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On the other hand. a magnetic field that penetrates deep into a star causes nearly uniform rotation (Mossetal.1990;Charbonneau&MacGregor1992;Maeder&Meynet2003). | On the other hand, a magnetic field that penetrates deep into a star causes nearly uniform rotation \citep{moss,uhlaci,mame}. |
. The possible role of stellar winds in the mechanism of the rotational instability is supported by the fact that stronger period variations are detected inOri... which is more luminous than aand. consequently. has a stronger stellar wind (Krtiéka&Kubát2001). | The possible role of stellar winds in the mechanism of the rotational instability is supported by the fact that stronger period variations are detected in, which is more luminous than and, consequently, has a stronger stellar wind \citep{krku}. |
. Moreover. this could help to explain the fact that the younger star rrotates more slowly thanVir. | Moreover, this could help to explain the fact that the younger star rotates more slowly than. |
. All these considerations should be the subject of additional physical modelling. | All these considerations should be the subject of additional physical modelling. |
Nevertheless. the aim of this letter is more unpretentious — we only want to draw the attention to an interesting property of the rotation of some upper MS stars. | Nevertheless, the aim of this letter is more unpretentious – we only want to draw the attention to an interesting property of the rotation of some upper MS stars. |
Our study investigated the nature of the rotational period variations in two well-observed mCP stars παπάOri. | Our study investigated the nature of the rotational period variations in two well-observed mCP stars and. |
. Contrary to the results of previous studies. we show that the period changes are not monotonic — intervals of rotational deceleration alternate with intervals of rotational acceleration. all on the timescale of several decades. | Contrary to the results of previous studies, we show that the period changes are not monotonic – intervals of rotational deceleration alternate with intervals of rotational acceleration, all on the timescale of several decades. |
These results explain the spin-down time paradox of these stars. according to which the spin-down time was significantly shorter than the age of these stars. | These results explain the spin-down time paradox of these stars, according to which the spin-down time was significantly shorter than the age of these stars. |
On the other hand. this unexpected behaviour of two fairly dissimilar mCP stars poses a strong challenge for any theoretical models. | On the other hand, this unexpected behaviour of two fairly dissimilar mCP stars poses a strong challenge for any theoretical models. |
how the global neutral fraction evolves: this technique has been used is countless reionization studies. | how the global neutral fraction evolves; this technique has been used is countless reionization studies. |
Aside from the seemingly magical "+" in the previous sentence. one of the main additiona uncertainties are the feedback processes: how do sources impact the current and future generations of sources? | Aside from the seemingly magical $\rightarrow$ ” in the previous sentence, one of the main additional uncertainties are the feedback processes: how do sources impact the current and future generations of sources? |
Answering this question is non-trivial. especially in the early universe. | Answering this question is non-trivial, especially in the early universe. |
Radiative. chemica and mechanical feedback can effect the ionizing efficiencies of the first. highly biased sources. | Radiative, chemical and mechanical feedback can effect the ionizing efficiencies of the first, highly biased sources. |
Molecular hydrogen (H2) cooling can allow very small Qwith virial temperatures. ων. of several hundred Kelvin) halos to host astrophysical sources. | Molecular hydrogen $_2$ ) cooling can allow very small (with virial temperatures, $\Tvir$, of several hundred Kelvin) halos to host astrophysical sources. |
At later times (though still prior to the bulk of reionization. 2). as the H» dissociative background builds-up and the contribution of molecularly-coolec halos declines. feedback enters the regime of atomically-cooled halos. Jim10! K. This regime is more straightforward to model as negative radiative feedback emerges as a single dominan mechanism. especially on large scales. | At later times (though still prior to the bulk of reionization, \citealt{HRL97}) ), as the $_2$ dissociative background builds-up and the contribution of molecularly-cooled halos declines, feedback enters the regime of atomically-cooled halos, $\Tvir \gsim 10^4$ K. This regime is more straightforward to model as negative radiative feedback emerges as a single dominant mechanism, especially on large scales. |
