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High-S/N spectra are used for a complete spectral fit using a power law model. | High-S/N spectra are used for a complete spectral fit using a power law model. |
A detailed description of the data analysis pipeline is giver in Appendix AppendixA:. | A detailed description of the data analysis pipeline is given in Appendix \ref{Crab_pipeline}. |
. The complete pipeline used for our study of GRO J1655-40 consists of steps 1. 2. 4. and 5 described there. | The complete pipeline used for our study of GRO J1655-40 consists of steps 1, 2, 4, and 5 described there. |
RXTE monitored the whole outburst of GRO J1655-40 since its discovery (?). | RXTE monitored the whole outburst of GRO J1655-40 since its discovery \citep{Markwardt_2005_GRO}. |
The dataset is composed of 490 observations. performed between 2005-02-26 and 2005-11-11. | The dataset is composed of 490 observations, performed between 2005-02-26 and 2005-11-11. |
Each observation has a typical integration time of ~ 1.5 ks. for a total observing time of ~664 ks. | Each observation has a typical integration time of $\sim$ 1.5 ks, for a total observing time of $\sim$ 664 ks. |
Spectral data extracted from the complete dataset have been kindly made available to the community by the MIT group .. Spectra for both source and background as well as response | Spectral data extracted from the complete dataset have been kindly made available to the community by the MIT group Spectra for both source and background as well as response |
There is currently a growing interest in the acceleration of non-thermal particles at highly relativistic shocks. | There is currently a growing interest in the acceleration of non–thermal particles at highly relativistic shocks. |
There are three classes of velativistic sources: bevoud the well-established extra-Galactic (Blazars) aud Galactic (superluminal) sources. both of which exhibit superluninal notions. it is uow also well-established that Gamuna Ray Bursts (GRBs) display highly relativistic expausious. with Lorentz [actors well in excess of 100. | There are three classes of relativistic sources: beyond the well–established extra–Galactic (Blazars) and Galactic (superluminal) sources, both of which exhibit superluminal motions, it is now also well-established that Gamma Ray Bursts (GRBs) display highly relativistic expansions, with Lorentz factors well in excess of $100$. |
Other classes of relativistic sources may include Soft Gamuna Ray Repeaters (SCRs). whose recurrent explosions are largely super—Edclinetou. aud special SNe similar to SN 1998bw. which displayed marginally Newtouiau expausion 1) when optical emission lines became detectable. about a mouth after the explosion. | Other classes of relativistic sources may include Soft Gamma Ray Repeaters (SGRs), whose recurrent explosions are largely super–Eddington, and special SNe similar to SN 1998bw, which displayed marginally Newtonian expansion $\approx 6\times 10^4\; km\; s^{-1}$ ) when optical emission lines became detectable, about a month after the explosion. |
With the ciscovery of GRBs’ alterglows. it has now become feasible to derive the euergy spectral iudex & of electrous accelerated. at the forward shock. as a function of the varying (decreasing) shock Lorentz factor 5. provided simultaneous wide-band spectral coverage is available. | With the discovery of GRBs' afterglows, it has now become feasible to derive the energy spectral index $k$ of electrons accelerated at the forward shock, as a function of the varying (decreasing) shock Lorentz factor $\gamma$, provided simultaneous wide–band spectral coverage is available. |
With the launch of the USA/Italy/Uls mission SWIFT. these data will become available for a statistically significant uumber of bursts. testing directly inodels for particle acceleration at relativistic shocks. | With the launch of the USA/Italy/UK mission SWIFT, these data will become available for a statistically significant number of bursts, testing directly models for particle acceleration at relativistic shocks. |
