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Iu our Monute-Carlo calculations below. we adopt the latest stellarmass function obtained from the S-COSAIOS survey (Spitzer-Cosmic Evolution Survey) for galaxies at redshift ranges 02.OL OL0.6.0.6 O88. 0.8LO. and 1.02 2010).. aud that obtaiued from the Sloan Digital Sky Survey (SDSS) for galaxies iu the local universe (2~ 0.1)2010)... respectively. | In our Monte-Carlo calculations below, we adopt the latest stellarmass function obtained from the S-COSMOS survey (Spitzer-Cosmic Evolution Survey) for galaxies at redshift ranges $0.2 - 0.4$, $0.4-0.6$, $0.6-0.8$, $0.8-1.0$, and $1.0-1.2$ , and that obtained from the Sloan Digital Sky Survey (SDSS) for galaxies in the local universe $z\sim 0.1$ ), respectively. |
According to equation (1)) and the stellar muass functions. our sinulations generate a large nuniber of mereiug pairs of ealaxies over redshift 0 to 1.2. | According to equation \ref{eq:mrg}) ) and the stellar mass functions, our simulations generate a large number of merging pairs of galaxies over redshift $0$ to $1.2$. |
The masses of the two progenitors ave assigned by M,—max(aM+Lip)Al ancl AM,o=wilHawa+-.DEA). respectively,+ where e is+ randomly selected according to its distribution function f(r). During the mereiug process of two galaxies. whether sjeuificaut nuclear activities at their centers cau be trigeered depend on the two factors: (1) whether MDITs are mutiallv hosted iu both unelei or whether the MDBIIS are massive enoueh: aud (2) whether sufficient easeous materials cau be quickly delivered iuto the vicinity of the AIBUs. | The masses of the two progenitors are assigned by $M_{*,1}=\max(\frac{xM_*}{1+x},\frac{M_*}{1+x})$ and $M_{*,2}={\rm min}(\frac{xM_*}{1+x},\frac{M_*}{1+x})$, respectively, where $x$ is randomly selected according to its distribution function $f(x)$ During the merging process of two galaxies, whether significant nuclear activities at their centers can be triggered depend on the two factors: (1) whether MBHs are initially hosted in both nuclei or whether the MBHs are massive enough; and (2) whether sufficient gaseous materials can be quickly delivered into the vicinity of the MBHs. |
Those two factors are closely related to theprogenitor morphologics of the mereiug pairs. | Those two factors are closely related to theprogenitor morphologies of the merging pairs. |
Observatious have shown that the masses of ATBIs. A.. in theceuters of nearby galaxies are tightly correlated with the stellar inass of the spheroidal conrponeuts of the ealaxies Aejue 11). | Observations have shown that the masses of MBHs, $M\bh$ , in thecenters of nearby galaxies are tightly correlated with the stellar mass of the spheroidal components of the galaxies $M\bulge$ -11), |
from the detailed neutrino-radiation hydrodynamics calculation described in the previous subsection, are utilized to determine the actual rates as a function of time. | from the detailed neutrino-radiation hydrodynamics calculation described in the previous subsection, are utilized to determine the actual rates as a function of time. |
These reaction rates depend on the distance from the proto-neutron star and the mean energy and luminosity of neutrinos emitted from the proto-neutron star. | These reaction rates depend on the distance from the proto-neutron star and the mean energy and luminosity of neutrinos emitted from the proto-neutron star. |
It depends on the structure and the evolution, which is sensitive to the EOS, and the precise evolution history of ejected matter including the early phase of the core bounce. | It depends on the structure and the evolution, which is sensitive to the EOS, and the precise evolution history of ejected matter including the early phase of the core bounce. |
Thus, this is different from the treatment of other nuclear reaction rates, which are determined only by local thermodynamic conditions and density and temperature. | Thus, this is different from the treatment of other nuclear reaction rates, which are determined only by local thermodynamic conditions and density and temperature. |