In the presence of an ionizing background. radiation. the IGM is photo-heated to a temperature of = 107K. raising the cosmological Jeans mass. which could suppress gas accretion onto small-mass halos (e.g. 22223). | In the presence of an ionizing background radiation, the IGM is photo–heated to a temperature of $\gsim 10^4$ K, raising the cosmological Jeans mass, which could suppress gas accretion onto small-mass halos (e.g. \citealt{Efstathiou92, SGB94, TW96, HG97}) ). |
Early work on this subject (so-called "Jeans mass filtering) suggested that an ionizing background would completely suppress star formation in low—redshift "dwarf galaxy” halos with circular velocities rein35km s.l. and partially suppress star-formation in halos with 35 kms+ Zorn < 100 kms! (2). | Early work on this subject (so-called “Jeans mass filtering”) suggested that an ionizing background would completely suppress star formation in low–redshift “dwarf galaxy” halos with circular velocities $v_{\rm circ} \lsim~35~\kmps$ , and partially suppress star–formation in halos with 35 $\kmps$ $\lsim$ $v_{\rm
circ}$ $\lsim$ 100 $\kmps$ \citep{TW96}. |
Many reionization studies assume prescriptions of gas suppression based on these results (e.g. 22222)). | Many reionization studies assume prescriptions of gas suppression based on these results (e.g. \citealt{RS03, OM04, WC07, Iliev07, WBH08}) ). |
However. more recent studies (22). find that at <3. more compact halo profiles. increased cooling efficiencies. and shorter exposure times to the ultraviolet background (UVB) could lessen the importance of negative radiative feedback. | However, more recent studies \citep{KI00, Dijkstra04} find that at $z \gsim 3$, more compact halo profiles, increased cooling efficiencies, and shorter exposure times to the ultraviolet background (UVB) could lessen the importance of negative radiative feedback. |
We expand on these works by exploring a wider parameter space. placed in a broader context of an inhomogenious reionization with a patchy UVB due to both source clustering and modulation by HII regions. | We expand on these works by exploring a wider parameter space, placed in a broader context of an inhomogenious reionization with a patchy UVB due to both source clustering and modulation by HII regions. |
? studied the effects of such radiative feedback in cosmological simulations: however. these simulation boxes Were of necessity very small and only included a single reionization model. | \citet{Gnedin00filter} studied the effects of such radiative feedback in cosmological simulations; however, these simulation boxes were of necessity very small and only included a single reionization model. |
reionization. We do not attempt to self-consistently model feedback during reionization: such a thing is beyond the capacity of current simulations. especially given our poor understanding of the first generations of astrophysical sources and their environments. | We do not attempt to self-consistently model feedback during reionization; such a thing is beyond the capacity of current simulations, especially given our poor understanding of the first generations of astrophysical sources and their environments. |
Instead we statistically present the effects of an inhomogeneous UVB on the suppression of gas content in low-mass galaxiescooling. i.e. Ty=101 K. We do this using a tiered approach: using numerical simulations (22). to calibrate very large scale. high resolution "semi-numerical" simulations (2).. | Instead we statistically present the effects of an inhomogeneous UVB on the suppression of gas content in low-mass galaxies, i.e. $\Tvir \gsim 10^4$ K. We do this using a tiered approach: using numerical simulations \citep{TW95, Dijkstra04} to calibrate very large scale, high resolution “semi-numerical” simulations \citep{MF07}. |
In order to keep things general. we explore wide swaths of parameter space keeping assumptions minimal. | In order to keep things general, we explore wide swaths of parameter space keeping assumptions minimal. |
This paper is organized as follows. | This paper is organized as follows. |