Furthermore. since GRBs must also clearly accelerate protous. the same tides A& may cletermine the spectrum of ultra high euerey cosmic rays observed at Earth. | Furthermore, since GRBs must also clearly accelerate protons, the same index $k$ may determine the spectrum of ultra high energy cosmic rays observed at Earth. |
However. util recently. both the lack of astrophysical motivation aud the difficulty inherent in treating highly anisotropic distribution fuucetions have stiffened research on this topic. | However, until recently, both the lack of astrophysical motivation and the difficulty inherent in treating highly anisotropic distribution functions have stiffened research on this topic. |
Early | Early |
class have been discovered receutlv (see Matt 1997. for a review]. | class have been discovered recently (see Matt \cite{matt_b} for a review). |
As anticipated above. Sy2s in carly spectroscopic studies were mostly selected. frou former allsky Xταν surveys. henee these studies were generally biased for Xorav bright objects. | As anticipated above, Sy2s in early spectroscopic studies were mostly selected from former all–sky X–ray surveys, hence these studies were generally biased for X–ray bright objects. |
Very ikelv. this selectiou criterion resulted iu a bias in favor of low Nyy Sv2s. | Very likely, this selection criterion resulted in a bias in favor of low $_H$ Sy2s. |
Later on. hard Xray spectroscopic studies (nostlv bv means of ASC'A) probed fainter samples of ACNS (οιο, Turner et al. 1997a)). | Later on, hard X–ray spectroscopic studies (mostly by means of ASCA) probed fainter samples of AGNs (e.g. Turner et al. \cite{turner_a}) ). |
IIowever. many of he Sy2s observed by ASCA were selected amongst sources known to show broad liepA in polarized Light (Awaki ct al. 1997)). | However, many of the Sy2s observed by ASCA were selected amongst sources known to show broad lines in polarized light (Awaki et al. \cite{awaki}) ). |
This selection criterion nüehlt introduce a bias for low Ny as well. | This selection criterion might introduce a bias for low $_H$ as well. |
ITeisler et al. (19973) | Heisler et al. \cite{heisler}) ) |
showed that the detectability of polarized broad lines is related to the obscuration of the nuclear reeglon. | showed that the detectability of polarized broad lines is related to the obscuration of the nuclear region. |
Smunrizius. former Nrav spectroscopic surveys were seriously biased against heavily obscured Syv2s aud. therefore. they are not suitable to study the real distribution of the absorbing column deusities Nyy. The knowledge of the distribution of Ny iu Sv2s is Huportant to understand the nature of their obscuring medium. Which has implications for the unified model. | Summarizing, former X–ray spectroscopic surveys were seriously biased against heavily obscured Sy2s and, therefore, they are not suitable to study the real distribution of the absorbing column densities $_H$ The knowledge of the distribution of $_H$ in Sy2s is important to understand the nature of their obscuring medium, which has implications for the unified model. |
Also. the distribution of Nyy is relevaut to the svuthesis of the Xrav background. | Also, the distribution of $_H$ is relevant to the synthesis of the X–ray background. |
Indeed. obscured ACNs are thought to contribute to most of the high energy (> 2 keV) extragalactic backeround (Comastri et al. 1995.. | Indeed, obscured AGNs are thought to contribute to most of the high energy $>$ 2 keV) extragalactic background (Comastri et al. \cite{comastri}, |
Madau et al. 1005). | Madau et al. \cite{madau}) ). |
We have undertaken a program of observations with BeppoSAX. the ItalianDutch Xray satellite. aimed at studving the hard X.rav properties of weak Sy2s aud at assessing the vreal” distribution of their absorbing coluun densities. | We have undertaken a program of observations with BeppoSAX, the Italian–Dutch X–ray satellite, aimed at studying the hard X–ray properties of weak Sy2s and at assessing the “real” distribution of their absorbing column densities. |