In order to calculate the nucleosynthesis evolution of ejected matter within a postprocessing approach, the dynamic evolution is required in radial Lagrangian mass zones. | In order to calculate the nucleosynthesis evolution of ejected matter within a postprocessing approach, the dynamic evolution is required in radial Lagrangian mass zones. |
For this reason, the evolution determined with the radiation hydrodynamics code AGILE-BOLTZTRAN (seeMezzacappa&Bruenn1993a,b,c;Liebendorferetal.2001a,b) which is based on an adaptive grid, was mapped on a Lagrangian grid of 120 mass zones. | For this reason, the evolution determined with the radiation hydrodynamics code AGILE-BOLTZTRAN \citep[see][]{MezzacappaBruenn:1993a, MezzacappaBruenn:1993b,
MezzacappaBruenn:1993c,Liebendoerfer:etal:2001a,Liebendoerfer:etal:2001b}
which is based on an adaptive grid, was mapped on a Lagrangian grid of 120 mass zones. |
This provides the Lagrangian evolution of physical quantities, such as density, temperature, electron fraction, and velocity of the ejected material, and in addition the neutrino fluxes experienced as a function of time. | This provides the Lagrangian evolution of physical quantities, such as density, temperature, electron fraction, and velocity of the ejected material, and in addition the neutrino fluxes experienced as a function of time. |
The mass zones which are ejected in the explosion are classified in three different categories, related to their ejection process and thermodynamic quantities. | The mass zones which are ejected in the explosion are classified in three different categories, related to their ejection process and thermodynamic quantities. |
As listed in Table1, zones 001 to 120, given with the ejection timescale after bounce and the final which is Y, at the end of NSE (tj)(below T=9 GK), Yengsmcover the material from the surface of the inner core at 1.48000 Mo to layers with a corresponding mass of 1.49482 Mo. | As listed in Table, zones 001 to 120, given with the ejection timescale after bounce $t_{\rm{ej}}$ ) and the final $Y_{e,\rm{NSE}}$ which is $Y_e$ at the end of NSE (below $T=9$ GK), cover the material from the surface of the inner core at 1.48000 $M_{\odot}$ to layers with a corresponding mass of 1.49482 $M_{\odot}$. |
They are also shown with respect to their mass as well as Y,-distribution in Figure 1.. | They are also shown with respect to their mass as well as $Y_e$ -distribution in Figure \ref{fig-init}. |
'These zones coincide with the matter discussed at the end of Section 2.2, where at high densities a Y; decrease down to even 0.02 has been noticed. | These zones coincide with the matter discussed at the end of Section 2.2, where at high densities a $Y_e$ decrease down to even $0.02$ has been noticed. |
The evolution of these mass zones in time is displayed in Figure2. | The evolution of these mass zones in time is displayed in Figure. |
Zones 001 to 014 are not ejected within 0.5 s after the core bounce. | Zones 001 to 014 are not ejected within $0.5$ s after the core bounce. |
These zones preserve the original low-Y. obtained during collapse and shock wave propagation. | These zones preserve the original $Y_e$ obtained during collapse and shock wave propagation. |
As they are not ejected, we ignore them in the further nucleosynthesis discussion, plus all matter originating from regions at smaller radii. | As they are not ejected, we ignore them in the further nucleosynthesis discussion, plus all matter originating from regions at smaller radii. |
In Figure 2, they are displayed in black. | In Figure , they are displayed in black. |
Zones 015 to 019 are ejected in the so- NDW, shown in red. | Zones 015 to 019 are ejected in the so-called NDW, shown in red. |
Their Y, is strongly affected by neutrino interactions, turning this matter proton-rich. | Their $Y_e$ is strongly affected by neutrino interactions, turning this matter proton-rich. |
Zones 020 to 050 (displayed in green) have stalled from infall after shock formation and are ejected thereafter due to neutrino heating and dynamic effects (Fischeretal. 2011).. | Zones 020 to 050 (displayed in green) have stalled from infall after shock formation and are ejected thereafter due to neutrino heating and dynamic effects \citep{Fischer:etal:2011}. |