In 32.1. and 32.2. we describe our hydrodynamic collapse simulations and semi-numerical cosmological simulations. respectively. | In \ref{sec:mark_sims} and \ref{sec:my_sims}, we describe our hydrodynamic collapse simulations and semi-numerical cosmological simulations, respectively. |
In $3.1. we present the results from our collapse simulations covering a wide range of parameter choices. while in $3.2 we present parametrized distributions of UV fluxes from our cosmological simulations. | In \ref{sec:fcoll} we present the results from our collapse simulations covering a wide range of parameter choices, while in \ref{sec:dist} we present parametrized distributions of UV fluxes from our cosmological simulations. |
In $3.3. we combine these results to quantify the importance of UV feedback during reionization. | In \ref{sec:feedback} we combine these results to quantify the importance of UV feedback during reionization. |
In S+ we discuss the assumptions and uncertainties in our approach. | In \ref{sec:ass} we discuss the assumptions and uncertainties in our approach. |
Finally. in S5.we present our conclusions. | Finally, in \ref{sec:conc} we present our conclusions. |
We quote all quantities in comoving units. with the exception of flux. where we denote proper units with a prefix 'p'. | We quote all quantities in comoving units, with the exception of flux, where we denote proper units with a prefix 'p'. |
We adopt the background cosmological parameters (Oy. O51. O,. n. σκι Ho) = (0.76. 0.24. 0.0407. 0.96. 0.76. 72 km + 1. matching the three-year results of theWAZAP satellite (2).. which in turn are consistent with the recent five-year data release (2).. | We adopt the background cosmological parameters $\Omega_\Lambda$, $\Omega_{\rm M}$, $\Omega_b$, $n$, $\sigma_8$, $H_0$ ) = (0.76, 0.24, 0.0407, 0.96, 0.76, 72 km $^{-1}$ $^{-1}$ ), matching the three–year results of the satellite \citep{Spergel07}, which in turn are consistent with the recent five-year data release \citep{Komatsu08}. . |
As we outlined above. we use spherically-symmetrie simulations to model the collapse of gas onto halos with masses M2107A.. | As we outlined above, we use spherically-symmetric simulations to model the collapse of gas onto halos with masses $M \gsim 10^8 \Msun$. |
With these numerical simulations. we are able to determine the amount of gas which collapses onto a halo of mass Ad at 2. under the presence of a UVB with specitie intensity οι turned on at redshift τω. | With these numerical simulations, we are able to determine the amount of gas which collapses onto a halo of mass $M$ at $z$, under the presence of a UVB with specific intensity $J_{21}$ turned on at redshift $\zon$. |
We then generate large-scale halo. ionization and parametrized flux fields. using semi-numerical techniques. | We then generate large-scale halo, ionization and parametrized flux fields, using semi-numerical techniques. |
Combining these results. we investigate what fraction of halos experience strong negative radiative feedback. | Combining these results, we investigate what fraction of halos experience strong negative radiative feedback. |
We describe our methodology and numerical techniques in greater detail below. | We describe our methodology and numerical techniques in greater detail below. |
We investigate the impact of a photoionizing flux on the ability of gas to cool and collapse onto dark matter halos by performing I-D hydrodynamical simulations which include both gas and dark matter. | We investigate the impact of a photoionizing flux on the ability of gas to cool and collapse onto dark matter halos by performing 1-D hydrodynamical simulations which include both gas and dark matter. |
We use the code that was originally written by ? and modifiedcharacterized to study the impact of photoionization feedback on the formation of high-redshift low-mass galaxies (see Dijkstra et al 2004 for a more detailed description). | We use the code that was originally written by \citet{TW95} and modifiedcharacterized to study the impact of photoionization feedback on the formation of high-redshift low-mass galaxies (see Dijkstra et al 2004 for a more detailed description). |
To summarize our calculations: an externally generated UVB impinges upon a sphericaly symmetric collapsing dark matter halo of mass AZ that would have collapsed to r=0 at redshift 2. in the absence of gas pressure. | To summarize our calculations: an externally generated UVB impinges upon a spherically symmetric collapsing dark matter halo of mass $M$ that would have collapsed to $r=0$ at redshift $z$, in the absence of gas pressure. |