As described in the next section. BeppoSAN is au excellent ool to pursue this goal. since it combines high seusitivitv (required. to observe weak ACNs) aud a wide spectral coverage (0.1.300 keV. required to identify aud discutanele various spectral componcuts). | As described in the next section, BeppoSAX is an excellent tool to pursue this goal, since it combines high sensitivity (required to observe weak AGNs) and a wide spectral coverage (0.1–300 keV, required to identify and disentangle various spectral components). |
Sv2s suitable for this study were drawn out of Maioliuo Ricke’s (1995)) sample. | Sy2s suitable for this study were drawn out of Maiolino Rieke's \cite{maiolino_a}) ) sample. |
This sample is extracted from the Revised Shapleyv-Aines (RSA) catalog of galaxies (that is limited to uaeuitude Br<13.2. Saudage Tanuuaun LOST)). and Sevfert ealaxies are selected according to their optica lines. | This sample is extracted from the Revised Shapley-Ames (RSA) catalog of galaxies (that is limited to magnitude $B_T < 13.2$, Sandage Tammann \cite{sandage}) ), and Seyfert galaxies are selected according to their optical lines. |
As discussed in Maiolino Ricke. this sample is much less biased than others. both iu terms of huninosity of the Sevtert nuclei aud in terms of properties of their hos ealaxies. | As discussed in Maiolino Rieke, this sample is much less biased than others, both in terms of luminosity of the Seyfert nuclei and in terms of properties of their host galaxies. |
Out of the 51 Sy2s in the Maioliuo Ricke sample 22 have already been observed by. ASCA. | Out of the 54 Sy2s in the Maiolino Rieke sample 22 have already been observed by ASCA. |
We selected 8 of the remaining 32 Sv2s. based on their (arrow} line flux: fo naxiuze the chauces of detection we chose the sources showing the highest |OITI| flux. | We selected 8 of the remaining 32 Sy2s, based on their (narrow) line flux: to maximize the chances of detection we chose the sources showing the highest [OIII] flux. |
Maiolino Ricke also show that. although the [OITI] Ine is e1uitfted on scales much larger than the putative oc-eale torus. the [OITI] is not a completely isotropic indicator of the nuclear huuünositw. for the host ealaxy disk might obscure wart of the NLR. | Maiolino Rieke also show that, although the [OIII] line is emitted on scales much larger than the putative pc-scale torus, the [OIII] is not a completely isotropic indicator of the nuclear luminosity, for the host galaxy disk might obscure part of the NLR. |
Iowever. once the ΟΠΗ flix is corrected for the extinction. deduced from he Baluer decrement. it should provide an indication of he nuclear activity that is independent of the pc-scale obscuration due to the torus. | However, once the [OIII] flux is corrected for the extinction deduced from the Balmer decrement, it should provide an indication of the nuclear activity that is independent of the pc-scale obscuration due to the torus. |
As a consequence. although our [OTH)|-based selection criterion might introduce a bias in our sauple forinfrinsicallg |uninous sources. it avoids biases against liehly obscured Sy2 uuclei. thus overcoming limitations of former survevs. | As a consequence, although our [OIII]-based selection criterion might introduce a bias in our sample for luminous sources, it avoids biases against highly obscured Sy2 nuclei, thus overcoming limitations of former surveys. |
Preliminary results of this survey were published iu Salvati et al. (1997... 10051). | Preliminary results of this survey were published in Salvati et al. \cite{salvati}, \cite{salvati2}) ). |
In this paper we report aud discuss results of BeppoSAN observations for all of the 8 Seyfert 2s in our [OTT] selected sample. | In this paper we report and discuss results of BeppoSAX observations for all of the 8 Seyfert 2s in our [OIII] selected sample. |
A more thorough statistical analysis. obtained by mereine our BeppoSAX data with data in the literature. is preseuted iu Bassani et al. ( | A more thorough statistical analysis, obtained by merging our BeppoSAX data with data in the literature, is presented in Bassani et al. ( |