The adjacent zones 051 to 120 are ejected in a prompt way, due to the shock wave originating from the deconfinement phase transition (displayed in blue). | The adjacent zones 051 to 120 are ejected in a prompt way, due to the shock wave originating from the deconfinement phase transition (displayed in blue). |
We clearly see the division of matter which is ejected in a prompt fashion(blue), matter which is coasting | We clearly see the division of matter which is ejected in a prompt fashion(blue), matter which is coasting |
funelions for the sources. | functions for the sources. |
Finally. in Section 6 we discuss (he implications of our findings on our understanding of Cvguus OB? aud compare it {ο other massive SETs. | Finally, in Section 6 we discuss the implications of our findings on our understanding of Cygnus OB2 and compare it to other massive SFRs. |
The location of Cygnus OD2 in the Galactic Plane introduces the non-trivial issue of contamination [rom foreground ancl background sources. | The location of Cygnus OB2 in the Galactic Plane introduces the non-trivial issue of contamination from foreground and background sources. |
This problem is further complicated bv the high and variable extinction toward and throughout the Cvenus region. | This problem is further complicated by the high and variable extinction toward and throughout the Cygnus region. |
There are many methods Chat have been used to select members of a star forming region while minimizing (he contanunation Irom the field population. but one of the most effective ancl efficient is to use X-rav enission as a (racer of vouth. since voung pre-MS stars are orders of magnitude more luminous in A-ravs (han their ALS equivalents (7).. | There are many methods that have been used to select members of a star forming region while minimizing the contamination from the field population, but one of the most effective and efficient is to use X-ray emission as a tracer of youth, since young pre-MS stars are orders of magnitude more luminous in X-rays than their MS equivalents \citep{prei05}. |
In. Paper 1 à catalogue of 1696 sources X-rav sources was presented. extracted from observations of (svo fields in the center of Cve OD2. | In Paper 1 a catalogue of 1696 sources X-ray sources was presented, extracted from observations of two fields in the center of Cyg OB2. |
The catalogue also included optical aud θαΕΠ associations taken from the IPILAS (INTPhotometricHaSurvey. ?).. 2ALASS (TwoMicron.AllSkySurvey.7) and UNIDSS (UnitedIxingdomInfraredDeepSkySurvey.ο) photometric catalogues. | The catalogue also included optical and near-IR associations taken from the IPHAS \citep[INT Photometric H$\alpha$ Survey,][]{drew05}, , 2MASS \citep[Two Micron All Sky Survey,][]{skru06} and UKIDSS \citep[United Kingdom Infrared Deep Sky Survey,][]{luca08} photometric catalogues. |
This work is based on this catalogue and principally the 1501 sources with either optical or near-IR counterparts (hat represents (he main stellar sample. | This work is based on this catalogue and principally the 1501 sources with either optical or near-IR counterparts that represents the main stellar sample. |
The completeness of (his catalogue is dependent on both the X-ray detections and the associations al other wavelengths. | The completeness of this catalogue is dependent on both the X-ray detections and the associations at other wavelengths. |
In Section 4.2. we show. by comparison wilh deeper N-rax stellar catalogues. that the completeness limit for (his catalogue is ~1 M... | In Section \ref{s-xlf} we show, by comparison with deeper X-ray stellar catalogues, that the completeness limit for this catalogue is $\sim$ 1 $_{\odot}$. |
In Paper 1 il was shown that the near-IB. photometry combined from 2MÁÀSS and UNIDSS observations will allow pre-MS stars down to 0.2 M. (o be detected al an extinction of “ly~7. | In Paper 1 it was shown that the near-IR photometry combined from 2MASS and UKIDSS observations will allow pre-MS stars down to 0.2 $_{\odot}$ to be detected at an extinction of $A_V \sim 7$. |
