mass relation of Weidner et al. |
2010). | 2010). |
All stars are started on a zero-age main sequence with a uniform composition of hydrogen. X. = 0.7. helium. Y = 0.28. ancl metallicity. Z = 0.02. | All stars are started on a zero-age main sequence with a uniform composition of hydrogen, $X$ = 0.7, helium, $Y$ = 0.28, and metallicity, $Z$ = 0.02. |
Stellar evolution is computed according to the algorithm described by Aarseth (2003). | Stellar evolution is computed according to the algorithm described by Aarseth (2003). |
Primordial binaries were not included but binary. aid multiple svstem lormation was allowed and observed. | Primordial binaries were not included but binary and multiple system formation was allowed and observed. |
External perturbations were represented by a fixed galactic tidal field. | External perturbations were represented by a fixed galactic tidal field. |
The tidal radius is given by the expression (e.g.. Aarseth 2003) where ο and B are (he Oorts constants of Galactic rotation. | The tidal radius is given by the expression (e.g., Aarseth 2003) where $A$ and $B$ are the Oort's constants of Galactic rotation. |
For a star located al that distance from the cluster. the central attraction of the eluster is balanced by the Galactic tidal force. | For a star located at that distance from the cluster, the central attraction of the cluster is balanced by the Galactic tidal force. |
The OorUs constants are chosen to be A= 14.4. B=-12.0 kms 'kpe ! (Binney Tremaine 2003). | The Oort's constants are chosen to be $A$ = 14.4, $B$ = -12.0 km $^{-1}$ $^{-1}$ (Binney Tremaine 2008). |
The cluster pair is assumed to move in a circular orbit at (he Solar Circle with no passing molecular clouds (see Aarseth 2003. pp. | The cluster pair is assumed to move in a circular orbit at the Solar Circle with no passing molecular clouds (see Aarseth 2003, pp. |
127-129. for additional details). | 127-129, for additional details). |
No escaping stars were removed [rom the calculations. | No escaping stars were removed from the calculations. |
The simulations presented here have been performed on a Dell Precision T5500 + Nvidia Tesla SLOTO svstem. | The simulations presented here have been performed on a Dell Precision T5500 + Nvidia Tesla S1070 system. |
The T5500 has 2 quad-core Intel Xeon E5540 processors at 2.53 Gllz. | The T5500 has 2 quad-core Intel Xeon E5540 processors at 2.53 GHz, |
to the neutron exposure during the s-process operation. | to the neutron exposure during the $s$ -process operation. |
Higher neutron exposures would produce strong Ba stars. while mild Ba stars have accreted matterial exposed to weaker neutron exposures. | Higher neutron exposures would produce strong Ba stars, while mild Ba stars have accreted matterial exposed to weaker neutron exposures. |
In the present work we report chemical abundances for two stars from the bulgelike sample of Michel Grenon (Grenon 1999, 2000). HD 11397 and HD 14282. | In the present work we report chemical abundances for two stars from the bulgelike sample of Michel Grenon (Grenon 1999, 2000), HD 11397 and HD 14282. |
Those stars belong to a kinematically selected sample of the solar neighborhood. the bulgelike stars (Pompéiia et al. | Those stars belong to a kinematically selected sample of the solar neighborhood, the bulgelike stars (Pompéiia et al. |
2002. 2003). | 2002, 2003). |
The bulgelike stars have very eccentric orbits. with eccentricities e > 0.25 and small pericentric distances with R, < 3.5 kpe. | The bulgelike stars have very eccentric orbits, with eccentricities e $>$ 0.25 and small pericentric distances with $_{\rm p}$ $\leq$ 3.5 kpc. |
Therefore they were probably born near the galactic bulge (Pompétia et al. | Therefore they were probably born near the galactic bulge (Pompéiia et al. |
2003. 2008). | 2003, 2008). |
The chemical abundances of the bulgelike stars have been previously studied (Pompéiia et al. | The chemical abundances of the bulgelike stars have been previously studied (Pompéiia et al. |