in prep.). | in prep.). |
A description of the BeppoSAN observatory is eiveu πι Boclla et al (1997a)). | A description of the BeppoSAX observatory is given in Boella et al. \cite{boella_a}) ). |
The pavload iustruineuts include four coaligned narrow field iustrunents: a Low Encrev Concentrator Spectrometer (LECS. Parmar ct al. 1997). | The payload instruments include four co–aligned narrow field instruments: a Low Energy Concentrator Spectrometer (LECS, Parmar et al. \cite{parmar}) ), |
three Medium Eucrex Concentrator Spectrometers (MECS. Boclla ot al. 199753). | three Medium Energy Concentrator Spectrometers (MECS, Boella et al. \cite{boella_b}) ), |
a Teh Pressure Cas Scintillation Proportional Counter (ΠΡΟ, Alanzo et al 1997)) | a High Pressure Gas Scintillation Proportional Counter (HPGSPC, Manzo et al. \cite{manzo}) ) |
and ai Phoswich Detector System (PDS. Froutera ct al 1997) | and a Phoswich Detector System (PDS, Frontera et al. \cite{frontera}) ). |
Both LECS aud MECS spectrometers have naiagiue capabilities (angular resolution ~ 1.2 arcnun. FWOAD and cover the 0.110 keV and 1.510 keV spectral bands. respecively (in the overlapping spectral region the MECS are three times more scusitive thau LECS). | Both LECS and MECS spectrometers have imaging capabilities (angular resolution $\sim$ 1.2 arcmin, FWHM) and cover the 0.1–10 keV and 1.5–10 keV spectral bands, respectively (in the overlapping spectral region the MECS are three times more sensitive than LECS). |
Their euergv resolution is about at 6 keV. IIPGSPC and PDS operate iu the 1120 keV aud 15300 keV spectral bands. respectively. | Their energy resolution is about at 6 keV. HPGSPC and PDS operate in the 4–120 keV and 15–300 keV spectral bands, respectively. |
Iu the overlapping region the PDS is more sensitive (bv a factor of 1.8). while the WPGSPC has superior energy resolution. | In the overlapping region the PDS is more sensitive (by a factor of 4–8), while the HPGSPC has superior energy resolution. |
Table 1 lists the sources observed iu our program so far. along with the onsource total integration time and net count rate (1.0. background subtracted) for each iustruinent. | Table \ref{tab_obs} lists the sources observed in our program so far, along with the on–source total integration time and net count rate (i.e. background subtracted) for each instrument. |
One of the MECS units stopped working iu May 1997. as a consequeuce ALCC-05-L8-002 was observed with two MECS units ouly. | One of the MECS units stopped working in May 1997, as a consequence MCG-05-18-002 was observed with two MECS units only. |
For all sources. but NGC 5613. LECS and NECS spectra were extracted from an aperture of δ aud Ll’ in radius. since these apertures were found to optimize the signaltonoise ratios for these faint sources. | For all sources, but NGC 5643, LECS and MECS spectra were extracted from an aperture of $'$ and $'$ in radius, since these apertures were found to optimize the signal–to–noise ratios for these faint sources. |
For NGC 5613 we chose au extraction radius of 1. since outside of this radius the observed fiw is affected by the emission of a nearby galaxy cluster (see Sect. L1)). | For NGC 5643 we chose an extraction radius of $'$, since outside of this radius the observed flux is affected by the emission of a nearby galaxy cluster (see Sect.\ref{results}) ). |
Another exception is NGC 1386. where we used for the LECS the same extraction aperture used for the MECS (i0. 0 iu radiux) | Another exception is NGC 1386, where we used for the LECS the same extraction aperture used for the MECS (i.e. $'$ in radius) |
The study of stellay populations of non-resolved svstcuis has greatly relied the models derived from evolutionlE populations svuthesisou technique. | The study of stellar populations of non-resolved systems has greatly relied on the models derived from evolutionary populations synthesis technique. |