T1ο ionizing radiation field is switched on at 244. and is characterized by a power-law spectrum of the form I(t)=οιHH)mmü21 ergs n 1 TI sr. where Mgds the ionization threshold of hydrogen. | The ionizing radiation field is switched on at $z_{\rm on}$, and is characterized by a power-law spectrum of the form $J(\nu)=J_{\rm 21}(\nu/\nu_H)^{-\alpha}\times 10^{-21}$ erg $^{-1}$ $^{-1}$ $^{-2}$ $^{-1}$, where $\nu_H$is the ionization threshold of hydrogen. |
We are interested in the fraction of the gas. fi;=Aus(p)/Matp0). (ie. the mass of collapsed gas divided by the mass of gas which would have collapsed in the absence of pressure) that is instantly available for star formation at ο. | We are interested in the fraction of the gas, $\fcoll \equiv M_{\rm gas}(p)/M_{\rm gas}(p=0)$, (i.e. the mass of collapsed gas divided by the mass of gas which would have collapsed in the absence of pressure) that is instantly available for star formation at $z$ . |
Therefore. we compute the fraction of gas which has collapsed at 2 (with an associated Hubble time /o.1) which differs from the calculations presented by e.g. Thoul Weinberg (1996). and Dijkstra et al. ( | Therefore, we compute the fraction of gas which has collapsed at $z$ (with an associated Hubble time $t_{\rm coll}$ ) which differs from the calculations presented by e.g. Thoul Weinberg (1996), and Dijkstra et al. ( |
2004) who computed the fraction of gas which was able to collapse up to 2/,.. | 2004) who computed the fraction of gas which was able to collapse up to $2t_{\rm coll}$. |
For all runs. we switch on the UVB at ται14. | For all runs, we switch on the UVB at $\zon=14$. |
The 5- WMAP polarization data place a constraint on the reionization redshift (assuming instantaneous reionization) of >=11.0+1.4 (οι. | The 5-yr polarization data place a constraint on the reionization redshift (assuming instantaneous reionization) of $z=11.0 \pm 1.4$ \citep{Dunkley08}. |
Thus it is highly unlikely that the majority of the universe was ionized before 2~14 (for some interpretations of the less-constrained 3-yr data. see. for example ??)). | Thus it is highly unlikely that the majority of the universe was ionized before $z\sim14$ (for some interpretations of the less-constrained 3-yr data, see, for example \citealt{HB06, CF06}) ). |
Therefore. our choice Of ton= Lis conservative. since most halos should be exposed to ionizing radiation at later times and so we overestimate the impact of radiative feedback. | Therefore, our choice of $\zon=14$ is conservative, since most halos should be exposed to ionizing radiation at later times and so we overestimate the impact of radiative feedback. |
Furthermore. following Dijkstra et al. ( | Furthermore, following Dijkstra et al. ( |
2004) wehave assumed a= 1.0. which is typically associated with accreting black holes (e.g. 29) and not with stellar sources which would be better characterized with a~5 €??).. | 2004) wehave assumed $\alpha=1.0$ , which is typically associated with accreting black holes (e.g. \citealt{Madau04})) and not with stellar sources which would be better characterized with $\alpha\sim5$ \citep{TW96, BL01}. . |
Because the simulations ignore self-shielding. lowering à is equivalent to increasing JJ», in terms of the photoionization and photoheating rates. | Because the simulations ignore self-shielding, lowering $\alpha$ is equivalent to increasing $J_{\rm 21}$ in terms of the photoionization and photoheating rates. |
Therefore. in our calculations we overestimate the impact of the photoionization feedback (see 4+ for àmore extended discussion). | Therefore, in our calculations we overestimate the impact of the photoionization feedback (see \ref{sec:ass} for amore extended discussion). |
The Galactic center (GC) has loug been suspected of harboring a central mass conceutration whose eravitational iufiueuce may be the cause of the high-velocity ionized gas streamers emitting in the 12 san [Ne Π line (Woμαι1976: Lacy.etal. 1979: for a recent review. sce Mezeer.Duschl&Zvlka1996)). | The Galactic center (GC) has long been suspected of harboring a central mass concentration whose gravitational influence may be the cause of the high-velocity ionized gas streamers emitting in the 12 $\mu$ m [Ne II] line \cite{wol76}; \cite{lacy79}; for a recent review, see \cite{mez96}) ). |
This ionized gas appears to be in orbit about a mass distribution 3«109AZ. within a few arcseconds (170011 pe at the GC) of the compac radio source Ser A* (Lo&Clausse11982: 1985)). | This ionized gas appears to be in orbit about a mass distribution $\sim 3\times 10^6\;M_\odot$ within a few arcseconds $^{\prime\prime}\approx 0.04$ pc at the GC) of the compact radio source Sgr A* \cite{lo83}; \cite{sera85}) ). |
