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The azimuthal averaciis can be oue without or after dejxojection of he ealactic isk using an cstimation of the ealaxws inclinationj | The azimuthal averaging can be done without or after deprojection of the galactic disk using an estimation of the galaxy's inclination. |
s Unless he galaxy shape is simple ancl we]-definecd. applving he latter method is subject to 1ncertainties ue to the iultiplicitv of factors that can adfect7. such as. projecte dust absorption or real disk distorlolis. | Unless the galaxy shape is simple and well-defined, applying the latter method is subject to uncertainties due to the multiplicity of factors that can affect, such as, projected dust absorption or real disk distortions. |
The effecs of secine (due o Eartlis atmosphere aud scatterie within the telescope) on the inteusity and shape parameters of the fitted ellipses can be very iuportant and alter the reul sin the 1mer galaxy regions. | The effects of seeing (due to Earth's atmosphere and scattering within the telescope) on the intensity and shape parameters of the fitted ellipses can be very important and alter the results in the inner galaxy regions. |
Although our analvsis is mainv concerned with the extended comporent. we have| applied. au Inner cut-off radius to our surface brighttess profiles so that the profile decomposition (deseri below) is not affected by secius effecs, | Although our analysis is mainly concerned with the extended component, we have applied an inner cut-off radius to our surface brightness profiles so that the profile decomposition (described below) is not affected by seeing effects. |
Usnally in such studies the mimi cut-off radius is taken to be lf or ο]al to the secing PSF. but a nore elaborae method is described by Frauxctal.1f150... | Usually in such studies the minimum cut-off radius is taken to be half or equal to the seeing PSF, but a more elaborate method is described by \cite{franx89}. |
Using their formulac approximated for all core radii and taking;2M qo equivalent to the eror in f16 local surface Damituess due to the seeing effects. weie have: where Fy is the second order moment of the seeing PSF. | Using their formulae approximated for small core radii and taking to be equivalent to the error in the local surface brightness due to the seeing effects, we have: where $F_{2}$ is the second order moment of the seeing PSF. |
For our data. we assunie a "reasonde value for. (seo Έναςeal.1980: Jorgensenetal. 1992.. | For our data, we assume a “reasonable” value for (see \cite{franx89}; \cite{jorgensen92}. |
The FWIIM for each object Is mmcasured on several stars of each nuage aud the mean values are listed in Tables [ aud 5 of Paper I. Requiring that the error iu the local surface brightucss | The FWHM for each object is measured on several stars of each image and the mean values are listed in Tables 4 and 5 of Paper I. Requiring that the error in the local surface brightness |
a single low-redshift bin of line emission. or doesn't change. 29 compared to 29 for a single low-redshift bin of continuum emission. | a single low-redshift bin of line emission, or doesn't change, 29 compared to 29 for a single low-redshift bin of continuum emission. |
The analysis presented in. Section. 5.2. demonstrated. the sensitivity of the FoM to both the location. anc number. of redshift: regimes. ie. our optimization clearly preferred a single. οι redshift” bin. | The analysis presented in Section 5.2 demonstrated the sensitivity of the FoM to both the location, and number, of redshift regimes, i.e., our optimization clearly preferred a single “low redshift” bin. |
“Pherefore. we study here the benefits of splitting this single "low. redshift” bin into multiple thinner redshift bins. assuming these thinner bins are contiguous over the full redshift range of the single wider bin. | Therefore, we study here the benefits of splitting this single “low redshift” bin into multiple thinner redshift bins, assuming these thinner bins are contiguous over the full redshift range of the single wider bin. |
We expect these thinner redshift’ bins to have worse measurements of d4(2) and (2). compared to the single wider bin (because their volumes will be smaller). but this clisacdvantage could be overcome by the increased number of distance measurements available for fitting the cosmologica models (we do not consider correlations between the errors of dillerent redshift bins). | We expect these thinner redshift bins to have worse measurements of $d_A(z)$ and $H(z)$, compared to the single wider bin (because their volumes will be smaller), but this disadvantage could be overcome by the increased number of distance measurements available for fitting the cosmological models (we do not consider correlations between the errors of different redshift bins). |