This approach is based on the spectrophometric properties of stars at. ideally. all evolutionary phases and takes iuto acconnt all phenomena that argely affect the evolution of a star (e.g. duass-loss). | This approach is based on the spectrophometric properties of stars at, ideally, all evolutionary phases and takes into account all phenomena that largely affect the evolution of a star (e.g., mass-loss). |
Over the wears. ever since the technique was first ∖⊳−implemented∖∖∖ (Tinsley↴⋅↴∖⇁↽≺⋅∖∖↽∢↼⊲≻1968.1972). a Mwide variety of nmodels based‘l ou differcutnveut uieredieutsBout ilave ;been constructed iand used .in the study of vounge2 and old. stellar populations. | Over the years, ever since the technique was first implemented \citep{tinsley68,tinsley72}, a wide variety of models based on different ingredients have been constructed and used in the study of young and old stellar populations. |
2005).. Not surprisingly. most of the work done up to date has vastly focused in the optical spectrophotomoetrie properties of stellar svsteius. aud until relatively recently it has expanded {ο other wavelengths (as far as the detailed analysis of spectral features ix concerned). and. ir sone cases. jucluded the effects of an interstellar medium (Silvactal,1908:Pauuzzoetal. 2005). | Not surprisingly, most of the work done up to date has vastly focused in the optical spectrophotometric properties of stellar systems, and until relatively recently it has expanded to other wavelengths (as far as the detailed analysis of spectral features is concerned), and, in some cases, included the effects of an interstellar medium \citep{silva98,panuzzo05}. |
At ultraviolet (UV) waveleneths. usnally divided into two segments, the fa-UV. (2002000 Aj) and the imid-UV (20003200 A3). the natural systems to look at are those whose underlying populations copiouslv eiut aud have their enüssion maxima du ⋅D . . _ ↴∖↴⋅↖↽↴ | At ultraviolet (UV) wavelengths, usually divided into two segments, the far-UV (1200–2000 ) and the mid-UV (2000–3200 ), the natural systems to look at are those whose underlying populations copiously emit and have their emission maxima in that window, i.e. star-forming systems. |
∖↴↑↸∖⋯↴∖↴⋜⋯∖↸∖⊼⊓⋅↸∖⋯↸∖↕⋅↖↽↕∐∏⋯↥⋅⋜⋯↕↕⊔⊔⋜⋯⋅↖↽⋜↧↴∖↴⊓⋅≺∏≻↕⋅↖↽↴∖↴↕↸⊳⋜↧↕ contexts (see.οι,Duzzoni2002).. if was eventually realized that also old aud intermediate age populations. which will be the main subject in this paper. deserve attention bv their own right. | While these systems are extremely important in many astrophysical contexts \citep[see, e.g.,][]{buz02}, it was eventually realized that also old and intermediate age populations, which will be the main subject in this paper, deserve attention by their own right. |
As au cxample we cin mention the countless studies motivated by the anexpected finding of a prominent far-UV flux excess iu the bulge of Andromeda (Code1969). | As an example we can mention the countless studies motivated by the unexpected finding of a prominent far-UV flux excess in the bulge of Andromeda \citep{code69}. |
. Aside of this ‘ar-UV fux excess. the mid-UWV still remains vastly quexylored. iu spite of the early suggestions that this wavelength region cau help iu lifting the so-called age-netallicity degeneracy (AMID) that plagues the optical spectrophotometzic⋅ properties⋅ of evolved populations. and that preveuts the uuivocalR determinationR: of: these xuiueters (Worthev-:1991:Dormanetal.⊀∙2003). | Aside of this far-UV flux excess, the mid-UV still remains vastly unexplored, in spite of the early suggestions that this wavelength region can help in lifting the so-called age-metallicity degeneracy (AMD) that plagues the optical spectrophotometric properties of evolved populations and that prevents the univocal determination of these parameters \citep{worthey94,dorman03}. |