À nore recent mapping of the 20 line emission from Ser A West wihan augularC» resolution of ~1" shows the preseuce of three dominaif kinenatie features. kuowu as the Western Are. the Northern Arm. and he Bar (Roberts&Coss 1993)). | A more recent mapping of the $\alpha$ line emission from Sgr A West with an angular resolution of $\sim 1^{\prime\prime}$ shows the presence of three dominant kinematic features, known as the Western Arc, the Northern Arm, and the Bar \cite{rob93}) ). |
The former appears to be iu circuir rotation about Ser A* at a radius of ~1 pc. | The former appears to be in circular rotation about Sgr A* at a radius of $\sim 1$ pc. |
Its velocity of ~1jm kus Ἰ auplies that the enclosed mass is 73.5«109AZ... | Its velocity of $\sim 105$ km $^{-1}$ implies that the enclosed mass is $\sim 3.5\times 10^6\;M_\odot$ . |
Compleimeitary studies of the A2.17 μι Bre line cinission from this region (Icvbst.etal.1993)) vield a tnass of ~5AL... for the Norther arm and the ceutral Bax. whose dynamics recure a central concentration of ες109AY. within a radius of ~0.17 pe. | Complementary studies of the $\lambda2.17\;\mu$ m $\gamma$ line emission from this region \cite{herb93}) ) yield a mass of $\sim 5\;M_\odot$ for the Northern arm and the central Bar, whose dynamics require a central concentration of $\sim
4\times10^6\;M_\odot$ within a radius of $\sim 0.17$ pc. |
These carly mass determinations have been supported by subsequent measurements of the stellar velocities aud velocity dispersions at various distances from Ser ÀA*. | These early mass determinations have been supported by subsequent measurements of the stellar velocities and velocity dispersions at various distances from Sgr A*. |
Followiug the initial work bv Sellereu et al. ( | Following the initial work by Sellgren et al. ( |
1987). aud. Ricke Rieke (1988). Taller et al. ( | 1987), and Rieke Rieke (1988), Haller et al. ( |
1996) used the velocity dispersions of sars af zo. pe from Ser A* ο derive a compact mass of ~2«109AL... | 1996) used the velocity dispersions of stars at $\ga 0.1$ pc from Sgr A* to derive a compact mass of $\sim 2\times 10^6 \;M_\odot$. |
This is cousisteut with the value of ~2453.2«109M. derived by Geuzel et al. ( | This is consistent with the value of $\sim 2.5-3.2\times 10^6\;M_\odot$ derived by Genzel et al. ( |
1996). using the radial velocities aud velocity dis)orsions of 25 early-type stars aid of ~200 red giauts am supereiuits within the central 2 pe. | 1996), using the radial velocities and velocity dispersions of $\sim 25$ early-type stars and of $\sim 200$ red giants and supergiants within the central 2 pc. |
A third technique for tracing the ceutral ¢eravitational potential is based onthe acquisition of proper miotious for he ~50 brightest stars within fje radial range ~0.001Wl pe (Eckart&Cenzel 1996)). | A third technique for tracing the central gravitational potential is based onthe acquisition of proper motions for the $\sim
50$ brightest stars within the radial range $\sim 0.004-0.4$ pc \cite{eck96}) ). |
These stellar motions also seen to require a central dark mass of 2003«108AL... in good agreement with both the ionized eas kinematics aud the velociiN dis])orsioni nieasurenieits. | These stellar motions also seem to require a central dark mass of $2-3\times 10^6\;M_\odot$, in good agreement with both the ionized gas kinematics and the velocity dispersion measurements. |
Of course. showing tvat the GC ust contain a centralized mass concentration does nof necessarily inplv that this dark nat cris in the form of a compact object with a ew million solar masses. | Of course, showing that the GC must contain a centralized mass concentration does not necessarily imply that this dark matter is in the form of a compact object with a few million solar masses. |
It docs uot even imply that the 1unusual radio source Ser Α must be associated with it. | It does not even imply that the unusual radio source Sgr A* must be associated with it. |
However. it is possible to demonstrate that Ser A* is probably not stellaa-like. | However, it is possible to demonstrate that Sgr A* is probably not stellar-like. |
This is based οi the fact that a heavy object in dynamical equilibrium wih the surrounding stellar cluster will move skNel. so that a failure to detect proper motion in Ser A* nay be used to provide au independent estimae of its mass, | This is based on the fact that a heavy object in dynamical equilibrium with the surrounding stellar cluster will move slowly, so that a failure to detect proper motion in Sgr A* may be used to provide an independent estimate of its mass. |