In the north-western field. where some UIXIDSS data is lacking. it was estimated in Paper 1 that the 2NLASS observations reach à pre-MS depth of 1 AL. at Ay~7. though the lower extinction in this field means the observations will be shehtly deeper than (his. | In the north-western field, where some UKIDSS data is lacking, it was estimated in Paper 1 that the 2MASS observations reach a pre-MS depth of 1 $_{\odot}$ at $A_V \sim 7$, though the lower extinction in this field means the observations will be slightly deeper than this. |
Since not all voung stars emit N-ravs we must also be aware of the limits of emüssion mechanisms. | Since not all young stars emit X-rays we must also be aware of the limits of X-ray emission mechanisms. |
O and early B-ivpe stars (earlier (han (wpe D2. >10 AI. are expected to generate N-ravs in their racliatively-criven stellar winds (e.g.2).. while. al ages of ~5 Myr. lower mass stars («3.3 M.) have substantial convective envelopes Chat produce X-ray emission through magnetic dvnamo activity (e.g. 2).. | O and early B-type stars (earlier than type B2, $>$ 10 $_{\odot}$ are expected to generate X-rays in their radiatively-driven stellar winds \citep[e.g.][]{berg97}, while, at ages of $\sim$ 5 Myr, lower mass stars $<$ 3.3 $_{\odot}$ ) have substantial convective envelopes that produce X-ray emission through magnetic dynamo activity \citep[e.g.][]{pall81}. . |
Intermediate mass stars do not have deep convective lavers or strong stellar winds and are therefore not expected to emit | Intermediate mass stars do not have deep convective layers or strong stellar winds and are therefore not expected to emit X-rays. |
Studies that have observed X-ray emission [rom late D and A-type stars have attributed (his to coronal emission from unresolved late-tvpe companions (e.g.?).. | Studies that have observed X-ray emission from late B and A-type stars have attributed this to coronal emission from unresolved late-type companions \citep[e.g.][]{stel06}. |
The mass completeness regime can be determined by counting the number of spectroscopically known OD stars in our field of view with X-rav detections. | The intermediate-mass completeness regime can be determined by counting the number of spectroscopically known OB stars in our field of view with X-ray detections. |
We findthat the fraction ofOD | We findthat the fraction ofOB |
from hard to soft state. | from hard to soft state. |
We thank F. Mever ancl E. Mever-Iofneister for detailed discussions on the computational results aud [ον their reading through (he manuscript aud comments on it. | We thank F. Meyer and E. Meyer-Hofmeister for detailed discussions on the computational results and for their reading through the manuscript and comments on it. |
We also (hank FI. Liu for helpful discussions. | We also thank F.K. Liu for helpful discussions. |
Q.L. wishes to thank the hospitality of High Enerey Astronomy Group al YNAQO. which makes this work possible. | Q.L. wishes to thank the hospitality of High Energy Astronomy Group at YNAO, which makes this work possible. |
This work is partially supported by the National Natural Science Foundation of China (Grants-10533050. 10413001. 10525313). ihe DailtenJillua program of the Chinese Academy of Sciences. and the RFDP (Grant 20050001026). | This work is partially supported by the National Natural Science Foundation of China (Grants-10533050, 10473001, 10525313), the BaiRenJiHua program of the Chinese Academy of Sciences, and the RFDP (Grant 20050001026). |
and Ursa MajorllI. ? find very homogeneous [Mg/Ca] abundance ratios refmgea.fig)) and fairly homogeneous abundance trends for the a-elements refcamgfe.fig)). | and Ursa II, \citet{frebel2009CBU} find very homogeneous [Mg/Ca] abundance ratios \\ref{mgca.fig}) ) and fairly homogeneous abundance trends for the $\alpha$ -elements \\ref{camgfe.fig}) ). |
These findings seem to suggest that the ultra-faint dSph galaxies are different from the classical dSph galaxies. | These findings seem to suggest that the ultra-faint dSph galaxies are different from the classical dSph galaxies. |