Pherefore. we performed a ALCALC search as a Function of the number of low redshift bins anc discovered that the FoM quickly saturated at 2260 for al configurations with more than one “low redshift bins. i.e.. the optimal survey would split the single low-redshift regime into two equallysized bins. | Therefore, we performed a MCMC search as a function of the number of “low redshift” bins and discovered that the FoM quickly saturated at $\simeq260$ for all configurations with more than one “low redshift” bins, i.e., the optimal survey would split the single low-redshift regime into two equally–sized bins. |
These results should be revisitec for cillerent dark energy mocoels. | These results should be revisited for different dark energy models. |
We find that the best Figure-of-Moerit searching through all possible survey configurations is one which spends all its time at low redshift surveying an area of around. 6000. sq. | We find that the best Figure-of-Merit searching through all possible survey configurations is one which spends all its time at low redshift surveying an area of around 6000 sq. |
degs. targetting line-emission galaxies (survey A in Table SJ) | degs, targetting line-emission galaxies (survey A in Table \ref{binredshifttable}) ). |
The medium redshift of the bin is about 1.1. and.it will stretch from z~OSL4. | The medium redshift of the bin is about 1.1 andit will stretch from $z \sim 0.8
- 1.4$. |
Phe exposure time is 15 minutes per field-of-view. | The exposure time is 15 minutes per field-of-view. |
A additional high redshift bin is disfavoured. | A additional high redshift bin is disfavoured. |
This optimum is not highly peaked. as we can see by looking at Figure 4.. | This optimum is not highly peaked, as we can see by looking at Figure \ref{singlelowzbin}. |
Phe Hlattening of the FoM curve for area at about 6000. ddegs and the flat plateau at the top of the τμ curve indicates that deviations from these optimal values will result in only small changes in the Figure-of-Merit. also shown by the large width of the IHexibilitv bars in Table s.. | The flattening of the FoM curve for area at about 6000 degs and the flat plateau at the top of the $z_{low}$ curve indicates that deviations from these optimal values will result in only small changes in the Figure-of-Merit, also shown by the large width of the flexibility bars in Table \ref{binredshifttable}. |
This shows that these results are robust against moderate changes in the survey design. | This shows that these results are robust against moderate changes in the survey design. |
In the previous Section we focussed on optimizing the survey parameters to obtain the best constraints on the properties of dark energy. whilst. keeping the constraint parameters fixed. | In the previous Section we focussed on optimizing the survey parameters to obtain the best constraints on the properties of dark energy, whilst keeping the constraint parameters fixed. |
We now consider the case where the current constraint parameters are changed or new constraints are added. and the resulting elfects on the optimal survey FoM. and configuration. | We now consider the case where the current constraint parameters are changed or new constraints are added, and the resulting effects on the optimal survey FoM and configuration. |
The constraint parameters cannot. be considered: as simple survey design. parameters. as they are built in at an instrument level. | The constraint parameters cannot be considered as simple survey design parameters, as they are built in at an instrument level. |
Furthermore. the constraint parameters will be unbounded by FoM in one direction. | Furthermore, the constraint parameters will be unbounded by FoM in one direction. |
For exaniple a survey that runs for 5 vears will abwavs do better than a survey that lasts for only 3 vears. | For example a survey that runs for 5 years will always do better than a survey that lasts for only 3 years. |
However. the behaviour of the FoM will be different for different constraint parameters. | However, the behaviour of the FoM will be different for different constraint parameters. |
The Figure-of-Merit. of the best. survey will continue to scale as the total survey time is increased. but may quickly asvmptote to some maximal value in the case of the number of spectrograph fibres. | The Figure-of-Merit of the best survey will continue to scale as the total survey time is increased, but may quickly asymptote to some maximal value in the case of the number of spectrograph fibres. |