Discutangling the effects of age and chemical compositioul is particularly imiportaut when attempting to evaluate he characteristics of distant red objects for which. hrough optical observations only feasible with the current seneration of large telescopes. we can only access the rest-frame mid-UV flux | Disentangling the effects of age and chemical composition is particularly important when attempting to evaluate the characteristics of distant red objects for which, through optical observations only feasible with the current generation of large telescopes, we can only access the rest-frame mid-UV flux \citep[e.g.,][]{dunlop96}. . |
these two lauits. we can find the reκ depeudeuce of the parameters via the requiremeut (Lujxps | these two limits, we can find the redshift dependence of the parameters via the requirement $\avg{L n} \propto \csfr$. |
The vedshift historv of cosuic-star formation is well-known (IIoxiuchietal.2009).. aud we encode this in the diinensouless "shape function Sz) (0). | The redshift history of cosmic-star formation is well-known \citep{horiuchi}, and we encode this in the dimensionless “shape” function S(z) (0) . |
Wο then have L.GO/L.()= 5S(:)inu the case of pure hinunosity evolution. aud (2)fr.(0)=5(:) in the case of pure density evolution. | We then have $\tracer_*(z)/\tracer_*(0) = S(z)$ in the case of pure luminosity evolution, and $n_*(z)/n_*(0) = S(z)$ in the case of pure density evolution. |
At a given redshift. 0) | At a given redshift, eq. \ref{eq:mavg}) ) |
gives the scaling (Mj=GUηρεxCMascnifS(:i). | gives the scaling $\avg{\mgas} = \avg{\mgas \psi n}/\csfr \propto \avg{\mgas \psi n}/S(z)$ . |
For our luminosity function and Keunicutt-Scliuidt melation we fud "a local --value of (Mai.9=68s10?A.- | For our luminosity function and Kennicutt-Schmidt relation we find a local value of $\avg{\mgas}_{z=0} =6.8 \times 10^{9} \ \msol$. |
At ο we have CMconjx(112)SEAD). alc OUSx(L112)8£262). | At other redshifts we have $\avg{\mgas \psi n} \propto (1+z)^{-\beta} \tracer_*(z)^\omega S(z)$, and so $\avg{\mgas} \propto (1+z)^{-\beta} \tracer_*(z)^\omega$. |
Thus for the pire inuinositv case L.XS(2). we find that the gas mass stronely evolves as Lea.)WwOL[τσι9. in response o the stronely changing SER. | Thus for the pure luminosity case $\tracer_* \propto S(z)$, we find that the gas mass strongly evolves as $\avg{\mgas} \propto (1+z)^{-\beta} S(z)^\omega$, in response to the strongly changing SFR. |
Consequeuthy. the factor of 1t )rise nmi cosmic star-formation at +~| implics a net HHennma-rav Duiudnositv merease of a factor 230. | Consequently, the factor of $10$ rise in cosmic star-formation at $z \simeq 1$ implies a net gamma-ray luminosity increase of a factor $\simeq 30$. |
Ou the other haud. iu the pure density evolution case. galaxy SERs are coustaut. L(2)=const. so that the mean gas luas8 GMthe,(11:) actually With redshift. while colmoving uunber2 of star-forudus ealaxics Πιοοσακος. but the net enhancement at high redshift is SETHer than in the pure lhuninosity evolution case. | On the other hand, in the pure density evolution case, galaxy SFRs are constant, $\tracer_*(z) = const$, so that the mean gas mass $\avg{\mgas} \propto (1+z)^{-\beta}$ actually with redshift, while the comoving number of star-forming galaxies increases, but the net enhancement at high redshift is smaller than in the pure luminosity evolution case. |
This kev difference leads to the factor ~1 between the ECB predictions for the pure censity aud pure luuinosity evolution cases seen in Figure 1. | This key difference leads to the factor $\sim 4$ between the EGB predictions for the pure density and pure luminosity evolution cases seen in Figure 1. |
Our full numerical calenlation uses a Milly Way pionic source spectzni Whoseshape is derived frou Ποιο.&Eubliu (2001)... calibrated to observations by normalizing the 2100 MeV photon enuüssion por hydrogen atom to the result at iuteriuediate Caletie latitudes (Abdoetal.20095)... | Our full numerical calculation uses a Milky Way pionic source spectrum whose is derived from \citet{pfrommer}, , calibrated to observations by normalizing the $> 100$ MeV photon emission per hydrogen atom to the result at intermediate Galactic latitudes \citep{FermiMW09}. |
The coziuic SER is from ITorxiucliietal.(2009).. | The cosmic SFR is from \citet{horiuchi}. |