2003) anc their Ba. Zr. La and Y content have been determined. | 2003) and their Ba, Zr, La and Y content have been determined. |
HD 11397 and HD 14282 have shown abundance anomalies i their s-process content. with enhanced [Ba/Fe] and [La/Fe ratios when compared to the other bulgelike stars (see Fig. | HD 11397 and HD 14282 have shown abundance anomalies in their $s$ -process content, with enhanced [Ba/Fe] and [La/Fe] ratios when compared to the other bulgelike stars (see Fig. |
8 of Pompétia et al. | 8 of Pompéiia et al. |
2003). | 2003). |
HD 14282 also shows overabundant [Zr/Fe] ratios when compared to the other bulgelike stars. | HD 14282 also shows overabundant [Zr/Fe] ratios when compared to the other bulgelike stars. |
It order to properly classify the present stars and to perform a full analysis of their neutron-capture elements profile. we have inferred their abundances for C. N. La. Ba. Nd. Sm. Sr. Zr. Y. Mo. Ru. Ce. Pr. Gd. Dy. Hf and Pb. and for the +-process element Eu. | In order to properly classify the present stars and to perform a full analysis of their neutron-capture elements profile, we have inferred their abundances for C, N, La, Ba, Nd, Sm, Sr, Zr, Y, Mo, Ru, Ce, Pr, Gd, Dy, Hf and Pb, and for the $r$ -process element Eu. |
This work is divided as follows: in Sect. | This work is divided as follows: in Sect. |
2 we describe the observations and reduction procedures. the stellar parameters determination and the model atmospheres are described in Sect. | 2 we describe the observations and reduction procedures, the stellar parameters determination and the model atmospheres are described in Sect. |
3. the abundance calculations and their results are given in Sect. | 3, the abundance calculations and their results are given in Sect. |
4. in Sect. | 4, in Sect. |
5 we discuss our results. and in Sect. | 5 we discuss our results, and in Sect. |
6 we give a summary of the paper. | 6 we give a summary of the paper. |
Sample stars have been observed at the l.52m telescope of ESO (European Southern Observatory). La Silla. in September 1999. | Sample stars have been observed at the 1.52m telescope of ESO (European Southern Observatory), La Silla, in September 1999. |
The spectra were obtained using the FEROS spectrograph (Fiber-fed Extended Range Optical Spectrograph) with wavelength range 356 to 920 nm and a resolution of R = 48.000. | The spectra were obtained using the FEROS spectrograph (Fiber-fed Extended Range Optical Spectrograph) with wavelength range 356 to 920 nm and a resolution of R = 48,000. |
The detector is a back-illuminated CCD with 2048 X 4096 pixels of 15 ym size. | The detector is a back-illuminated CCD with 2048 X 4096 pixels of 15 $\mu$ m size. |
Reductions were performed using the DRS (online data reduction system of FEROS) and for a subsequent reduction. the TELLURIC. CONTINUUM. RVIDLINES and DOPCOR tasks of the IRAF package were applied. | Reductions were performed using the DRS (online data reduction system of FEROS) and for a subsequent reduction, the TELLURIC, CONTINUUM, RVIDLINES and DOPCOR tasks of the IRAF package were applied. |
Precise stellar parameters are fundamental for an acurate inference of the chemical abundances. | Precise stellar parameters are fundamental for an acurate inference of the chemical abundances. |
Chemical abundances are particularly sensitive to the temperature of the stellar atmosphere (Ty) and the surface gravity of the star (log ¢). | Chemical abundances are particularly sensitive to the temperature of the stellar atmosphere $_{\rm eff}$ ) and the surface gravity of the star (log g). |
Two other stellar parameters also play important role in the abundance determination: the metallicity of the star (|Fe/H]) and the microturbulence velocity (£). | Two other stellar parameters also play important role in the abundance determination: the metallicity of the star ([Fe/H]) and the microturbulence velocity $\xi$ ). |
First guesses for the stellar parameters were inferred from photometric data and distances from Hipparcos mission, | First guesses for the stellar parameters were inferred from photometric data and distances from Hipparcos mission. |