We consider three cases: (1) a extra constraint is imposed on the maximum area to be surveyed. motivated bv the size of the input. photometric catalogue: (2) the number of spectroscopic fibres is changed: and (3) the telescope aperture and field-of-view are changed. | We consider three cases: (1) a extra constraint is imposed on the maximum area to be surveyed, motivated by the size of the input photometric catalogue; (2) the number of spectroscopic fibres is changed; and (3) the telescope aperture and field-of-view are changed. |
Finally. we consider how the optimal design changes if constraints [rom other dark energy surveys (of both supernovae and baryon acoustic oscillations). which will have been completed: by the time that à WEMOS-like instrument is constructed. are included in the analysis. | Finally, we consider how the optimal design changes if constraints from other dark energy surveys (of both supernovae and baryon acoustic oscillations), which will have been completed by the time that a WFMOS-like instrument is constructed, are included in the analysis. |
The optimal surveys are those with the maxiniunm possible area. surveving thousands of square degrees. | The optimal surveys are those with the maximum possible area, surveying thousands of square degrees. |
Llowever. a spectroscopic survey requires an input catalogue of photometrically pre-selected galaxies. | However, a spectroscopic survey requires an input catalogue of photometrically pre-selected galaxies. |
Llow would the dark encreyv constraints be allectecl if the available area of the input catalogue is less than the optimal value? | How would the dark energy constraints be affected if the available area of the input catalogue is less than the optimal value? |
Looking at Table S [or a single redshift bin. we see that the Uexibility bounds place a lower limit on the total survey area of 3000 sq. | Looking at Table \ref{binredshifttable} for a single redshift bin, we see that the flexibility bounds place a lower limit on the total survey area of 3000 sq. |
degs.. | degs., |
compared to the optimal area of 6000 sq. | compared to the optimal area of 6000 sq. |
degs. | degs. |
Such imaging surveys. whilst not vet in existence. are in the planning stages (for example the Dark Energy Survey (DIES). see Abbott 2005). | Such imaging surveys, whilst not yet in existence, are in the planning stages (for example the Dark Energy Survey (DES), see Abbott 2005). |
In this Section we investigate the variation of the optimal l'igure-of-Merit. with the available number of spectrograph libres. assuming a single survey bin at low redshift. | In this Section we investigate the variation of the optimal Figure-of-Merit with the available number of spectrograph fibres, assuming a single survey bin at low redshift. |
. We consider both line and continuum-emission survey strategies. including both “pessimistic” and “optimistic” versions of the number counts caleulation. which correspond simply to decreasing and increasing the predicted. number of 2~1 ealaxies by 50%.. | We consider both line and continuum-emission survey strategies, including both “pessimistic” and “optimistic” versions of the number counts calculation, which correspond simply to decreasing and increasing the predicted number of $z \sim 1$ galaxies by . |
We assume that the fibres are able to | We assume that the fibres are able to |
Diffuse extragalactic background light (EBL) represents the sum total of all the photons produced by luminous matter over the lifetime of the Universe. | Diffuse extragalactic background light (EBL) represents the sum total of all the photons produced by luminous matter over the lifetime of the Universe. |
In the ultraviolet. optical. and near-infrared (IR) the EBL is directly attributable to star formation and active galactic nucleus (AGN) activity (forreviews.see222722.andreferences therein).. | In the ultraviolet, optical, and near-infrared (IR) the EBL is directly attributable to star formation and active galactic nucleus (AGN) activity \citep[for reviews, see][and references therein]{tyson1990,tyson1995,henry1991,henry1999,leinert1998}. |
However. in regions with significant dust opacity radiation generated by stars and AGN is reprocessed by dust into the IR as thermal emission (see.e.g.??).. leading to the expectation of a significant EBL component at A~10.— 1000/:m. Measurements of this Cosmic Diffuse IR Backround (CDIRB) were notoriously difficult to obtain. owing primarily to the presence of significant zodiacal and galactic foregrounds and the lack of access to this region of the spectrum from the ground. | However, in regions with significant dust opacity radiation generated by stars and AGN is reprocessed by dust into the IR as thermal emission \citep[see, e.g.][]{soifer1991,sanders1996}, leading to the expectation of a significant EBL component at $\lambda \sim 10-1000\mu$ m. Measurements of this Cosmic Diffuse IR Backround (CDIRB) were notoriously difficult to obtain, owing primarily to the presence of significant zodiacal and galactic foregrounds and the lack of access to this region of the spectrum from the ground. |