Indeed. Draco. a classical dSph. has one odd star (?).. (Compare.e.g.. | Indeed, Draco, a classical dSph, has one odd star \citep{fulbright2004Draco}, \citep[Compare, e.g., the recent results in][which shows that eight of
the bright red giants in Draco all have normal Mg to Ca ratios.]{2009ApJ...701.1053C}. |
resolution cannot be as high as in the simulations which focus on single haloes. | resolution cannot be as high as in the simulations which focus on single haloes. |
This paper is structured in the following wav. | This paper is structured in the following way. |
We summarize the halo model and give relevant. background information in re[sec:Background.. | We summarize the halo model and give relevant background information in \\ref{sec:Background}. . |
In re[sec:Methodc.. we give our method and results. | In \\ref{sec:Method}, we give our method and results. |
In refsec:Comparison.. we compare our results to other mocels and in refsec:Conclusions we summarize our conclusions. | In \\ref{sec:Comparison}, we compare our results to other models and in \\ref{sec:Conclusions} we summarize our conclusions. |
We begin by reviewing the formalism of the halo model to compute the non-linear dark matter power spectra. following the notation in ?.. | We begin by reviewing the formalism of the halo model to compute the non-linear dark matter power spectra, following the notation in \citet{2000MNRAS.318..203S}. |
The haloes are characterized by their mass M. density profile simulationspCr.AZ). number density ÁCÀAZ) and bias bCAL) | The haloes are characterized by their mass $M$, density profile $\rho(r,M)$, number density $n(M)$ and bias $b(M)$. |
V-body suggest a amily of density profiles: where ος is a characteristic density. r£, is the radius where the profile has an effective power law index of 9. and Lo<a<1 (??) . | $N$ -body simulations suggest a family of density profiles: where $\rho_s$ is a characteristic density, $r_s$ is the radius where the profile has an effective power law index of $-2$, and $-1.5<\alpha<-1$ \citep{1997ApJ...490..493N,1998ApJ...499L...5M}. |
llere we use aΞ 1. since power spectra are not sensitive to the inner parts of the halo. | Here we use $\alpha = -1$ , since power spectra are not sensitive to the inner parts of the halo. |
As isconventional. we re-parametrize o, and r, in equation (1)) in terms of a halo mass and concentration. | As isconventional, we re-parametrize $\rho_s$ and $r_s$ in equation \ref{eqn:rho(r)}) ) in terms of a halo mass and concentration. |
We define the virial radius ré of a halo to be the radius ola sphere with some characteristic mean density. cliscussed xdlow. | We define the virial radius $r_{\rm vir}$ of a halo to be the radius of a sphere with some characteristic mean density, discussed below. |
Insisting that the mass contained inside roi is M fixes ps lor a given εν, | Insisting that the mass contained inside $r_{\rm vir}$ is $M$ fixes $\rho_s$ for a given $r_s$. |
We introduce the concentration parameter. he ratio 6—ru | We introduce the concentration parameter, the ratio $c \equiv r_{\rm vir}/r_s$. |
There are several definitions of virial racdus used in the literature. | There are several definitions of virial radius used in the literature. |
In this paper. the virial radius is defined as the radius of a halo-centreed: sphere which las a mean density 180 times the mean density of the universe. | In this paper, the virial radius is defined as the radius of a halo-centreed sphere which has a mean density 180 times the mean density of the universe. |
This may be denoted rjsoo. | This may be denoted $r_{180\Omega}$. |
Other authors set rig o the radius inside which the mean density is 200 times heσας density (rosy). à measure independent of O (7. example).. | Other authors set $r_{\rm vir}$ to the radius inside which the mean density is 200 times the density $r_{200}$ ), a measure independent of $\Omega$ \citep[for example]{1997ApJ...490..493N}. |
Still others use à spherical modeΙ o estimate the radius of a virialized halo Por . | Still others use a spherical collapse model to estimate the radius of a virialized halo $r_{\Delta}$ ). |