The final stellar parameters were calculated as follows: temperatures were derived requiring that Fe I lines with different excitation potentials give the same iron abundance: gravities and metallicities were inferred by forcing the agreement between Fe | and Fe II abundances: microturbulence velocities were calculated by requiring no slope in the [Fe/H] vs. EW (equivalent width) plot. | The final stellar parameters were calculated as follows: temperatures were derived requiring that Fe I lines with different excitation potentials give the same iron abundance; gravities and metallicities were inferred by forcing the agreement between Fe I and Fe II abundances; microturbulence velocities were calculated by requiring no slope in the [Fe/H] vs. EW (equivalent width) plot. |
A detailed description of the entire procedure is given in Pompéiia et al. ( | A detailed description of the entire procedure is given in Pompéiia et al. ( |
2008. in preparation). | 2008, in preparation). |
The adopted model atmospheres are an updated version of the plane-parallel MARCS model atmospheres with standard composition (Gustafssonetal..2003). | The adopted model atmospheres are an updated version of the plane-parallel MARCS model atmospheres with standard composition \citep{gus03}. |
. The final parameters and respective uncertainties are given in Table I. | The final parameters and respective uncertainties are given in Table 1. |
Abundance analysis has been performed by matching the synthetic profile with the observed spectrum. | Abundance analysis has been performed by matching the synthetic profile with the observed spectrum. |
The line synthesis code is an updated version of the code by Monique Spite (1967) (e.g. Cayrel et al. | The line synthesis code is an updated version of the code by Monique Spite (1967) (e.g. Cayrel et al. |
2001: Barbuy et al. | 2001; Barbuy et al. |
2003). | 2003). |
The line list and atomic references are given in. Table 3. | The line list and atomic references are given in Table 3. |
Hyperfine structure (HFS) has been applied for La. Ba. Eu and Pb. | Hyperfine structure (HFS) has been applied for La, Ba, Eu and Pb. |
The references for the HFS are: Rutten (1978) for Ba II. Lawler et al. ( | The references for the HFS are: Rutten (1978) for Ba II, Lawler et al. ( |
20012) for La IL. Lawler et al. ( | 2001a) for La II, Lawler et al. ( |
2001b) for Eu IL and Biémmont et al. ( | 2001b) for Eu II, and Biémmont et al. ( |
2000) for Pb I. Uncertainties on abundances were calculated for both HD 11397 and HD 14282. by verifying how much the variation of lo of the atmospheric parameters affects the output value of the synthesis program. here logA,,. and also by taking into account the standard deviation of the abundances for which more than 2 lines are available. | 2000) for Pb I. Uncertainties on abundances were calculated for both HD 11397 and HD 14282, by verifying how much the variation of $\sigma$ of the atmospheric parameters affects the output value of the synthesis program, here $\log{A_p}$, and also by taking into account the standard deviation of the abundances for which more than 2 lines are available. |
Table 3. shows the values for this calculation. | Table \ref{errab} shows the values for this calculation. |
Under the simplifying hypothesis of independent errors. the uncertainty of the output value is given by where AAy. AA,,. ΔΑ. and AA;. are the differences in A, due to the uncertainties in the temperature. metallicity. logg. and microturbulent velocity respectively and 4Η) is the standard deviation of the average. | Under the simplifying hypothesis of independent errors, the uncertainty of the output value is given by where $\Delta A_T$, $\Delta A_{mt}$, $\Delta A_l$, and $\Delta A_\xi$, are the differences in $A_p$ due to the uncertainties in the temperature, metallicity, $\log g$, and microturbulent velocity respectively and $sdm$ ' is the standard deviation of the average. |