The launch of theExplorer (COBE:foranoverview.see?) finally revealed a CDIRB comparable in brightness to the optica EBL. providing a complete census of obscured star formation and AGN activity across cosmic time (222222). | The launch of the \citep[COBE: for an overview, see][]{boggess1992} finally revealed a CDIRB comparable in brightness to the optical EBL, providing a complete census of obscured star formation and AGN activity across cosmic time \citep{hauser1998,kelsall1998,aredt1998,dwek1998,fixsen1998,hauser2001}. |
Because the emission mechanisms that generate the CDIRB are intimately connected to galaxy formation and evolution. i provides a powerful observational constraint on models ο). | Because the emission mechanisms that generate the CDIRB are intimately connected to galaxy formation and evolution, it provides a powerful observational constraint on models \citep{partridge1967}. |
A particularly popular technique for modeling the CDIRB has been backwards evolution (e.g..22222222). in which a parameterized fi to the evolution of the IR luminosity function — using low-redshif | A particularly popular technique for modeling the CDIRB has been backwards evolution \citep[e.g.,][]{rowanrobinson2001,rowanrobinson2009,lagache2003,lagache2004,lagache2005,xu2003,franceschini2008,finke2009b}, in which a parameterized fit to the evolution of the IR luminosity function – using low-redshift |
Vaucouleurs et al. ( | Vaucouleurs et al. ( |
1991) and have μαστάος BYx12. | 1991) and have magnitudes $B^0_T < 12$. |
1 selected those galaxies which have been detected at both GO aud LOO 41 by TRAS (Iuapp et al. | I selected those galaxies which have been detected at both 60 and 100 $\mu$ m by IRAS (Knapp et al. |
1989). | 1989). |
IT = 50 iu Ape 1 is ο throughout this paper. | $_0$ = 50 Km $^{-1}$ $^{-1}$ is assumed throughout this paper. |
For each galaxy the FIR huuinositv. the flux ratio. the dust temperature and the dust mass frou FIR data (hereafter FIR dust mass) have been derived aud [listed in Table 1. | For each galaxy the FIR luminosity, the flux ratio, the dust temperature and the dust mass from FIR data (hereafter FIR dust mass) have been derived and listed in Table 1. |
The FIR fluxes by Knapp et al. ( | The FIR fluxes by Knapp et al. ( |
1989) have been corrected taking iuto accom the contribution of hot ciretuustellar dust (CJ95) aud have Όσσα used for the computation of: the FIR hunuinositv in the band 120 yan (Uelou ct al. | 1989) have been corrected taking into account the contribution of hot circumstellar dust (GJ95) and have been used for the computation of: the FIR luminosity in the band 40-120 $\mu$ m (Helou et al. |
1985). the flux ratio and the dust color temperature (Wenning et al. | 1985), the flux ratio and the dust color temperature (Henning et al. |
1990). | 1990). |
The uncertainties in the three quantities depend on the fux uncertainties. which are estimated to be in the range 1056-3054. (IKnapp et al. | The uncertainties in the three quantities depend on the flux uncertainties, which are estimated to be in the range $\%$ $\%$ (Knapp et al. |
1989). | 1989). |
The chemical aud physical properties of ie dust eramus play a critical rolle iu the evaluation of the uncertainties on temperature and mass. | The chemical and physical properties of the dust grains play a critical rôlle in the evaluation of the uncertainties on temperature and mass. |
On the other hand. the present paper is mainly addressed to ie dependeuce of the dust mass on the choice of a outieular temperature distribution. aud Twill not discuss 1e problem of the dust properties (which will be studied oeji a forthcoming oper). | On the other hand, the present paper is mainly addressed to the dependence of the dust mass on the choice of a particular temperature distribution, and I will not discuss the problem of the dust properties (which will be studied in a forthcoming paper). |
A spectral iudex a=1 for je dust is used in all the computations. | A spectral index $\alpha=1$ for the dust is used in all the computations. |
The FIR dust nasses listed in Table 1 are evaluated from the FIR fluxes Nw means of the single temperature and the temperature distribution uodels respectively. | The FIR dust masses listed in Table 1 are evaluated from the FIR fluxes by means of the single temperature and the temperature distribution models respectively. |
The dust niasses computed by adopting a siugle temperature a1 a temperature distribution are compared in Fie... | The dust masses computed by adopting a single temperature and a temperature distribution are compared in Fig.1. |
