Introductory Summer Research Group Meeting Spectroscopy and X-rays 28 May 2001 Members of the group: David Cohen Allison Adelman David Conners Carie Cardamone Elliot Reed Rachel Sapiro and Matti Klock Geneviève de Messières Kate Penrose

Project Summaries Allison – ROSAT X-ray survey of O stars David – Inertial fusion modeling, ultraviolet spectroscopy of hot star winds Carie – spectral analysis of Chandra spectrum of  Sco and XMM spectrum of  CMa Elliot – spectral analysis of Chandra spectrum of  Sco Rachel – line ratio analysis and modeling of UV spectral data from SSX (Michael Brown’s experiment), hydrodynamical and spectral modeling of gas cell photoionization experiments on the Z-machine at Sandia National Lab Matti – hydrodynamical modeling of table-top laser ablation experiments both in the Moscatelli lab and at Sandia National Lab

contact information David: Allison: David: Carie: Elliot: Rachel: Matti: Geneviève: Kate: Geneviève (Chandra) and Kate (fusion) are at home this summer, but are both continuing to work part-time on their projects -- those of you whose work overlaps with theirs should feel free to get in touch with them if you have any questions

Our individual projects are part of a bigger pursuit Understanding the interactions of very strong electromagnetic radiation fields (i.e. light) with matter. Contexts are (1) hot star winds, (2) laser-produced fusion plasmas, (3) X-ray photoionized plasmas. We want to use spectroscopy to determine the properties of (i.e. diagnose) experimental and astrophysical systems. In practice, this means determining temperatures, ionization states, densities, velocities, incident radiation field spectra, … This is an intermediate step in gaining a global understanding of the physical systems in question.

The physical systems we’re trying to understand are (1) stellar winds (supersonic outflows from the surfaces of stars) (2) X-ray binaries/X-ray nebulae (gas clouds surrounding black holes and neutron stars) (3) laser fusion experiments

Solar-type magnetic heating? Or massive stellar wind shock heating? On the stellar wind side, the big specific question we’re trying to answer is: How do hot stars make their X-rays? magnetic/coronal mechanism vs. wind shocks

High intensity light invariably means that things get hot… Hot things (above about half-a-million degrees) means X-rays (remember, there’s a relationship between color/wavelength and temperature) Spectroscopy provides information, through its dependence on atomic states and the Doppler shift, about the microphysics, as well as bulk properties, of gases (plasmas, if they’re ionized). A spectrum carries a huge amount of information: “If a picture is worth a thousand words, then a spectrum is worth a thousand pictures.” -- anon.

The spectral resolution of x-ray telescopes has improved many hundred-fold over the past decade, enabling us to detect and resolve emission lines in hot stars like  Sco. Spectral resolution, defined: Where  is the width of the narrowest feature detectable, and is the wavelength of that feature. Thus, the resolution is a (rough) measure of the total number of independent “channels” or pieces of information in a spectrum.

ROSAT (1993), resolution E/  E ~ 2 - With only a few independent pieces of information, the shape of this smooth-looking spectrum tells us only about the relative amount of ‘hard’ to ‘soft’ X-rays (i.e. those with energies above 0.5 keV to those with energies below this level).

ASCA (1997) resolution E/  E ~ 40 - Some of the strongest line complexes are just visible above the pseudo-continuum of blended weak lines. The number of independent elements of information is just not great enough to distinguish individual emission lines.

Chandra (2000) resolution E/  E ~ We can now see numerous emission lines; and not just see them, but resolve them, measuring their intrinsic widths.

Background Astronomers measure wavelength,, in Angstrom Units (Å) Remember, frequency,, and wavelength of any wave phenomenon are related by: =c Where c is the characteristic wave velocity (speed of light for light, duh…) Frequency is directly proportional to photon energy: E=h And inversely proportional to wavelength: E=hc/ When we talk about single photons, the unit of energy is the electron volt (eV) 12.4 Å = 1000 eV or 1 keV; X-rays are taken to be any photons with energies greater than 100 eV (0.1 keV) and wavelengths, therefore, below 124 Å Note that since temperature is proportional to energy, we can relate light of a given wavelength to a characteristic temperature through Boltzmann’s constant, k: E=kT Where 1 keV corresponds to 11.6 million degrees