The average value of A, (Αρ) is obtained by averaging the individual abundances of several lines. | The average value of $A_p$ $A_{pm}$ ) is obtained by averaging the individual abundances of several lines. |
Applying a propagation of errors and taking into account the uncertainty calculated with Eq. |.. | Applying a propagation of errors and taking into account the uncertainty calculated with Eq. \ref{erapinst}, |
the uncertainty on Ay, is where 7; 1s the number of lines for which AA;. AA,,,. AA;. and AA, were computed. | the uncertainty on $A_{pm}$ is where $n_l$ is the number of lines for which $\Delta A_T$ , $\Delta A_{mt}$ , $\Delta A_l$, and $\Delta A_\xi$ were computed. |
The uncertainty on the logarithm of Apm IS The abundanceΑμ log e(X) is related to the output of the synthesis program by log e(X) = logAj, + [Fe/H]. | The uncertainty on the logarithm of $A_{pm}$ is The abundance $\log\epsilon$ (X) is related to the output of the synthesis program by $\log\epsilon$ (X) = $\log{A_{pm}}$ + [Fe/H]. |
Therefore. the uncertainty 1s The relation between the abundance excess relative to iron [X/Fe] and the output value of thesynthesisprogram ts |X/Fe] | Therefore, the uncertainty is The relation between the abundance excess relative to iron [X/Fe] and the output value of thesynthesisprogram is [X/Fe] |
terpartIsaplanofthe112(2007)spulsuSct Tere. we are reporting a 100 ksNALAL observation towards the sky region around the X-ray (ROSAT) source IRNS J150131.1-273932. | }\tikzmark{mainBodyEnd1}
\offprints{A. A. Nucita}
\date{Submitted: XXX; Accepted: XXX}
{
\abstract
% context heading (optional)
% {} leave it empty if necessary
{Low-mass X-ray binaries are peculiar binary systems composed of a compact object and a low-mass star. Recently, a new class of these systems, known as symbiotic $X$-ray binaries (with a neutron star with a M-type giant companion), has been discovered.}
% aims heading (mandatory)
{Here, we present long-duration ${\it XMM}$ observations of the source 1RXS J180431.1-273932.}
% methods heading (mandatory)
{Temporal and spectral analysis of the source was performed along with a search for an optical counterpart. We used a Lomb-Scargle periodogram analysis for the period search and evaluated the confidence level using Monte-Carlo simulations.}
% results heading (mandatory)
{The source is characterized by regular pulses so that it is most likely a neutron star. A modulation of $494.1\pm0.2$\,s (3$\sigma$ error) was found with a confidence level of $>$99\%. Evidence of variability is also present, since the data show a rate of change in the signal of $\sim -7.7\times 10^{-4}$ counts s$^{-1}$ hr$^{-1}$. A longer observation will be necessary in order to determine if the source shows any periodic behavior. The spectrum can be described by a power law with photon index $\Gamma\sim 1$ and a Gaussian line at 6.6\,keV. The X-ray flux in the 0.2--10\,keV energy band is $5.4\times 10^{-12}$ erg s$^{-1}$ cm$^{-2}$.
The identification of an optical counterpart (possibly an M6III red-giant star with an apparent visual magnitude of $\simeq 17.6$) allows a conservative distance of $\sim 10$ kpc to be estimated. Other possibilities are also discussed.}
% conclusions heading (optional), leave it em\tikzmark{mainBodyStart2}mpty\tikzmark{mainBodyEnd2} \tikzmark{mainBodyStart3}if\tikzmark{mainBodyEnd3} \tikzmark{mainBodyStart4}necessary\tikzmark{mainBodyEnd4}
{Once the distance was estimated, we got an $X$-ray luminosity of $L_X\ut<6\times 10^{34}$ erg s$^{-1}$, which is consistent with the typical $X$-ray luminosity of a symbiotic LMXB system.}\tikzmark{mainBodyStart5}}\tikzmark{mainBodyEnd5}
\tikzmark{mainBodyStart6}} Here, we are reporting a $100$ ks observation towards the sky region around the $X$ -ray (ROSAT) source 1RXS J180431.1-273932. |
It was observed iu. October 2005 (Observation ID 30597) with both the EPIC MOS aud PN cameras (Turneretal.2001:Strüder20013) operating with a thin filter mode. | It was observed in October 2005 (Observation ID 30597) with both the EPIC MOS and PN cameras \citealt{mos,pn}) ) operating with a thin filter mode. |
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