Both methods lead to a temperature dependence. as expected when the dust amount is derived by using thermal euission. | Both methods lead to a temperature dependence, as expected when the dust amount is derived by using thermal emission. |
Iu fact. the colder the dust erains. the larger should be their total uuuiber in order to produce a even FIR cunission. | In fact, the colder the dust grains, the larger should be their total number in order to produce a given FIR emission. |
The dust nasses obtained with a temperature distribution are larger than those cerived from a siugle-teiipperrattiurre model | The dust masses obtained with a temperature distribution are larger than those derived from a re model. |
This confinis that IRAS neasurenients allow to determine the warm (30-51 Is) dust amount onlv. while the contribution of the cold dust is neglected. | This confirms that IRAS measurements allow to determine the warm (30-54 K) dust amount only, while the contribution of the cold dust is neglected. |
Moreover. it tums out that. iu the oeseut sample. the FIR huuinositv does not depoeud ou the color temperature. lus confiiuuug that the FIR Cluission is a result of different contributions and caunot )e properly explained by a source with a snele equilibrium enmperature. | Moreover, it turns out that, in the present sample, the FIR luminosity does not depend on the color temperature, thus confirming that the FIR emission is a result of different contributions and cannot be properly explained by a source with a single equilibrium temperature. |
The two models give masses differing ly actors from 2 to 6. | The two models give masses differing by factors from 2 to 6. |
This wide range is due cither to he shape of the temperature distribution aud/or to the uncertainties in the dust parameters. | This wide range is due either to the shape of the temperature distribution and/or to the uncertainties in the dust parameters. |
Ou the other laud. he ratio between the two dust mass evaluations (Fig. | On the other hand, the ratio between the two dust mass evaluations (Fig. |
2) shows a general temperature dependence which suggests hat colder galaxies have a larger amount of missed dust. | 2) shows a general temperature dependence which suggests that colder galaxies have a larger amount of missed dust. |
= 110010) and the larduess of the radiation fick derived with IRS suggest that most of the mid-IR huninosity is produced by very voung massive star formation. | $\lesssim$ pc) and the hardness of the radiation field derived with IRS suggest that most of the mid-IR luminosity is produced by very young massive star formation. |
Following Righy&Ricke(2001). who studie the evolution of various mid-IR line ratios with time using the CLOUDY photoionization model. we fine that the [NOETI]/|NOeII].. | and SIV]/|STIT lonizing indices measured À3ITI|/LAxTI],in the 9980125 host are compatible with a starburst episode vounger than MN. | Following \citet{Rigby04b} who studied the evolution of various mid-IR line ratios with time using the CLOUDY photoionization model, we find that the [NeIII]/[NeII], [ArIII]/[ArII] and [SIV]/[SIII] ionizing indices measured in the 980425 host are compatible with a starburst episode younger than $\sim$ Myr. |
This is consistent with the estimate obtaies by Tanuneretal.(2006) based on the umber of WolfBavet stars observed in this region (i.c.. ADINIr). | This is consistent with the estimate obtained by \citet{Hammer06} based on the number of Wolf-Rayet stars observed in this region (i.e., Myr). |
The mid-IR cinissiou line diagnostics that cau be used to characterize the jiouizine conditions of star-fornuue environments are obviously not independent from cach other. | The mid-IR emission line diagnostics that can be used to characterize the ionizing conditions of star-forming environments are obviously not independent from each other. |
A ununber of correlations have ecu established between ΓΝΟΠΙΓ |AxETI|/[AxTI], aud STV]/|ST. as well as with ΝΟ.|SIV|/[NeII| and with optical line ratios like [OITII]A5Q007/|]OTII|A3727 (c.g.2009:Bernard-Salasetal. 2009). | A number of correlations have been established between [NeIII]/[NeII], [ArIII]/[ArII] and [SIV]/[SIII], as well as with [SIV]/[NeII] and with optical line ratios like $\lambda$ $\lambda$ 3727 \citep[e.g.,][]{Giveon02,Wu06,Dale06,Groves08,Gordon08,Hao09,Bernard09}. |
. The mid-IR spectra properties of the WR region are fully consistent witFPR these different relationships. as well as with the roug correlation that exists between the metallicity of galaxie: and the harduess of their radiation field (Wwotal.2006:Cordonetal. 2008). | The mid-IR spectral properties of the WR region are fully consistent with these different relationships, as well as with the rough correlation that exists between the metallicity of galaxies and the hardness of their radiation field \citep{Wu06,Gordon08}. |