The photons of different wavelength, that together constitute a spectrum, are produced by a handful of specific atomic processes Much of the power of spectroscopy comes from this macro-micro connection (observed spectrum vs. movement of electrons inside atoms) -- What is an atom? To us, it is a central electrostatic potential plus a bunch of quantum mechanically determined bound energy levels and a “continuum” of unbound states. Think potential energy: energy is like a staircase continuum ground state In both these cases, energy is released an atom

Let’s consider some specific atomic processes and see which ones lead to either the production of photons (emission) or the destruction of photons (absorption) simplified 2-level atom (plus continuum) key  previous position of an electron current position of an electron movement of an electron movement of a photon photon e external electron

1. Radiative excitation key  previous position of an electron current position of an electron movement of an electron movement of a photon photon e external electron  Photon must be of exactly the right energy Photon is destroyed (absorbed)

2. Radiative de-excitation (“spontaneous emission”) key  previous position of an electron current position of an electron movement of an electron movement of a photon photon e external electron  Photon energy reflects energy difference of electron bound states Photon is produced (emitted)

3. Collisional excitation key  previous position of an electron current position of an electron movement of an electron movement of a photon photon e external electron Free electron need only have more energy than the energy difference between states Free electron is not destroyed, but loses some kinetic energy (note potential for cooling…) e e

4. Collisional de-excitation key  previous position of an electron current position of an electron movement of an electron movement of a photon photon e external electron Free electron is not destroyed, but gains some kinetic energy e e

Collisional excitation followed by spontaneous (radiative) emission is the most important process in X-ray-emitting cosmic plasmas e e  It is a cooling process (taking kinetic energy from free electrons -- i.e. heat) and converting it to photons, which can escape the plasma (and bring information to our telescopes about the conditions in the emitting plasma).

These basic atomic processes lead to a wealth of phenomena in X-ray spectra. An X-ray spectrum reflects the physical properties of the source (the electron density and temperature, radiation flux and temperature, bulk motions). These properties can be used to discriminate among different physical models of a laboratory or astrophysical object.

The following slides are supplemental – read at your own risk

The context of X-ray spectral observations of hot stars Hot stars are thought not to have outer convection zones, magnetic fields, or the associated magnetic dynamo and corona that our sun has. Thus their discovery 20 years ago as relatively strong soft X-ray sources was a surprise. Hot stars do have strong radiation-driven winds. These winds are subject to a line-driving instability which can lead to shock heating of the wind plasma. Although this mechanism has been assumed to produce the observed X-rays, the numerical simulations do not do a very good job of quantitatively reproducing the observed X-rays. The new high-resolution X-ray spectra of hot stars holds out the potential to discriminate between these two general theories of stellar X-ray emission; or possibly to inspire the development of a new theory.

energy 1s 2 1 S 1s2s 3 S 1s2p 3 P 1s2p 1 P F R I A partial energy level diagram for helium-like ions, such as Si +12 (see the FIR panel inset on the Chandra spectrum figure). The resonance transition (R) is strongest, but the intercombination (I) and forbidden (F) lines can also be strong. Electrons in the long- lived, or metastable, 3 S state can be collisionally or radiatively excited to the 3 P level (transition energy of order 10 eV), making the F/I ratio a good diagnostic of density and radiation field (effectively distance from the star). The forbidden line of the silicon feature in the  Sco spectrum is strong, indicating that the X-ray emitting plasma is both far from the star and low density…but how far and how low?

Line profiles are affected by the hot plasma’s spatial and velocity distribution, as well as the degree of attenuation by an overlying cold stellar wind. In the next panel I show synthetic line profiles for a family of coronal models (left) and wind-shock models (right). Line profiles as a function of the scaled wind velocity (x=c  / o v  ) are shown for different instrumental resolutions in each sub-panel. The panels have wind attenuation increasing downward. Note that as the wind attenuation increases less and less of the red side of the profiles are seen. This is because the redshifted wind necessarily comes from the back side of the star, which suffers the most absorption. The coronal models assume an X-ray source that is strongly concentrated at the base of a slowly accelerating wind, while the wind-shock models assume that the source is distributed within the supersonic outflow beyond some minimum radius. The narrow features seen in  Sco would seem to point to a mostly coronal origin to these X-rays; but the strong forbidden lines indicate that the hot plasma is not very close to the surface of the star.

    CoronalWind-shock Instrumental broadening  v  = 0, 0.1, 0.3, 0.5    x lxlx