For theMilkv Way SER. used to normalize the cosmic-ray flux/SER ratio. we use the recent estimate of Robitaille&Whitney(2010) (UMw=LAL./vr. a factor of 3 lower than earlier work). | For the Milky Way SFR, used to normalize the cosmic-ray flux/SFR ratio, we use the recent estimate of \citet{robitaille} $\psi_{\rm MW} = 1 {M_\odot/\rm yr}$, a factor of 3 lower than earlier work). |
Figure 1 shows our results for fhe normal galaxy contribution to the ECD. | Figure 1 shows our results for the normal galaxy contribution to the EGB. |
We pkt predictions for the Iuuitiug cases of pure ήπιοsity aud of pure density evolutiou. | We plot predictions for the limiting cases of pure luminosity and of pure density evolution. |
The uucertaiulos 111 the model inputs. Slmed in quadrature. propagate into the displaved ene baud that applies fo each curve. which we estimate to 2be a factor of 1000, resulting from uncertainties of: iu pionic enüissivitv (Abdoetal.20095)... in t16 normalization of the Galactic star-formation rate (Rovitaille&Whitney2010).. 1i the cosmic star-Orlation rates (IIoxinchietal.20Oy. aud in he huninosity scaling in eq. (2)). | The uncertainties in the model inputs, summed in quadrature, propagate into the displayed error band that applies to each curve, which we estimate to be a factor of $10^{\pm 0.3}$, resulting from uncertainties of: in pionic emissivity \citep{FermiMW09}, in the normalization of the Galactic star-formation rate \citep{robitaille}, in the cosmic star-formation rates \citep{horiuchi}, and in the luminosity scaling in eq. \ref{eq:mavg}) ). |
The true svsteinatic TLCvtadntv would also reflect the icealizations dn our nocel Giniversal cosmic-ray spectra and confinement). | The true systematic uncertainty would also reflect the idealizations in our model (universal cosmic-ray spectra and confinement). |
These errors are hard to estimate but iu any case ΠΠ] hat the uncertainty ranee in Figure Lis a lower bound otje error budget. | These errors are hard to estimate but in any case imply that the uncertainty range in Figure 1 is a lower bound to the error budget. |
Within errors. our predictions for both Παπάας mocels all at or below the level of the data. where the data ποσα to support the pure Iuninositv evolution case that explains uecarly the cutire signal. | Within errors, our predictions for both limiting models fall at or below the level of the data, where the data seem to support the pure luminosity evolution case that explains nearly the entire signal. |
Comparing ceutral values. this model gives z50% of theFermi EGD <10 GeV. Thus. unresolved normal galaxies make a substantial aud likely doniuaut contribution to the observed EGD. without overpredicting the signal. | Comparing central values, this model gives $\approx 50\%$ of the EGB $\la 10$ GeV. Thus, unresolved normal galaxies make a substantial and likely dominant contribution to the observed EGB, without overpredicting the signal. |
Eveu the pure density evolution case accounts for a minima of of the EGB around 0.3 GeV: this provides alower liuüt to the normalealaxy signal. | Even the pure density evolution case accounts for a minimum of of the EGB around 0.3 GeV; this provides a limit to the normal-galaxy signal. |
Thus. anvother EGB sources (Stecker&Salamou1996:Dennuer2007:Abdoctal.2010d:Faucher-Cagnere&Loeb2010) ust contribute no more than the remaining SO of the data. | Thus, any EGB sources \citep{ss96,dermer07,fermicounts,mspulsars}
must contribute no more than the remaining $80\%$ of the data. |
Tudeed. the LAT teamupper lanit to the blazar ECB contribution shown in Figure 1 is comparable to our lower limit (Abdoetal.2010d). | Indeed, the LAT team limit to the blazar EGB contribution shown in Figure 1 is comparable to our limit \citep{fermicounts}. |
. The spectral shapes of the two limiting cases are very similar: the peak in E772/dE lies at ~0.3 GeV because the bulk of the signal comes from 2~1. | The spectral shapes of the two limiting cases are very similar: the peak in $\eobs^2 dI/d\eobs$ lies at $\sim 0.3$ GeV because the bulk of the signal comes from $z \sim 1$. |