. Moreover. the clectron density in the WR region cau be coustraimec frou the wo Sulfur III cussion lines detected with IRS. | Moreover, the electron density in the WR region can be constrained from the two Sulfur III emission lines detected with IRS. |
The flux ratio Als. A33.53 ~ 00.8. corresponds o AN.~ 200c02 (Civeonetal.20023. | The flux ratio $\lambda$ $\lambda$ $\sim$ 0.8 corresponds to $N_e\sim200$ $^{-3}$ \citep{Giveon02}. |
This ds slightly lareer than the estimate that was derived from he optical spectimm of the WR region using the SITI|AGT16A /|SITIJAGT31A. fux ratio OM.=158 7. Ihuunnmeretal. 2006)) and it is also higher than he densities measured in the ΤΠ regions of 1101 (SIH|AT8.71/|SITIJA33.53 Cordon.etal. 2008)). | This is slightly larger than the estimate that was derived from the optical spectrum of the WR region using the $\lambda$ $\lambda$ flux ratio $N_e = 158$ $^{-3}$, \citealt{Hammer06}) ) and it is also higher than the densities measured in the HII regions of 101 $\lambda$ $\lambda$ $\sim$ 0.2, \citealt{Gordon08}) ). |
This could be related to ~the 00.2.effect of the stellar winds xoduced. by the WolfRavet populations. which iav iive ciiciently compressed the eas aud the ISM of the WR reeion in the host of 9980125. | This could be related to the effect of the stellar winds produced by the Wolf-Rayet populations, which may have efficiently compressed the gas and the ISM of the WR region in the host of 980425. |
Finally. the equivalent widths of the PAID features characterizing the WR region fall within expectations eiven the metallicity aud the harduess of the radiation field previously measured. | Finally, the equivalent widths of the PAH features characterizing the WR region fall within expectations given the metallicity and the hardness of the radiation field previously measured. |
In star-forming galaxies the relative streneth of PAIIS over the VSG continua is known to remain coustant over a wide range of environments, but it rapidly decreases with harder lonizing radiations. | In star-forming galaxies the relative strength of PAHs over the VSG continuum is known to remain constant over a wide range of environments, but it rapidly decreases with harder ionizing radiations. |
This is particularly apparent when the harduess of the radiation field reaches a threshold of 3 (Gordonetal.2008:Lebouteiller 1b).. ~or NOI]11similarly when the ebhenmical abundance gets lower than - llog|O/II|SN.8.3 (ee...Ποιοςeeclbrachtetal. 2008). | This is particularly apparent when the hardness of the radiation field reaches a threshold of $\sim$ 1–3 \citep{Gordon08,Lebouteiller11}, or similarly when the chemical abundance gets lower than $\sim$ 8–8.3 \citep[e.g.,][]{Houck04b,Engelbracht05,Wu06,Engelbracht08}. |
. The properties of the WR region in the 9980125. host ealaxy correspond to this reehme of transition between extreme metal-poor regionis and euvironments with solar metallicity. aud the PATIS aro clearly detected albeit with smaller EWs than observed in sources with softer radiation fields. | The properties of the WR region in the 980425 host galaxy correspond to this regime of transition between extreme metal-poor regions and environments with solar metallicity, and the PAHs are clearly detected albeit with smaller EWs than observed in sources with softer radiation fields. |
Looking at retfig:plot pee... wealsonotethatthe’. PAIL inter-baud ratio ds substautially smaller than observed in the mid-IR spectimm of the prototypical starburst ealaxy T7711 but it is once again in agreement with the trend that has been secu between this quantity and the ionizius index measured with the [NeITI]/[NOIH] line ratio (Simithctal.2007). | Looking at \\ref{fig:plot_spec}, we also note that the PAH inter-band ratio is substantially smaller than observed in the mid-IR spectrum of the prototypical starburst galaxy 7714 but it is once again in agreement with the trend that has been seen between this quantity and the ionizing index measured with the [NeIII]/[NeII] line ratio \citep{Smith07}. |
.. Given that the feature can be produced by larger PAIIs than the ones responsible for the enidssion (Drame&Li2007). the WR regiou could thus be experienciug a selective destruction of the small PATIs that cut at short mid-IR wavelengths (Smithetal.2007). | Given that the feature can be produced by larger PAHs than the ones responsible for the emission \citep{Draine07}, the WR region could thus be experiencing a selective destruction of the small PAHs that emit at short mid-IR wavelengths \citep{Smith07}. |
(Fem The total IR huuuositv derived in Lt (log |Lin/L.]e 88.7) places tιο WR region at the very bright end of the Iuninuosits> function of WIL regions observed in the local Universe (Bradleyetal.2006:Leeetal.2011). | 0.7cm The total IR luminosity derived in 4 (log $_{\rm IR}$ $_{\odot}$ $\sim$ 8.7) places the WR region at the very bright end of the luminosity function of HII regions observed in the local Universe \citep{Bradley06,Lee11}. |