mThese Lead o dillerent concentrations and. masses. | These lead to different concentrations and masses. |
.=0.3 ancl A=OF. ουSPAXl'asuo- for the same halo with Crsno=10 one meCogo6. ὃν~8S. ουfALisue~0.72 and Aly/Alixsug"m OST. | For $\Omega=0.3$ and $\Lambda=0.7$, $r_{200} < r_\Delta < r_{180\Omega}$, so for the same halo with $c_{180\Omega}=10$ one has $c_{200} \sim 6$, $c_\Delta\sim 8$, $M_{200}/M_{180\Omega} \sim 0.72$ and $M_\Delta/M_{180\Omega} \sim 0.87$ . |
Assuming the NEW halo profile it is straight to translate values for concentrations and masses between these conventions. | Assuming the NFW halo profile it is straightforward to translate values for concentrations and masses between these conventions. |
We choose equo as it permits the use of the universal mass function. (7).. | We choose $c_{180\Omega}$ as it permits the use of the universal mass function \citep{2002astro.ph..7185W}. |
To avoid cumbersome notation. in this paper we refer to ejxoo as c and AMixoo as M. | To avoid cumbersome notation, in this paper we refer to $c_{180\Omega}$ as $c$ and $M_{180\Omega}$ as $M$. |
For the mass dependence of concentration we choose a power law parametrization where ος and 3 are free parameters. and AZ; is the non-linear mass scale. which we will define shortly. | For the mass dependence of concentration we choose a power law parametrization where $c_0$ and $\beta$ are free parameters, and $M_*$ is the non-linear mass scale, which we will define shortly. |
This is of course not the most general parametrization. but as we will show below it sullices for the current dynamical range. | This is of course not the most general parametrization, but as we will show below it suffices for the current dynamical range. |
The halo number density is written in terms of the multiplicity function via where p is the mean matter density of the universe and mass is written in terms of peak height where δν(ο) is the spherical over-cdensity which collapses at z (ὃς&L.68) and ofAl) is the rms Uuctuation in the matter density smoothed on a scale 2=(3Mπρ). | The halo number density is written in terms of the multiplicity function via where $\bar{\rho}$ is the mean matter density of the universe and mass is written in terms of peak height where $\delta_c(z)$ is the spherical over-density which collapses at $z$ $\delta_c \approx 1.68$ ) and $\sigma(M)$ is the rms fluctuation in the matter density smoothed on a scale $R=(3M/4\pi \bar{\rho})^{1/3}$. |
eh(Al,)ὃν=1 defines the non-linear mass scale. | $\nu(M_*) = 1$ defines the non-linear mass scale. |
A cosmology =03. \=07. Ll=02. and ox=0.9 has AL=107h1ALL. | A cosmology with $\Omega_0 = 0.3$, $\Lambda = 0.7$, $\Gamma=0.2$, and $\sigma_8=0.9$ has $M_* \approx 10^{13}\ h^{-1}\ M_{\sun}$. |
Te power spectrum has two terms. | The power spectrum has two terms. |
The first corresponds to correlations in density between: pairs of points where each member of the pair lies in a cillerent halo. and so is named the halo-halo (hh) term. | The first corresponds to correlations in density between pairs of points where each member of the pair lies in a different halo, and so is named the “halo-halo” (hh) term. |
Phe second corresponds to correlations between pairs in the same halo. and is known as the one-halo or Poisson (I) term. | The second corresponds to correlations between pairs in the same halo, and is known as the one-halo or Poisson (P) term. |
For convenience in calculating convolutions. we work in Fourier space. introducing the Fourier transform of the halo profile. normalized bv the virial mass. The mass of the NEW profile is. logarithmically. divergent. so in order to evaluate this integral. we must impose a cutoll. | For convenience in calculating convolutions, we work in Fourier space, introducing the Fourier transform of the halo profile, normalized by the virial mass, The mass of the NFW profile is logarithmically divergent, so in order to evaluate this integral, we must impose a cutoff. |
Ehis does not have to be at the virial radius. since we know that haloes are not completely truncated there. | This does not have to be at the virial radius, since we know that haloes are not completely truncated there. |