Using VLBI.Backer (1991) derived a lower mass lint of ~202000AL, . Which appears to rule out the possibility that Ser A* | Using VLBI,Backer (1994) derived a lower mass limit of $\sim 20-2000\;M_\odot$ , which appears to rule out the possibility that Sgr A* |
The standard paracdiem kuown as the ACDAL cosimoloey includes dark compoucuts: dark matter aud dark energy (2).. | The standard paradigm known as the $\Lambda$ CDM cosmology includes dark components: dark matter and dark energy \citep{Komatsu2011}. |
Weal: lensing of backerouud galaxies bv foreground large-scale structures. the so-called "cosmic shear. has been recognized as a potentially powerful tool for probing the distribution of dark matter as well as the uature of dark energv (?). | Weak lensing of background galaxies by foreground large-scale structures, the so-called “cosmic shear”, has been recognized as a potentially powerful tool for probing the distribution of dark matter as well as the nature of dark energy \citep{Albrecht2006}. |
Reports of significant detections of cosmic shear have been made by various eroups B | Reports of significant detections of cosmic shear have been made by various groups \citep{Bacon2000,VanWaerbeke2000,Wittman2000,Hamana2003,Jarvis2006,Semboloni2006,Fu2008,Schrabback2010}. |
y analyzing the cosimic shear data. one can directly measure the power spectrum of the matter density fluctuations on cosmological scales. which contain a wealth of cosimological information such as neutro lasses aud dark energy. equation-of-state parameters (?777j).. | By analyzing the cosmic shear data, one can directly measure the power spectrum of the matter density fluctuations on cosmological scales, which contain a wealth of cosmological information such as neutrino masses and dark energy equation-of-state parameters \citep{Jarvis2006,Semboloni2006,Ichiki2009,Schrabback2010}. |
Consequently. it is a main goal of cosinoloey to infer and constrain these quantities from observations. | Consequently, it is a main goal of cosmology to infer and constrain these quantities from observations. |
To do this. a nuuber of ambitious surveys are planned. such as the Hyper Suprimc-Cam Weak. Leusing Survey (7)tinal. the Dark Encrey Survey(DES)!.. the Large Synoptic Survey Telescope(LSST?.. Euclid (?) and the Wide-Field Infrared Survey Telescope(WEIRST)?. | To do this, a number of ambitious surveys are planned, such as the Hyper Suprime-Cam Weak Lensing Survey \citep{Miyazaki2006}, the Dark Energy Survey, the Large Synoptic Survey Telescope, Euclid \citep{Refregier2010}, and the Wide-Field Infrared Survey Telescope. |
. Most weak-leusing information is contaimed i siall angular scales auc therefore weak-lensing statistics are uonlinear aud non-CGaussiou 2?(2?).. | Most weak-lensing information is contained in small angular scales and therefore weak-lensing statistics are nonlinear and non-Gaussian \citep{Munshi2008,Takada2009,Sato2009,Sato2011a}. |
Tf we aim to exploit the full information. we have to reat the nonlinear effects to accurately model the weak leusiug statistics. | If we aim to exploit the full information, we have to treat the nonlinear effects to accurately model the weak lensing statistics. |
Furthermore. oue has to use an appropriate likelihood function with Ooeiveu mareinalOo distributious. otherwise the obtained results would besystematically biased (2?).. | Furthermore, one has to use an appropriate likelihood function with given marginal distributions, otherwise the obtained results would besystematically biased \citep{Sato2010,Sato2011}. |
Iu a firs colupanion paper (C7.hereafterpaperD.. we studied the Fourier-3pace weak-leusing statistics such as the weak-leusing power spectrum aud bispectrum. and found that our model proposed by ?7.. which combines rturbation theory with halo models. agrees better with ray-tracing simulations than previously published fitting ornmlas and phenomenological models. | In a first companion paper \citep[][hereafter paper I]{Valageas2011f}, we studied the Fourier-space weak-lensing statistics such as the weak-lensing power spectrum and bispectrum, and found that our model proposed by \citet{Valageas2011d,Valageas2011e}, which combines perturbation theory with halo models, agrees better with ray-tracing simulations than previously published fitting formulas and phenomenological models. |