. Cousicdering the pliaysical size of the optical light distribution measured in the IIST image of the host (©LOOpe.Fruboetal.2000:Tlamumer2006).. this environment is also one of the sinele isolated ΠΠ regions with the highest star formation deusitv identified so far. | Considering the physical size of the optical light distribution measured in the HST image of the host \citep[$\lesssim$\,100\,pc, ][]{Fynbo00,Hammer06}, this environment is also one of the single isolated HII regions with the highest star formation density identified so far. |
For mstauce. the brightest star-forming conrplex iu the Milky Wav. WI9A. releases up to Lig~2.7«10 LEO within a size less than ~50ppe in diameter AVard-Thompson&Robson19903.. while the total huuinositv of the DDoradus nebula in the Large Maeclanic Cloud oulv reaches Ligτν10 LL. (Werneretal1978). | For instance, the brightest star-forming complex in the Milky Way, W49A, releases up to $_{\rm IR} \sim 2.7\times 10^7$ $_{\odot}$ within a size less than $\sim$ pc in diameter \citep{Ward_Thompson90}, while the total luminosity of the Doradus nebula in the Large Magelanic Cloud only reaches $_{\rm IR} \sim 4\times 10^7$ $_{\odot}$ \citep{Werner78}. |
Yet. some of the intrinsic characteristics of this environment have already been observed in other extragalactic sources. which suggests that such complex of star formation does uot represent a unique object in the local Universe. | Yet, some of the intrinsic characteristics of this environment have already been observed in other extragalactic sources, which suggests that such complex of star formation does not represent a unique object in the local Universe. |
For instance. isolated IIT regious with sinibku οΕν spectral slope and ionization iudices have been found in the eiaut spiral ealaxy AIILOL (Cordonetal.2008). | For instance, isolated HII regions with similar mid-IR spectral slope and ionization indices have been found in the giant spiral galaxy 101 \citep{Gordon08}. |
Amone these regions. 55161 does contain a population of WolfRavet stars (Schaereretal.1999) and it has a Hux of muuJy. which given the distance of LLOL MMpe) corresponds to a mid-IR huninosity lareer than what we measured iu the host of 9980125. | Among these regions, 5461 does contain a population of Wolf-Rayet stars \citep{Schaerer99} and it has a flux of $\sim$ mJy, which given the distance of 101 Mpc) corresponds to a mid-IR luminosity larger than what we measured in the host of 980425. |
The mid-IR properties of the WR region also vescuuble what is seen in Blue Compact Dwart (BCDs) galaxies with comparable oxvecn abundance. which often exhibit a steeply rising hot dust contimmiun with prominent ionic lines aud very little absorption bv the silicate features (Wiretal.2006). | The mid-IR properties of the WR region also resemble what is seen in Blue Compact Dwarf (BCDs) galaxies with comparable oxygen abundance, which often exhibit a steeply rising hot dust continuum with prominent ionic lines and very little absorption by the silicate features \citep{Wu06}. |
. Similarly. some of the eiaut. Super Star Clusters in the overlap region of the Antennae exhibit comparable or steeper mid-IR confini cluission as well as higher star formation rates, although these clusters iav be spatially more extended and their radiation field is typically a bit softer than observed in the WR region (Brandletal.2009). | Similarly, some of the giant Super Star Clusters in the overlap region of the Antennae exhibit comparable or steeper mid-IR continuum emission as well as higher star formation rates, although these clusters may be spatially more extended and their radiation field is typically a bit softer than observed in the WR region \citep{Brandl09}. |
. What is nonetheless very unusual in the host of 9980125 is the large fraction (15410 This characteristic of the WR region ids actually consistent with what can be observed in recent starbursts. despite the appareutlv modest extinction derived. froii the silicate features aud the Uvdrogen Baler decremeut (see rop)) | What is nonetheless very unusual in the host of 980425 is the large fraction $\pm$ This characteristic of the WR region is actually consistent with what can be observed in recent starbursts, despite the apparently modest extinction derived from the silicate features and the Hydrogen Balmer decrement (see \\ref{sec:irs_prop}) ). |
La faet. cergugounganddust embeddedsteHarclasterswith muchhighereatinetiouleg. | In fact, very young and dust-embedded stellar clusters with much higher extinction (e.g., |
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