Instead. NEW profile typically continues (ο 31. | Instead, NFW profile typically continues to $2-3r_{\rm vir}$. |
In this regime there is already. some overlap between the haloes. so for the purpose of correlations one can count the same mass element in more than one halo. | In this regime there is already some overlap between the haloes, so for the purpose of correlations one can count the same mass element in more than one halo. |
Thus the mass function integrated over all the haloes may even exceed the mean censity of the universe (but it can also be below it. since it is not πουτος that all the matter should be inside a halo). | Thus the mass function integrated over all the haloes may even exceed the mean density of the universe (but it can also be below it, since it is not required that all the matter should be inside a halo). |
The halo-halo contribution is where b(7) is the (linear) bias of a halo of mass Al(y)this for which we use the A/N-bodvy fit of7.. | The halo-halo contribution is where $b(\nu)$ is the (linear) bias of a halo of mass $M(\nu)$, for which we use the $N$ -body fit of \citet{1999MNRAS.308..119S}. |
Since we want term to reproduce the linear power spectrum on Large scales we impose this constraint onto the form for bf) (?).. | Since we want this term to reproduce the linear power spectrum on large scales we impose this constraint onto the form for $b(\nu)$ \citep{2000MNRAS.318..203S}. |
The Poisson contribution is Ifa cutoll beyond the virial racius is used. the mass weighting should rellect the increase. | The Poisson contribution is If a cutoff beyond the virial radius is used, the mass weighting should reflect the increase. |
The treatment of the Poisson term on large scales is only approximate. | The treatment of the Poisson term on large scales is only approximate. |
Mass ancl momentum conservation require that on very large. scales the non-linear term should. seale as A3. rather than as a constant implied. by equation. 7.. so the contribution from this term on large scales is overestimated [for αςεςd. | Mass and momentum conservation require that on very large scales the non-linear term should scale as $k^4$, rather than as a constant implied by equation \ref{eqn:Poisson}, so the contribution from this term on large scales is overestimated for $k\ll r_{\rm vir}^{-1}$. |
‘This partially compensates the increase in power from the matter outside the virial radius and for this reason we chose to useFue as the radial cutoll for the halo. | This partially compensates the increase in power from the matter outside the virial radius and for this reason we chose to use$r_{\rm vir}$ as the radial cutoff for the halo. |
The halo-halo term is only approximate. since we do not include the exclusion of haloes. which would suppress the term. | The halo-halo term is only approximate, since we do not include the exclusion of haloes, which would suppress the term. |
Because of these approximations one wouldnot expect the halo model to be | Because of these approximations one wouldnot expect the halo model to be |
remaining fraction not accounted for by our fits. this is still not a competitive method for measuring the total CFIRB. | remaining fraction not accounted for by our fits, this is still not a competitive method for measuring the total CFIRB. |
We find clear evidence of breaks in the slope of the ditferential number counts at approximately 10—20 mJy in all bands. which have been hinted at by previous analyses. | We find clear evidence of breaks in the slope of the differential number counts at approximately 10–20 mJy in all bands, which have been hinted at by previous analyses. |
Where they overlap. our fits agree well with otherHerschel results. | Where they overlap, our fits agree well with other results. |
Comparing with a selection of literature models. however. we find that no model entirely reproduces our observed number counts. | Comparing with a selection of literature models, however, we find that no model entirely reproduces our observed number counts. |
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