Tn this second oxeper. we study the realspace weak-leusing statistics. which are more often used for the statistical analysis of actual measurements than Fouricr-space statistics. vecamse observations are wade i configuration space. | In this second paper, we study the real-space weak-lensing statistics, which are more often used for the statistical analysis of actual measurements than Fourier-space statistics, because observations are made in configuration space. |
Previous works have already shown that on small scales the halo model provides a good description of he two-. three- and four-point correlations or snoothled uonienuts of the cosmic shear (using sole approximations) (222). whereas ao stochastic halo model cau recover he probability distribution function of the uunsinoothed convergence (27).. | Previous works have already shown that on small scales the halo model provides a good description of the two-, three- and four-point correlations or smoothed moments of the cosmic shear (using some approximations) \citep{Takada2002,Takada2003,Benabed2006}, whereas a stochastic halo model can recover the probability distribution function of the unsmoothed convergence \citep{Kainulainen2011,Kainulainen2011a}. |
Tere we include all "one-halo. του and "threc-halo terms. as well as one-loop orturbative results. and we compare these with larger-scale simulations. | Here we include all “one-halo”, ``two-halo'' and “three-halo" terms, as well as one-loop perturbative results, and we compare these with larger-scale simulations. |
This vields a ereater accuracy and allows us to compare these different contributions. from very large to simall scales. | This yields a greater accuracy and allows us to compare these different contributions, from very large to small scales. |
This should be useful for practical purposes because these different terms have different theoretical accuracies aud probe different regimes of eravitational clustering. hence it is important to know their relative impact ou weak-leusing probes as a function of angular scale. | This should be useful for practical purposes because these different terms have different theoretical accuracies and probe different regimes of gravitational clustering, hence it is important to know their relative impact on weak-lensing probes as a function of angular scale. |
This paper is organized as follows. | This paper is organized as follows. |
In Sect. | In Sect. |
2 we brictiv recall how coufiguratiou-space woeak-leusius statistics are colmputed from polyspectra of the 3D amatter density field. | \ref{Analytic} we briefly recall how configuration-space weak-lensing statistics are computed from polyspectra of the 3D matter density field. |
We describe our ummerical simulations aud the data analysis in Sect. ὃν, | We describe our numerical simulations and the data analysis in Sect. \ref{Numerical}. |
Then. we present detailed commparisous between the simmlation results. previous models. aud our theoretical predictions for two-point functions in Sect. | Then, we present detailed comparisons between the simulation results, previous models, and our theoretical predictions for two-point functions in Sect. |
| and three-point functions iu Sect. 5.. | \ref{Lensing-two-point}
and three-point functions in Sect. \ref{Lensing-three-point}. |
We study the relative iuportance of the different contributions arising from "one-halo. "two-halo. or "three-halo terms iu Sect. 6.. | We study the relative importance of the different contributions arising from “one-halo”, “two-halo”, or “three-halo” terms in Sect. \ref{contributions}. |
Then. we check the robustuess of our model when we vary the cosmological parzuueters in Sect. | Then, we check the robustness of our model when we vary the cosmological parameters in Sect. |
7 aud we brietiv PAπαν multi-scale ποιοιτς iu Sect. &.. | \ref{Cosmology} and we briefly study multi-scale moments in Sect. \ref{multi-scale}. . |
Finally. we couclude oe1 Sect. 9. | Finally, we conclude in Sect. \ref{Conclusion}. . |
fluctuations. probably because they are obtained frou the scatter between just four spectra. | fluctuations, probably because they are obtained from the scatter between just four spectra. |
For this reason I use the theoretical errors throughout. | For this reason I use the theoretical errors throughout. |
The first analysis was to see whether differeutial spectroplotometry in the relative spectra is reliable. | The first analysis was to see whether differential spectrophotometry in the relative spectra is reliable. |
rofres1 shows this for the three lauds jl. aand ki. | \\ref{res1} shows this for the three bands j1, and k1. |
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