These models predict that the EGB turus over for Ezx GeV. a testable prediction of our model. | These models predict that the EGB turns over for $\eobs \la 0.3$ GeV, a testable prediction of our model. |
For hadronic cluission. the ligh-cnerey spectral iudex is the same as the underlving proton spectral iudex. here s-padSer=2.75: this is somewhat steeper than theFermi sinele-power-law fit δω]2.11£0.05. | For hadronic emission, the high-energy spectral index is the same as the underlying proton spectral index, here $s_{\gamma,\rm had} = s_{\rm cr} = 2.75$ ; this is somewhat steeper than the single-power-law fit $s_{\rm obs} = 2.41 \pm 0.05$. |
Consequently. our predictions at high energies (210 CGoV) fall below the data. | Consequently, our predictions at high energies $\ga 10$ GeV) fall below the data. |
If normal galaxies had a of cosmic-rav spectral indices. the resulting ECB spectu would steepeu at high energies where the hardest sources would dominate. developing afa Tudeed. theFeri EGB data suggests a slight flattening of slope arouud E>10 GeV. which might lint at such a transition. | If normal galaxies had a of cosmic-ray spectral indices, the resulting EGB spectrum would steepen at high energies where the hardest sources would dominate, developing a. Indeed, the EGB data suggests a slight flattening of slope around $E \ga 10$ GeV, which might hint at such a transition. |
A ealaxy with characteristic lhuumositv L. has £°¢>100MeV)=1.1«107s1, | A galaxy with characteristic luminosity $\tracer_*$ has $L_\gamma^*(>100 {\rm MeV}) = 1.4 \times 10^{43} \ {\rm s^{-1}}$. |
Such objects have flux F if they lie at distances kr;=(Leο---Alpe(10on2sΤΕΙ 3, ThosFerm? | Such objects have flux $F$ if they lie at distances $r_* = (L_\gamma^*/4\pi F)^{1/2} =
11 \, {\rm \ Mpc} \ ({10^{-9} \ \rm cm^{-2} \ s^{-1}}/F)^{1/2}$ . |
should eventually resolve NGE) bir? nn(0)3 =51ESV?(19) normal galaxies. cousisteut with 23 detectious to date (theLMC.SAIC.andperhapsM31:Abdoetal.20106:al. 2010).. | Thus should eventually resolve N(>F) r_*^3 n_*(0)/3 = 5 normal galaxies, consistent with 2–3 detections to date \citep[the LMC, SMC, and perhaps M31;][]{fermi-LMC,fermi-SMC,
fermi-M31}. |
Our results do not account for starburst galaxies. nor for inverse-C'oniptou emission from any star-forming ealaxies: these nist contribute to the star-forming EGB. aud could have lid spectra dominating =10 GeV. We have also neglected gamma-ray attenuation by extragalactic background light (importantatEz30GeV:e.g.YoSteckeretal.2006.andreferences therein). | Our results do not account for starburst galaxies, nor for inverse-Compton emission from any star-forming galaxies; these must contribute to the star-forming EGB, and could have hard spectra dominating $\ga 10$ GeV. We have also neglected gamma-ray attenuation by extragalactic background light \citep[important at $E \ga 30$ GeV; e.g.,][and references therein]{sms}. |
We will address these issues in future work. | We will address these issues in future work. |
The amplitude and configuration of magnetic fields ina ealaxy have an additional effect on the scaling of cosimic-rav flux with SER. | The amplitude and configuration of magnetic fields in a galaxy have an additional effect on the scaling of cosmic-ray flux with SFR. |
Coufirmation that normal galaxies colmprise the bulk of the Fermi signal would constitute a uniqueprobe of the evolution of these magnetic fields between the redshift of peak star formation aud todas. | Confirmation that normal galaxies comprise the bulk of the Fermi signal would constitute a uniqueprobe of the evolution of these magnetic fields between the redshift of peak star formation and today. |
Because of their ubiquity. normalgalaxies produce the sinallestauisotropies in the ECB. far less than blazirsor other| proposed sources. | Because of their ubiquity, normalgalaxies produce the smallestanisotropies in the EGB, far less than blazarsor other proposed sources. |
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